{
    "0402/astro-ph0402528_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:introduction}  Quite often in the history of science major discoveries came from people who were not looking for them. This was certainly the case of the first observation of the Cosmic Microwave Background (CMB) radiation by Penzias and Wilson in 1964-65 \\cite{penzias65}: they were lucky enough to find themselves at the right time and place, and also skilled enough to pursue to the very end the hint that Nature was offering them. But besides luck and skill two major factors were decisive in this discovery.  First, the right technology was starting to become available. Advances in millimetre-wave technology after the end of World-War II gave a boost to the young field of radio-astronomy, begun with the pioneer works of Jansky and Reber \\cite{reber40, reber44}. The horn antenna and low noise receiver at Bell Labs was probably at the time the most sensitive microwave instrument on Earth.  Second, in the mid 60's the scientific stage was ready for this breakthrough. In 1929 Edwin Hubble's observation of the recession of galaxies revolutionised the vision of the Universe introducing the concept of ``evolution'' at cosmological scales \\cite{hubble29}. About a decade before Hubble's discovery, Alexander Friedmann showed that Einstein's General Relativity equations were compatible with a variety of scenarios, both static (as developed by Einstein by introducing the ``Cosmological Constant'', $\\Lambda$) and dynamic.  In the mid 1940's, George Gamow and his group developed a physical theory of the early Universe that was compatible with both General Relativity and Hubble's observations. They extrapolated physics back in time to the point where temperature and density were sufficiently high to support nuclear fusion, and studied the production of heavy elements from primordial protons. A remarkable side-prediction of the calculations of Gamow, Halper and Hermann (see, e.g., \\cite{gamow46}) was the presence of a photon field in the ``primordial fireball'', now red-shifted and cooled to very low temperatures by cosmic expansion. Penzias and Wilson, unaware of Gamow's results, observed for the first time this relic, isotropic, radiation at $\\sim 3$~K in the range of microwaves. Their discovery was more than an experimental confirmation of a theoretical prediction: it showed that it was possible to measure directly (and precisely) properties of the Universe in a very young state, when all the processes were still in the linear regime, before structures had formed yet.  Experimental and theoretical research in cosmology has been growing steadily since then, and CMB observations have been paying a major role in this growth. In the first decade or so after the CMB discovery, major efforts were aimed at understanding the nature of the CMB radiation itself and characterising its frequency spectrum. By the early 80's the Hot Big Bang prediction of a highly isotropic background with a nearly planckian spectrum was remarkably supported by observation. The scientific community gradually became more and more interested in the study of spectral distortions and spatial anisotropies of the CMB. In particular, anisotropies were believed to be present in order to explain the existence of local non-uniformities in the present-time matter distribution.  A new phase was opened up by the COBE mission in the early 90's. The FIRAS instrument \\cite{fixsen96} measured the CMB spectrum to be planckian at the level of 99.97\\% in the frequency range 60-600~GHz with a temperature of $T_0 = 2.725\\pm 0.002$~K. Coupled with sub-orbital measurements at low frequencies \\cite{bersanelli94, deamici91, bensadoun93, smoot87, sironi91, salvaterra02} very stringent constraints to spectral distortion parameters were placed, leading to tight upper limits on energy injections in the early Universe. FIRAS demonstrated that ``precision cosmology'' is possible with accurate measurements of the CMB, and confirmed the maturity achieved by microwave and sub-mm technology.  A second breakthrough came from COBE with the DMR instrument which provided the first unambiguous detection of anisotropies at the level of $\\Delta T/T \\approx 10^{-5}$ on large angular scales ($\\sim 7^\\circ$) \\cite{smoot92}. This result immediately stimulated many new experiments aiming at measuring the CMB angular distribution with increasing resolution and sensitivity. This explosion of experimental effort was motivated by the realisation that accurate measurements of the statistics of CMB anisotropies, reflected in its angular power spectrum from large scales to about $10'$, yield powerful constraints on fundamental parameters of cosmology such as the Hubble constant, $H_0$, the baryon density, $\\Omega_b$, the dark energy density, $\\Omega_\\Lambda$, etc. To date, more than 20 independent projects have been carried out with different technologies and from a variety of ground-based and balloon-borne experiments, recently with remarkable precision.  In early 2003 the NASA space mission WMAP has released the first full-sky map of the CMB with sub-degree angular resolution, setting tight constraints on many cosmological parameters. The ESA Planck mission, to be launched in 2007, will exploit CMB temperature anisotropy measurements to its fundamental limits imposed by unavoidable cosmic variance and astrophysical foregrounds, and will likely open a new phase in CMB science aimed at a precise measurement of anisotropies in the CMB polarisation state.  Although technological advances have allowed a steady increase in the sensitivity of CMB measurements, it has also increased the level of complexity of instruments and satellites. For example, the quest for high sensitivity leads to highly sophisticated cooling systems and to multi-feed arrays, which represent new challenges for the thermal and optical design of the instruments. The combination of instrument complexity and high sensitivity leads to very severe requirements in terms of rejection of systematic effects. In the second and third generation CMB space missions (namely WMAP and Planck) the systematic error control has become one of the most (if not {\\em the} most) critical experimental challenge, often pushing the understanding of current technologies into poorly known grounds.  Many excellent CMB reviews have been published covering both theoretical and experimental aspects (see, e.g., \\cite{bersanelli02} and references therein); however, the rapid evolutions in CMB science and in its related technologies often renew the need of state-of-the-art analysis. In this paper we present an overview of the main experimental issues of current efforts devoted to precision CMB measurements and discuss the main implications for the future challenges represented by precision polarisation anisotropy measurements. In particular the paper presents a discussion about the control of systematic effects in second and third generation space CMB experiments.  After a brief summary of the CMB theoretical background (Sect.~\\ref{sec:theoretical_background}) and of the main astrophysical limitations (Sect.~\\ref{sec:astrophysical_limitations}) we will review the main issues concerning the control of systematic errors in high-precision CMB measurements from space, with examples taken from the (still growing) experience formed in the context of the Planck mission. Sect.~\\ref{sec:cmb_experiments} gives a short account of the evolution of CMB experiments from COBE to Planck, highlighting the deep relationship between progress of technology  and increase in the scientific achievements as well as experimental challenges. In the last section we provide a discussion of our foreseen scenarios in CMB science after Planck. ",
        "conclusions": "After the discovery of CMB temperature anisotropies by COBE, a great experimental effort has been accomplished with ground-based and balloon-borne experiments and recently by the WMAP space mission, leading to a good determination of the angular power spectrum up to sub-degree scales. The acoustic nature of the spectrum, expected by theory, is now beautifully matched by observation. This is a truly remarkable achievement. The recent detection of polarisation E-modes and the TE correlation represents a further striking confirmation of the soundness of standard cosmology and provides new incentive to the experimental growth of CMB research. However, these rewarding results are probably only a foretaste of the precision observation attainable by the forthcoming generations of experiments, culminating with the Planck mission. The Planck survey will generate an unprecedented set of multi-frequency temperature and polarisation data, leading to major improvement in the determination of cosmological parameters and to a new profound verification of our cosmological understanding.  Beyond Planck, it is likely that CMB observations will concentrate on precision polarisation measurements (especially searching for gravitational wave signatures in the B-mode spectrum) and deep sub-arcmin imaging of secondary anisotropies. Large arrays of cryogenic detectors, now under study, should be capable of reaching the extreme instrument sensitivity required. However, this will not be enough. The experience of the past decades has shown that CMB measurements are typically limited by systematic effects, either of instrumental or astronomical nature, some of which have been discussed in this review. These limitations will need to be explored at a much deeper level in future high-precision enterprises.  Precision measurements of the CMB are able to shed light on very high energy phenomena occurring in the primordial cosmic environment, as well as on the physical history of the Universe. Today, observations of the CMB promise to remain one of the most powerful cosmological probes for yet many years in the future. \\\\  \\noindent{\\bf Acknowledgements}  The completion of this review has been greatly helped, either directly or indirectly, by the work of many people, in particular by the Planck Science Team, by the Planck-LFI consortium and by the Boomerang and MAXIMA teams. We would also like to thank the WMAP science team for figure permission and  L.A. Popa, D. S\\'aez, L. Toffolatti, L. Danese and G. De Zotti for useful discussions."
    },
    "0402/astro-ph0402002_arXiv.txt": {
        "abstract": "We investigate the required redshift accuracy of type Ia supernova and cluster number-count surveys in order for the redshift uncertainties not to contribute appreciably to the dark energy parameter error budget.  For the SNAP supernova experiment, we find that, without the assistance of ground-based measurements, individual supernova redshifts would need to be determined to about 0.002 or better, which is a challenging but feasible requirement for a low-resolution spectrograph. However, we find that accurate redshifts for $z<0.1$ supernovae, obtained with ground-based experiments, are sufficient to immunize the results against even relatively large redshift errors at high $z$. For the future cluster number-count surveys such as the South Pole Telescope, Planck or DUET, we find that the purely statistical error in photometric redshift is less important, and that the irreducible, systematic bias in redshift drives the requirements. The redshift bias will have to be kept below 0.001-0.005 per redshift bin (which is determined by the filter set), depending on the sky coverage and details of the definition of the minimal mass of the survey.  Furthermore, we find that X-ray surveys have a more stringent required redshift accuracy than Sunyaev-Zeldovich (SZ) effect surveys since they use a shorter lever arm in redshift; conversely, SZ surveys benefit from their high redshift reach only so long as some redshift information is available for distant ($z\\gtrsim 1$) clusters. ",
        "introduction": "Two of the most promising methods to measure cosmological parameters, in particular those describing dark energy, are distance measurements of type Ia supernovae (SNe Ia) and number counts of clusters of galaxies in the universe. SNe Ia have provided original direct evidence for dark energy (Riess et al.\\ 1998, Perlmutter et al.\\ 1999) (for earlier, indirect evidence, see Krauss and Turner\\ 1995 or Ostriker and Steinhardt\\ 1995) and are currently the strongest direct probe of the expansion history of the universe (Tonry et al.\\ 2003, Knop et al.\\ 2003). Their principal strength is the simplicity of relating the observable -- which is essentially the luminosity distance -- to cosmological parameters, and also the fact that each supernova redshift-magnitude pair provides a distinct measurement of a combination of those parameters. Number-counts, on the other hand, use the fact that galaxy clusters are the largest collapsed structures in the universe that have undergone a relatively small amount of post-processing.  Their distribution in redshift can be reliably calculated in a given cosmological model. The evolution of cluster abundance is principally sensitive to the comoving volume and growth of density perturbations (Haiman, Mohr \\& Holder 2001) and this cosmological probe is expected to reach its full potential with upcoming and future wide-field surveys.  Rapid improvement in the accuracy of measuring cosmological parameters implies that various systematic uncertainties, previously ignored, now have to be controlled and understood quantitatively. In the case of supernova measurements, an example is provided by the proposed SuperNova/Acceleration Probe (SNAP) satellite (Akerlof et al.\\ 2004) whose goals for measuring the equation of state of dark energy $w$ and its variation with redshift $dw/dz$ drive the requirements on the systematic control that are considerably more stringent than those attainable with current surveys. Similarly, the principal systematic difficulty in cluster counts is in establishing the relation between observable quantities (X-ray temperature or Sunyaev-Zeldovich flux decrement), and the cluster's mass which is necessary for comparison with theory. The mass-temperature relation, for example, is known to have a considerable scatter and is currently poorly determined, with fairly large intrinsic statistical errors and considerable systematic disagreements between different authors (see e.g. Fig. 2 in Huterer \\& White 2002). The cleanest way to include the mass-observable relation might be to determine it from the survey itself (this is known as ``self-calibration''; Levine, Schulz \\& White 2002, Majumdar \\& Mohr 2003, Hu 2003, Lima \\& Hu 2004), but this will almost certainly lead to degradations in parameter accuracies. Future surveys will require a careful accounting of all systematics -- theoretical and observational.  In this paper we concentrate on one of the most basic ingredients of supernova and cluster count measurements: the determination of redshift.  In the case of SNe Ia spectroscopic observations are necessary to identify the supernova type, and redshift is then supplied for free. Recently completed and ongoing surveys have sufficiently poor magnitude uncertainty and relatively low statistics and relatively weak control on known systematics, so that the spectroscopic redshift error is small enough for the redshifts to be considered perfectly known. However, as we shall see, future supernova observations require such accurate redshifts that even the spectroscopic accuracy is not {\\it a priori} guaranteed to be sufficient.  In the cluster count case, the situation is even more interesting, as spectroscopic observations will not be possible for all clusters, which may number in the tens of thousands.  One will therefore rely on photometric redshifts.  Although photometric redshifts are already impressively accurate (e.g. Fern\\'{a}ndez-Soto et al.\\ 2002, Csabai et al.\\ 2003, Collister \\& Lahav 2003, Vanzella et al.\\ 2003), we shall find that their bias (the difference between the mean photometric value and the true value at any redshift) needs to be kept exceedingly small for the redshift error not to contribute appreciably to the total error budget. Our analysis is timely, as follow-up surveys to obtain cluster redshifts, such as that at Cerro-Tololo International Observatory, are about to get underway soon.  Our analysis also complements recent analysis of the effect of systematic errors on future SN Ia measurements (Kim et al.\\ 2003, Frieman et al.\\ 2003) and a variety of related analyses regarding the cluster number-count surveys (e.g. Bartlett 2000, Holder \\& Carlstrom 2001, White, Hernquist \\& Springel 2002, White, van Waerbeke \\& Mackey 2002, Benson, Reichardt \\& Kamionkowski 2002, White 2003).  The paper is organized as follows. In section~\\ref{sec:method} we outline the procedure to include the redshift uncertainty in the standard Fisher-matrix parameter estimation. In Sec.~\\ref{sec:sne} we discuss the redshift requirements for future supernova surveys, while in Sec.~\\ref{sec:counts} we do the same for future cluster count surveys. We conclude in Sec.~\\ref{sec:concl}. Our fiducial model is a flat universe with matter energy density relative to critical of $\\Omega_M=0.3$ and the equation of state of dark energy $w=-1$.  Other cosmological parameters, necessary for the cluster abundance calculation, are discussed in Sec.~\\ref{sec:counts}. ",
        "conclusions": "\\label{sec:concl}  We considered how inexact redshifts affect future SNe Ia and number-count surveys. We treated the redshifts as additional parameters whom we assigned priors equal to their assumed measurement accuracy. Requiring that the redshift uncertainty do not contribute more than $\\sim 10\\%$ to the error budget in cosmological parameters, we imposed requirements on the redshift accuracy.  For a future survey that studies $\\sim 3000$ supernovae out to $z=1.7$ (e.g.\\ the SNAP space telescope) we find that, with accurate redshift measurements of $dz \\lesssim 0.001$ for $z<0.1$ supernovae, fairly poor redshift measurements can be tolerated at higher redshifts. Without this accurate measurement at low redshift, however, a fairly precise redshift measurement of $dz \\lesssim 0.002$ would be required over the full redshift range. Photometric redshifts are probably not an option, since spectral information is necessary to identify the SN type and control a variety of systematic errors.  Spectroscopy can be provided using sub-pixel interpolation of galaxy data from an on-board low-dispersion $R \\sim 100$ spectrograph (which is designed to measure broad supernova features).  Supplemental high-resolution ground-based observations using 10m-class telescopes, adaptive optics, and OH suppression can provide precise redshifts as necessary and to cross-check the redshifts from the low-dispersion spectrograph.  We thus conclude that redshift uncertainty will not significantly contribute to the error budget in the accurate measurement of dark-energy parameters that SNAP can deliver.  For future wide-field cluster count surveys, such as SPT, Planck or DUET, we find that the purely statistical errors are largely irrelevant as long as they are reasonably small (error of $\\lesssim 0.02$ per cluster) because they will average out due to the large number of clusters around any given redshift. However, the irreducible, systematic error that doesn't decrease with increasing number of clusters drives the redshift requirements. This irreducible redshift-independent error has to be kept below 0.001-0.005 per redshift bin. The widths of the redshift bins are determined by how the redshift signature (say, the 4000\\AA\\ break line) goes through the filter set of the redshift follow-up experiment, and here for illustration we assumed filters from the Sloan Digital Sky Survey. We found that the typical required redshift accuracy is more stringent for X-ray surveys since they have few clusters at $z\\gtrsim 1$ and therefore use a shorter lever arm in redshift.  SZ surveys benefit from their longer lever arm, but, of course, only if their high-redshift clusters have decent redshift information.  Obtaining redshifts for high-redshift clusters, therefore, should be an important goal of any redshift follow-up survey. While the photometric accuracy at redshifts greater than unity is highly uncertain at present, our analysis indicates that the lack of redshift information at $z\\gtrsim 2$ does not significantly degrade the cosmological constraints, while at redshifts $1\\lesssim z\\lesssim 2$ crude photometric information is sufficient to assure small degradation in constraints on $w$ (see Fig.~\\ref{fig:z_break}). With the current rate of progress in photometric redshift techniques, this should be a feasible goal within the next few years.   \\medskip We thank Jim Bartlett, Josh Frieman, Adrian Lee, and Tim McKay for useful discussions. We particularly thank Eric Linder for pointing out the importance of accurate redshifts for low-$z$ SNe, and Jim Annis and Martin White for comments on an early draft of the paper. DH and LMK are supported by the DOE grant to CWRU. AK was supported by the Director, Office of Science, of the U.S.  Department of Energy under Contract No. DE-AC03-76SF00098."
    },
    "0402/astro-ph0402234_arXiv.txt": {
        "abstract": "The Atacama Cosmology Telescope (ACT) project is described. This multi-institution collaboration aims to produce arcminute-resolution and micro-Kelvin sensitivity maps of the microwave background temperature over 200 square degrees of the sky in three frequency bands. We give a brief overview of the scientific motivations for such a map, followed by a design outline of our six-meter custom telescope, an overview of our proposed bolometer array detector technology, and site considerations and scan strategy. We also describe associated optical and X-ray galaxy cluster surveys. ",
        "introduction": "With results from WMAP in hand, it is clear that the near-term future of microwave background measurements will be primarily a push towards smaller angular scales and polarization.  As described at this meeting, many small-scale experiments are currently underway, under construction, or in the planning stage. This paper describes an ambitious proposed collaboration, the Atacama Cosmology Telescope, which aims to combine new bolometer array technology with a custom-designed six-meter telescope to produce an instrument with a notable combination of sensitivity, angular resolution, and control of systematic errors.  Current information about the experiment is available at http://www.hep.upenn.edu/~angelica/act/act.html.  Before detailing the experimental and observational aspects of this effort, we give a brief summary of the main scientific questions we aim to address, which are covered in more detail elsewhere in these proceedings (see also \\cite{kos03}).  At angular scales smaller than around 4 arcminutes, corresponding to multipoles $l>3000$, nonlinear contributions begin to dominate the total microwave background temperature anisotropies. The major sources of temperature fluctuations on these scales include the Sunyaev-Zeldovich effect, the Ostriker-Vishniac effect, and gravitational lensing. All of these effects arise both from individual clusters of galaxies and from the large-scale matter distribution.  The largest amplitude signal will come from the thermal SZ galaxy cluster distortions. We expect to compile a large catalog of clusters selected by their SZ signals, which provides a cleaner cluster selection criterion than flux-limited optical or X-ray surveys. This catalog will be well-suited to measuring the cluster number density as a function of mass and redshift, which is a sensitive probe of the growth rate of structure since redshift $z=1$; in turn, the structure growth rate constrains dark energy and neutrino masses. We also expect to place significant constraints on cluster masses and peculiar velocities via their kinematic SZ and gravitational lensing signatures. For the diffuse signals, we aim to detect gravitational lensing of the microwave background and construct a projected mass map on scales of 15 arcminutes. We expect to have sufficient sensitivity to detect the Ostriker-Vishniac effect, which is sensitive to the redshift and spatial variation of reionization, and the Rees-Sciama effect, which is sensitive to the non-linear evolution of gravitational potentials.  ACT will provide a measurement of the power spectrum on all scales from $l=200$ to $l=10000$, probing the primordial fluctuation spectrum for departures from a power law or for features; either could arise from non-minimal models of inflation. Having a single experiment which probes a wide range of angular scales with good control of systematics is crucial for uncovering small departures from perfect power law primordial spectra.  ACT's scan strategy results in a significant area of the survey being in the galactic plane. A variety of interesting topics in galactic astrophysics can be addressed with such a map, particularly properties and distribution of dust. ",
        "conclusions": ""
    },
    "0402/astro-ph0402144_arXiv.txt": {
        "abstract": " ",
        "introduction": "It is well known that the most massive globular cluster in the Milky Way, $\\omega$ Cen, shows unique properties in its metallicity content, internal kinematics, and structure. For instance, $\\omega$ Cen shows a wide spread in metallicity unlike other Galactic globular clusters (e.g. Norris, Freeman, \\& Mighell 1996): its metallicity distribution is peaked at [Fe/H]$\\simeq -1.6$, along with a second smaller peak at [Fe/H]$\\simeq -1.2$ and a long tail extending up to [Fe/H]$\\simeq -0.5$. Also, the metal-rich stars in $\\omega$ Cen are largely enhanced in $s$-process elements relative to those in globular clusters and field stars with similar metallicities (e.g. Norris \\& Da~Costa 1995). This suggests that the ejecta from low-mass, asymptotic giant branch (AGB) stars had to be retained and incorporated into the next-generation stars. However, although $\\omega$ Cen is most massive ($5\\times 10^6$ M$_\\odot$), it is unable to retain the AGB ejecta, as shown by Gnedin et al. (2002). Thus, an isolated formation of $\\omega$ Cen is unlikely, because the enriched gas would easily be lost by encountering the Galactic disk. The most viable explanation for the uniqueness of $\\omega$ Cen is that it was once the dense nucleus of a dwarf galaxy, i.e., a nucleated dwarf (Freeman 1993). A gravitational potential of progenitor's dark matter would help retaining the enriched gas and let the cluster being self-enriched at least over a few Gigayears.  In this contribution, we pursue the possible kinematical evidence for the existence of such a dwarf galaxy, which was already disrupted in the past. Dinescu (2002) first investigated this issue, by examining the possible signature of the progenitor's tidal debris among nearby metal-poor stars in the catalog of Beers et al. (2000, B00). She identified a group of stars with $-2.0<$[Fe/H]$\\le-1.5$, which departs from the characteristics of the inner Galactic halo but has retrograde orbits similar to $\\omega$ Cen. Her simplified disruption model of the progenitor galaxy demonstrated that trailing tidal debris, having orbital characteristics similar to the cluster, can be found in the solar neighborhood, although the concrete spatial distribution and kinematics of the debris stars remain yet unclear.  Dinescu's work motivates us to undertake a more refined approach to the issue, i.e., to conduct an N-body simulation for the tidal disruption of $\\omega$ Cen's progenitor galaxy (Mizutani et al. 2003). We obtain the characteristic structure and kinematics of its debris stars and compare with various observations showing signatures of recent merging events in the Milky Way (Gilmore, Wyse, \\& Norris 2002, GWN; Kinman et al. 2003, K03; Chiba \\& Beers 2000, CB). In particular, we show that a recently identified stream of stars at radial velocity of $\\sim 300$ km~s$^{-1}$ (GWN) is a natural outcome of the current disruption model, without significantly modifying local halo kinematics near the Sun. ",
        "conclusions": "Our simple model of an orbiting dwarf galaxy that once contained $\\omega$ Cen predicts a sequence of tidal streams in retrograde rotation and their existence can be imprinted in kinematics of nearby stars, especially in the direction against Galactic rotation (GWN) and at the NGP (K03), while local halo kinematics remain unchanged. The simulated streams are mostly distributed inside the solar circle, as suggested from the current orbital motion of $\\omega$ Cen (DGvA; Dinescu 2002). In contrast to Sgr dwarf galaxy having polar orbit (Ibata et al. 1997), the orbit of $\\omega$ Cen's progenitor galaxy is largely affected by a non-spherical disk potential, where the orbital plane exhibits precession with respect to the Galactic Pole, causing self-crossing of tidal streams in the disk region (Fig. 1). The projection of the orbit perpendicular to the disk plane shows an 'X'-like feature, thereby leaving denser streams at high $|z|$ than at low $|z|$ for a given radius. These characteristic spatial distributions of the debris stars give rise to more significant effects of the debris at the NGP than in the solar neighborhood, as shown here, although the current simulation failed to reproduce the reported largely retrograde rotation at NGP from the $\\omega$ Cen debris alone; perhaps, other, yet unknown halo substructures must be considered to reproduce the observations.   Existing kinematic studies of Galactic stars to search for a signature of $\\omega$ Cen's progenitor galaxy are yet confined to nearby stars, where the significance of the debris streams is modest, as shown here. Searches of stars inside the solar circle are more encouraging (Fig. 1), in particular in the directions of $l \\sim 320^\\circ$ and $l \\sim 50^\\circ$, where we expect the presence of high-velocity streams at $v_{los} = 200 \\sim 300$ km~s$^{-1}$ and $-400 \\sim -300$ km~s$^{-1}$, respectively. Future radial velocity surveys of these fields including the sample of the Sloan Digital Sky Survey or planned Radial Velocity Experiment are worth exploring in this respect. Also, detailed abundance studies of candidate stream stars will be intriguing, because such stars may exhibit different abundance patterns from field halo stars, as found in dwarf galaxies (Shetrone, C\\^{o}t\\'{e}, \\& Sargent 2001).   In this work, we adopt a fixed external gravitational potential for the calculation of an orbiting dwarf galaxy and its dynamical evolution, thereby neglecting dynamical friction against the satellite. However, to set more refined limits on its dynamical history, it is required to fully take into account frictional effect on the orbit as well as dynamical feedback of a satellite on the structure of the Galactic disk. For instance, Tsuchiya, Dinescu, and Korchagin (2003) reported their numerical models for a satellite with strong orbital decay and succeeded to reproduce the current orbit of $\\omega$ Cen from the launch of the progenitor satellite at 58 kpc from the Galactic center. They also showed that the debris particles at 3 Gyr are already smeared out and distributed in a flattened disk without having significant stream-like features as obtained here. It is noted that whether or not stream-like features survive by the current epoch depends on when and how a satellite galaxy is disrupted by Galactic tides, or in other words, the observational information on such features and the comparison with simulation results are useful for placing important constraints on when a satellite merging occurs. For this purpose, we are currently undergoing full N-body simulations of several host-satellite systems using GRAPE5, where both host and satellite galaxies are represented by lively dark halos and stellar components. Our goals with the use of GRAPE5 are to set tight limits on the initial total mass, internal mass distribution, and orbital motion of a progenitor galaxy, as well as the timing of merging with the Galactic disk, based on the comparison between the observation (e.g., RAVE) and simulation results for stream-like structures in the Milky Way. More details will be reported elsewhere (Mizutani \\& Chiba, in preparation)."
    },
    "0402/astro-ph0402658_arXiv.txt": {
        "abstract": "{ In this {\\it Letter} we present a detailed study of the lensing configuration in the cluster Abell 2218. Four multiple-images systems with measured spectroscopic redshifts have been identified in this cluster. These multiple images are very useful to constrain accurately the mass distribution in the cluster core, but they are also sensitive to the value of the geometrical cosmological parameters of the Universe.  Using a simplified maximum likelihood analysis we find $0<\\Omega_{\\rm M}<0.30$ assuming a flat Universe, and $0<\\Omega_{\\rm M}<0.33$ and $w<-\\, 0.85$ for a flat Universe with dark energy. Interestingly, an Einstein-de Sitter model is excluded at more than 4$\\sigma$. These constraints are consistent with the current constraints derived with CMB anisotropies or supernovae studies. The proposed method constitutes an independent test of the geometrical cosmological parameters of the Universe and we discuss the limits of this method and this particular application to Abell 2218. Application of this method with more sophisticated tools and to a larger number of clusters or with more multiple images constraints, will put stringent constraints on the geometrical cosmological parameters. ",
        "introduction": "The present Cosmology framework is characterized by a number of parameters which sets the global geometry of the Universe, its history and dynamics. The quest for these parameters is a long-standing issue in Observational Cosmology and is still the main driver of a large number of experiments.  Combining constraints coming from the power spectrum of the CMB anisotropies and the luminosity distances of distant type Ia supernovae (SNIa), a new standard model of cosmology is emerging \\citep{spergel03}: a flat Universe with an accelerating expansion ($\\Omega_{\\rm M} \\simeq 0.27$ and $\\Omega_\\Lambda \\simeq 0.73$).  To quantitatively explain these results the concept of dark energy has been put forward, characterized by the ratio of pressure and energy density $w=P_{\\rm X}/\\rho_{\\rm X} \\, c^2$, which reduces to the vacuum energy (the cosmological constant) for $w=-1$.  There is however, no strong observational constraints on $w$ yet (\\citealt{spergel03} give only $w<-\\, 0.78$).  This new standard cosmology is getting very popular. Although the flatness of the Universe seems robust, the exact value of $\\Omega_{\\rm M}$ is still a matter of debate \\citep{bridle03,blanchard03} as it is essentially driven by the SNIa results which can be discussed \\citep{rowan02}.  In order to independently probe the large scale geometry of the Universe, we propose to explore the potential use of cluster lenses as a long range optical test bench. Preliminary analysis of this method was first detailed by \\citet{link98} using simple lens models. Recently, we extended their work using more detailed simulations of realistic clusters of galaxies \\citep{golse02a}. The basic idea of this method is that each set of multiple-images identified in a cluster lens strongly constrains the cluster potential.  As the scaling of the mass model depends on the ratio of the angular distances $D_{\\rm LS}/D_{\\rm OS}$, it will also depends on the geometrical cosmological parameters ($\\Omega_{\\rm M}$, $\\Omega_\\Lambda$ and $w$).  In order to constrain these parameters, the combination of several sets of multiple images in a single lens is mandatory to disentangle between the degeneracies in the lens model. With a minimum of 4 systems of multiple images with known spectroscopic redshifts, we showed \\citep{golse02a} that one can put reasonable constraints in the ($\\Omega_{\\rm M} , \\Omega_\\Lambda$) plane with some characteristic degeneracies in the fitted parameters \\citep{golse02a}.  In this {\\it Letter} we apply this lensing test to the well studied cluster-lens Abell 2218.  Section 2 describes details of the lens modeling and the cluster mass distribution. The results of the optimization are discussed in Section 3 and a conclusion is presented in Section 4.  When necessary we scale the physical parameters with $\\Omega_{\\rm M} = 0.3, \\ \\Omega_\\Lambda = 0.7, \\ h=0.65$ with Hubble constant $H_0 = 100 \\ h$ km s$^{-1}$ Mpc$^{-1}$. Thus at the cluster redshift $z=0.176$, $1''$ corresponds to a linear scale of $2.08\\,h^{-1}$ kpc. ",
        "conclusions": "We have shown in this paper that we can derive reasonable cosmological constraints from the very detailed analysis of the lensing configuration of the cluster of galaxies A2218. The necessary conditions for this study are {\\it simple}: deep multicolor {\\em HST} images of a well selected cluster-lens, identification of a minimum of 4 families of multiple-images systems and secure redshift measurement of each family, which ranges from $z=$0.702 to 5.576\\,. With these constraints, and provided the mass distribution can be modeled by the sum of a dominant component and smaller additional ones (all following a truncated PIEMD mass profile), the geometrical problem can then be solved. The cosmological constraints presented in this paper are of similar accuracy than those derived from Supernovae analysis. Interestingly, both analysis are purely geometrical and completely independent tests, but they are not sensitive to the same combination of distances thus providing nearly orthogonal constraints in the $(\\Omega_{\\rm M}, \\Omega_\\Lambda)$ plane.  This new method to constrain the cosmological parameters is very attractive, especially in view of the outstanding performances of the Advanced Camera for Surveys (ACS) on board of {\\em HST} and the development of Integral Field spectrography that allow to secure the redshift of many multiple images in a very efficient way. The very spectacular ACS images presented by \\citet{benitez02} on Abell 1689 show that in a very near future, we can use the proposed method as a very serious cosmological test, by focusing on those clusters with more than 4 multiple images with spectroscopic redshift. One advantage of this method is its relatively low-cost in terms of telescope time and relatively easy to implement - although progress is needed to thoroughly explore the parameter space of the mass models and to implement a fully comprehensive likelihood analysis. Such improvements are currently under investigation and will in a near future allow a better treatment of this exciting problem."
    },
    "0402/astro-ph0402372_arXiv.txt": {
        "abstract": "We present here new results on circumstellar nebulosity around SU Aurigae, a T-Tauri star of about 2 solar mass and 5 Myrs old at 152 pc in the J, H and K bands using high resolution adaptive optics imaging (0$\\farcs$30) with the Penn state IR Imaging Spectrograph (PIRIS) at the 100 inch Mt. Wilson telescope.  A comparison with HST STIS optical (0.2 to 1.1 micron) images shows that the orientation of the circumstellar nebulosity in the near-IR extends from PAs 210 to 270 degrees in H and K bands and up to 300 degrees in the J band. We call the circumstellar nebulosity seen between 210 to 270 degrees as 'IR nebulosity'. We find that the IR nebulosity (which extends up to 3.5 arcsecs in J band and 2.5 arcsecs in the K band) is due to scattered light from the central star. The IR nebulosity is either a cavity formed by the stellar outflows or part of the circumstellar disk. We present a schematic 3-dimensional geometrical model of the disk and jet of SU Aur based on STIS and our near-IR observations. According to this model the IR nebulosity is a part of the circumstellar disk seen at high inclination angles. The extension of the IR nebulosity is consistent with estimates of the disk diameter of 50 to 400 AU in radius, from earlier mm, K band interferometric observations and SED fittings. ",
        "introduction": "SU Aurigae (SU Aur) is a T-Tauri star located in the Taurus-Aurigae complex of dark molecular clouds at a distance of 152 pc (de Warf et al. 2003, Hipparcos catalog). Its spectral type is G2III, mass $\\sim$ 2$M_\\odot$ and age is about 4 to 5 million years (de Warf et al. 1998, de Warf et al. 2003). Recent observations by Nadalin et al. (2000) have shown short time variability in the B band magnitude that they attribute to proto-planetary materials orbiting in the circumstellar disk. Petrov et al. (1996) have found existence of a gas outflow from the star from their spectroscopic results.  The HST/STIS coronagraphic observations by Grady et al. (2002) from 0.2 to 1.0 microns have revealed fan like structures extending up to a distance of 12-15 arcsecs in the west to south-west direction. These, according to the authors, are mainly reflection nebulae scattering the light of the central star. In addition they also detected streaming filamentary structures going radially outwards from the star which could be due to gas outflow either from the star or from the parent molecular cloud.  Our motivation comes from the STIS images of SU Aur, and from the fact that no report of high spatial resolution near-infrared images of the star could be found in the literature to this date. Therefore, we observed the source in the near-infrared (J, H and K band images) using adaptive optics and applied the technique of PSF subtraction to investigate the circumstellar region of SU Aur with high spatial resolution (0$\\farcs$25 to 0$\\farcs$30). In this paper we present the results of our investigation (of SU Aur). Section 2 describes observations and data analysis procedure, section 3 introduces the results, in section 4 we discuss our results and present our conclusions in section 5. ",
        "conclusions": "We have presented here PSF subtracted high spatial resolution images (0$\\farcs$30 arcsecs) of the circumstellar region of SU Aurigae (between 1 to 4 arcsecs) in the J,H and K bands. These images show a distinct region bright in the near-IR which we call the IR nebulosity. The IR nebulosity is prominent between position angles 210 degrees to 270 degrees and extends up to 3.5 arcsecs in the J band and 2.5 arcsecs in the K band. We present two scenarios about the nature of the IR nebulosity: a) that it could be a cavity formed by the stellar outflows or b) that it is a part of the circumstellar disk observed at high inclination angles ($\\ge$65 degrees). We favor the later case more and we present a schematic 3-dimensional geometrical model describing the orientation of the disk and jet of SU Aur with respect to the observer. However, more observations are necessary with a large telescope equipped with AO and perhaps with coronagraphs and polarimetry to resolve spatial structures close to the star (down to 0$\\farcs$1 arcsecs) for detailed modeling."
    },
    "0402/astro-ph0402414_arXiv.txt": {
        "abstract": "{We report results of a serendipitous hard X-ray (3--20 keV), nearly all-sky ($|b|>10^\\circ$) survey based on RXTE/PCA observations performed during satellite reorientations in 1996--2002. The survey is 80\\% (90\\%) complete to a 4$\\sigma$ limiting flux of $\\approx 1.8$ (2.5) $\\times 10^{-11}$ erg s$^{-1}$ cm$^{-2}$ in the 3--20~keV band. The achieved sensitivity in the 3--8~keV and 8--20~keV subbands is similar to and an order of magnitude higher than that of the previously record HEAO-1 A1 and HEAO-1 A4 all-sky surveys, respectively. A combined $7\\times 10^3$~sq.~deg area of the sky is sampled to flux levels below $10^{-11}$ erg s$^{-1}$ cm$^{-2}$ (3--20 keV). In total 294 sources are detected and localized to better than 1~deg. 236 (80\\%) of these can be confidently associated with a known astrophysical object; another 22 likely result from the superposition of 2 or 3 closely located known sources. 35 detected sources remain unidentified, although for 12 of these we report a likely soft X-ray counterpart from the ROSAT all-sky survey bright source catalog. Of the reliably identified sources, 63 have local origin (Milky Way, LMC or SMC), 64 are clusters of galaxies and 100 are active galactic nuclei (AGN). The fact that the unidentified X-ray sources have hard spectra suggests that the majority of them are AGN, including highly obscured ones ($N_{\\rm H}>10^{23}$~cm$^{-2}$). For the first time we present a $\\log N$--$\\log S$ diagram for extragalactic sources above $4\\times 10^{-12}$ erg s$^{-1}$ cm$^{-2}$ at 8-20 keV. ",
        "introduction": "The deep surveys performed recently in the standard X-ray band (2--10~keV) with the Chandra and XMM-Newton observatories (e.g. \\cite{chandra}, \\cite{xmm}) have convincingly proved the extragalactic origin of the cosmic X-ray background (CXB). Hundreds of point sources detected in these surveys provide us with a wealth of information about the distant Universe. However, due to the very small sky coverage of these surveys, they are practically unsuitable for the study of the local Universe ($z\\la 0.3$). Medium-sensitivity ($10^{-13}$--$10^{-12}$~erg~cm$^{-2}$~s$^{-1}$) X-ray surveys, such as those performed with ASCA (\\cite{gis1}) and BeppoSAX (\\cite{giommi2000}), cover larger areas of the sky ($\\la 10^2$~sq. deg) but also cannot sample efficiently the Universe within $\\sim 500$~Mpc of us. In this regard, the results of the soft X-ray ($<2$~keV) all-sky survey carried out with the ROSAT observatory (e.g. \\cite{rbsc}) are extremely important, but these cannot be directly extrapolated into the $>2$~keV energy band. Therefore, our knowledge of the statistical properties of the local population of hard X-ray sources still rests largely on the snapshot of the whole sky taken in the 2--100~keV energy band more than 20 years ago by the different experiments on board the HEAO-1 observatory, A1 (\\cite{a1}), A2 (\\cite{a2}) and A4 (\\cite{a4}). It is only now that we have the possibility to undertake a new hard X-ray (3--20~keV) all-sky survey at similar (below 10~keV) and much better (above 10~keV) sensitivity provided by the RXTE observatory.  The Rossi X-ray Timing Explorer (RXTE, \\cite{rxte}) was launched at the end of 1995 and has now been successfully operating for more than 7 years. The mission was primarily designed to study the variability of X-ray sources on time scales from sub-milliseconds to years (e.g. \\cite{swank_ns}). The maneuvering capability of the satellite combined with the high photon throughput of its main detector (PCA) has also made it possible to carry out a series of Galactic Bulge scans aimed at the detection of new transient sources and following the long-term behavior of known X-ray sources (\\cite{craig00}). In addition, over its still continuing life time, RXTE/PCA has collected a large amount of data of slew observations covering almost the entire sky. In the current work we use these data to perform an all-sky survey in the 3--20 keV energy band. The relatively narrow field of view of the PCA instrument allows us to localize sources to better than 1 deg and thus effectively avoid source confusion after we restrict our consideration to Galactic latitudes $|b|>10^\\circ$.  We note that the RXTE/PCA slew data have previously been utilized to reconstruct the average spectrum of the CXB (\\cite{cxbpaper}). ",
        "conclusions": ""
    },
    "0402/astro-ph0402424.txt": {
        "abstract": "{ Photometric data for 593 Cepheids in the LMC, measured by Udalski et al. in the OGLE survey, augmented by 97 longer period Cepheids from other sources, are analyzed for the period- color (P-C) and period-luminosity (P-L) relations, and for the variations of amplitude, light curve shape, and period across the instability strip at constant absolute magnitude. Both the P-C and P-L relations have different slopes for periods smaller and larger than 10 days. The break at 10 days is also seen in the period-amplitude relations, and the compound Fourier combinations of $R_{21}$ and $\\Phi_{21}$ introduced by Simon and Lee.  The LMC Cepheids are bluer than Galactic Cepheids in the $B$,$V$, and $I$ color bands, part of which is due to differential Fraunhofer line blanketing and part to real differences in the temperature boundaries of the instability strip. The LMC strip is hotter by between $80\\;$K and $350\\;$K depending on the period. Hence, both the slopes and (necessarily) the zero points of the P-L relations in $B$, $V$, and $I$ must differ between LMC and the revised relations (also given here) for the Galaxy, and in fact they do. The LMC Cepheids are brighter by up to $0.5\\;$mag at $\\log P = 0.4$ (2 days) and fainter by $0.2\\;$mag at $\\log P = 1.5$ (32 days). These facts complicate the use of Cepheid as precision distance indicators until the reason is found (metallicity differences or other unknown differences) for the non-universality of the P-L and P-C relations.  The very large data base permits mapping of various Cepheid properties at different positions within the instability strip, both at constant period and at constant absolute magnitude over the range of $2 < P < 40$ days and $-2 > M_{V} > -5$. Amplitude of the light curves are largest near the blue edge of the strip for periods between 2 and 7 days and longer than 15 days. The sense is reversed for periods between 7 and 15 days. The shape of the light curves varies systematically across the strip. Highly peaked curves (of large amplitude) that necessarily have large values of $R_{21}$ of about $0.5$, occur near the blue edge of the strip. More symmetrical (small amplitude) light curves that have, thereby, small values of $R_{21}$, generally occur near the red edge of the strip. Consequently, there is a strong correlation of $R_{21}$ with color within the strip at a given absolute magnitude. Strong correlations also exist between color and period, and color and amplitude at given absolute magnitudes, for the same reason that has long been known for RR Lyrae stars, based on the sloping lines of constant period in the CMD combined with the variation of amplitude, and now $R_{21}$, with color. The highly peaked light-curve shapes and large amplitudes (indicating a non-linear regime that is overdriven out of the linear regime) near the blue edge of the strip, show that the energy~driver for the pulsation (i.e.\\ the negative dissipation) is strongest at the blue edge. % ",
        "introduction": "% We showed in Paper~I \\citep{Tammann:etal:03} that the period-luminosity (P-L) relation for Cepheids in the Galaxy, LMC, and SMC have significantly different slopes. If true, by necessity the slopes of the ridge-line (mean) period-color (P-C) relations in the three galaxies must also differ. We saw that this requirement was met, based on the extensive new photometric data by \\citet{Udalski:etal:99b,Udalski:etal:99c} for LMC and SMC (\\citeauthor*{Tammann:etal:03}, Fig. 16; and Sect.~\\ref{sec:PL:galaxy:comparison} here) from the OGLE project and that of \\citet{Berdnikov:etal:00} for the Galaxy.  The problem posed by this result is severe. With the calibrated Galactic P-L relation (as revised in Sect.~\\ref{sec:PL:galaxy:revised} here) fixed by two independent methods (main sequence fittings and the Baade-Becker-Wesselink (BBW) kinematic expansion method), and in the LMC by its distance determined by a variety of non-Cepheid methods (\\citeauthor*{Tammann:etal:03}, Table~6), our LMC relation is {\\em brighter\\/} than that for the Galaxy by $-0\\fm37$ in $V$ at $\\log P = 0.4$ ($P = 2.5^{d}$), and {\\em fainter\\/} by $+0\\fm11$ at $\\log P = 1.6$ ($P = 40^{d}$). The differences in $I$ at the same periods are $-0\\fm32$ and $+0\\fm19$. The LMC and Galactic P-L relations cross at $\\log P \\approx 1.3$ ($P = 20^{d}$) {\\em if\\/} our adopted distance modulus of LMC at $(m-M)^{0} = 18.54$ is correct.  A principal purpose for this series, of which this is the second paper of a projected four paper brief, is to analyze in tedious detail the new LMC and SMC data obtained by the OGLE consortium and to compare them with the Galaxy. Our purpose is to ferret out clues for the cause of the differences in the P-L and P-C relations and of the different position in the luminosity-temperature diagram of the Cepheids in the three galaxies. In the present discussion here we follow the methods of \\citeauthor*{Tammann:etal:03}.  The prime suspects for at least part of the variations from galaxy to galaxy are the established metallicity differences, for the reasons suggested in \\citeauthor*{Tammann:etal:03}. At the most elementary level, the purely technical effects on the colors of the differential blanketing effects for different metallicities affect the color-color relations, and color-period relations. And, as mentioned above, once there is a difference in the slopes of the color-period relations for whatever series of reasons, the P-L relations must also differ.  We saw in \\citeauthor*{Tammann:etal:03} that real temperature differences exist in the ridge-line $L_{\\rm bol}$, $\\log T_{\\rm e}$ plane (the HR diagram) between the Galaxy and LMC Cepheids. Is this also due to the effect of metallicity differences in the much more complicated physics of the pulsation and the position of the instability strip? If so, this, of course, would be a deeper reason than the simple technical effect of Fraunhofer blanketing. One of the purposes of this paper is to extend the discussion of these points from \\citeauthor*{Tammann:etal:03} here in Sect.~\\ref{sec:InstabilityStrip:MTeff}.  A second purpose is to explore the properties of the instability strip by the same technique made manifest by the horizontal branch of globular clusters as it threads the instability strip nearly at constant luminosity. The large data base from the new precision photometric data of \\citet{Udalski:etal:99b,Udalski:etal:99c} for the LMC Cepheids permit a variety of studies of the Cepheid properties as a function of position (color) in the strip, similar to such studies for the RR Lyrae stars.  % Section 2 Section~\\ref{sec:Data} discusses the OGLE data of \\citet{Udalski:etal:99b,Udalski:etal:99c} for LMC Cepheids, defining the sample, the Cepheid reddenings, and the errors. % Section 3 The period-color relations in $(B\\!-\\!V)^{0}$ and $(V\\!-\\!I)^{0}$, with comparisons with the Galaxy, are in Sect.~\\ref{sec:PC}. % Section 4 The P-L relations for LMC in $B$, $V$, and $I$, with comparisons with an updated calibration of the Galactic Cepheids are in Sect.~\\ref{sec:PL}. % Section 5 Discussion of the break in the P-L and P-C relations at $\\log P = 1.0$, and the abnormal behavior of the Fourier components and amplitudes of the light-curves, also at 10 days, is in Sect.~\\ref{sec:break}. % Section 6 Section~\\ref{sec:InstabilityStrip} is an extended discussion of the position of the instability strip in the HR diagram ($M_{V}$-color plane). Comparison with the Galaxy (Fig.~15 of \\citeauthor*{Tammann:etal:03}) is made there. The break in the slope of the $M_{V}$-color relations in $(B\\!-\\!V)^{0}$ and $(V\\!-\\!I)^{0}$ at 10 days period is manifest. Also in that section, the slopes of the lines of constant period in the color-magnitude diagram (CMD) are derived in different period ranges, as are the correlations of amplitude with position (color) in the strip.  % Section 7 Section~\\ref{sec:PLC} shows the period-luminosity-color (PLC) relations in $B$, $V$, and $I$, based on these aforementioned properties of the LMC Cepheids. % Section 8 The instability strip in the luminosity-temperature ($\\log L\\!-\\!\\log T_{\\rm e}$) plane is derived in Sect.~\\ref{sec:InstabilityStrip:MTeff} where the effects of different metallicities and helium abundances, based on models in the literature, are discussed. % Section 9 Section~\\ref{sec:conclusion} is a summary, and states the dilemma caused by these results in the using of Cepheids as precision distance indicators in our program to calibrate the Hubble constant by means of Cepheid distances to galaxies that have produced type Ia supernovae \\citep{Parodi:etal:00,Saha:etal:01}.   % ****************************************************************** % 2. The Data for LMC Cepheids % ****************************************************************** ",
        "conclusions": "\\label{sec:conclusion} % The consequences of the differences in the slopes of the P-L relations for the Galaxy, LMC, and SMC weakens the hope of using Cepheids to obtain precision galaxy distances. Until we understand the reasons for the differences in the P-L relations and the shifts in the period-color relations, (after applying blanketing corrections for metallicity differences), we are presently at a loss to choose which of the several P-L relations to use (Galaxy, LMC, and SMC) for other galaxies.  Although we can still hope that the differences may yet be caused only by variations in metallicity, which can be measured, this can only be decided by future research such as survey programs to determine the properties of Cepheids in galaxies such as M\\,33 and M\\,101 where metallicity gradients exist across the image. But until we can prove or disprove that metallicity difference is the key parameter, we must provisionally {\\em assume\\/} that this is the case, and use either the Galaxy P-L relations in Sect.~\\ref{sec:PL:galaxy:revised} (Eqs.~\\ref{eq:gal:PL:B}$-$\\ref{eq:gal:PL:I}), or the LMC P-L relations in Sect.~\\ref{sec:PL:LMC} (Eqs.~\\ref{eq:PL:B}$-$\\ref{eq:PL:I:ge1}), or those in the SMC from the forthcoming Paper~III, in deriving galaxy distances from Cepheids. This has now in fact been done elsewhere for M\\,83 \\citep{Thim:etal:03}.  There are eleven principal research points in this paper.  \\noindent (1) The period-color (P-C) relations in $(B\\!-\\!V)^{0}$ and $(V\\!-\\!I)^{0}$ have significantly different slopes for periods smaller and larger than 10 days (Figs.~\\ref{fig:PC:BV}a and \\ref{fig:PC:VI}b). The slope of the longer period Cepheids is steeper than those with $P < 10$ days at the significance level of $3.4\\sigma$ in $(B\\!-\\!V)^{0}$ and $4.1\\sigma$ in $(V\\!-\\!I)^{0}$.  \\noindent (2) The LMC Cepheids are bluer than Galactic Cepheids by $0\\fm07$ in $(B\\!-\\!V)^{0}$ at $\\log P=0.4$, increasing to $0\\fm12$ at $\\log P=1$. The corresponding differences in $(V\\!-\\!I)^{0}$ are $0\\fm03$ and $0\\fm10$ (Eqs.~\\ref{eq:PC:BV:lt1}$-$\\ref{eq:PC:VI:ge1} for LMC compared with Eqs.~3 and 5 of \\citeauthor*{Tammann:etal:03}). Blanketing differences due to the lower metallicity of LMC accounts for only half of the observed color differences. A temperature difference (Figs.~\\ref{fig:logPTe} and \\ref{fig:LMC:Teff}) accounts for the remainder, with the LMC Cepheids being hotter by between $350\\;$K and $80\\;$K, depending on the period. Theoretical models of the position of the fundamental blue edge as functions of metallicity and helium abundance cannot account for the temperature difference (Eq.~\\ref{eq:logTe:YZ}) if $Z$ and $Y$ increase in lock step. An explanation is possible if $Y$ and $Z$ are anticorrelated (Sect.~\\ref{sec:InstabilityStrip:MTeff}), but this is highly counterintuitive based on the models of nucleosynthesis.  \\noindent (3) The P-L relations for LMC in $B$, $V$, and $I$ are all non-linear, each with a break at $\\log P = 1$ (Eqs.~\\ref{eq:PL:B:lt1}$-$\\ref{eq:PL:I:ge1}). The slope differences are significant at the at least $3\\sigma$ level. The reality of the break in the P-L relations at 10 days is supported by the breaks in the P-C relations (item 1). Necessarily, breaks in the P-C relations require breaks in the P-L relations.  \\noindent (4) The slopes of the LMC P-L relations in the $B$, $V$, and $I$ pass bands are different  from the slopes of the corresponding relations in the Galaxy as updated in \\ref{sec:PL:galaxy:revised}. For Cepheids with $\\log P<1$ the P-L slope differences between LMC and the Galaxy vanish in $B$, but amount to $2\\sigma$ and $3\\sigma$ in $V$ and $I$. The significance of the slope differences is even larger for long-period Cepheids with $\\log P>1$, i.e.\\ it increases from $3\\sigma$ in $B$ to $5\\sigma$ in $V$ and $I$. Since most available galaxy distances from Cepheids are based on long-period Cepheids ($\\log P_{\\rm median}\\!\\sim\\!1.4$) the choice of which P-L relation is used has a non-negligible effect on the derived absorption-corrected moduli by about $0\\fm2$ (cf. Eq.~40 in \\citeauthor*{Tammann:etal:03}) or $\\sim\\!10\\%$ in distance.  \\noindent (5) The absolute magnitudes of LMC Cepheids are {\\em brighter\\/} than those in the Galaxy by $0\\fm42$ to $0\\fm32$ in $B$, $V$, and $I$ at $\\log P=0.4$, becoming {\\em fainter\\/} by $0\\fm06$ in $V$ and by $0\\fm14$ in $I$ at $\\log P=1.5$ (Eqs.~\\ref{eq:PL:B:lt1}$-$\\ref{eq:PL:I:ge1} compared with \\ref{eq:gal:PL:B}$-$\\ref{eq:gal:PL:I}). All attempts to determine a Cepheid distance of LMC by means of a Galactic P-L relation \\citep[e.g.][]{Fouque:etal:03,Groenewegen:Salaris:03,Storm:etal:04} are therefore frustrated. The resulting values of $(m-M)^{0}_{\\rm LMC}$ would vary between $\\sim\\!18.16$ and $\\sim\\!18.63$ (still somewhat dependent on pass band) depending on whether Cepheids with $\\log P=0.4$ or $\\log P=1.5$ are compared. (The {\\em adopted\\/} LMC distance of $18.54$ in the present paper rests on a compilation of distance determinations which are {\\em independent\\/} of the P-L relation of Cepheids; see \\citeauthor*{Tammann:etal:03}, Table~6). -- There are now, however, good prospects to re-establish Cepheids as important distance indicators of LMC by means of their individual BBW distances, which depend little on metallicity. Work towards this aim is in progress (W.~Gieren, private communication).  \\noindent (6) The reality of the break in Cepheid properties at 10 days is shown in various correlations. (a) The Fourier components $R_{21}$ and $\\Phi_{21}$, introduced by \\citet{Simon:Lee:81}, show discontinuities at 10 days (Fig.~\\ref{fig:break:Rphi}). (b) A similar discontinuity exists in the period-amplitude relation, both for LMC and the Galaxy (Fig.~\\ref{fig:break:Vamp}). (c) The slope of the ridge-line of the color-magnitude diagram for the instability strip (Fig.~\\ref{fig:InstabilityStrip}) changes at 10 days, as in the theoretical HR ($L/T_{\\rm e}$) diagram (Fig.~\\ref{fig:LMC:Teff}). (d) Changes at 10 days also occur in other correlations such as the slopes of the lines of constant period (Fig.~\\ref{fig:alpha} and Sect.~\\ref{sec:InstabilityStrip:CPL}), the character of the color-amplitude correlation in the period range of 10 to 20 days (Figs.~\\ref{fig:Bamp} and \\ref{fig:BampBV}), and in the $R_{21}$ vs. $\\log P$ relation with the data binned by absolute magnitude (Fig.~\\ref{fig:logPRphi21}).  \\noindent (7) The loci of constant period in the CMD (Fig.~\\ref{fig:InstabilityStrip}) slope toward fainter magnitudes as the color changes across the strip from blue to red. Hence, at a given period there is a variation in absolute magnitude causing the systematic scatter in the P-L relations in $B$, $V$, and $I$. The mapping of this effect is particularly strong here using the LMC data because the data base over the entire period range from 2 to 40 days is superbly large from the OGLE project. The slopes of the lines of constant period in the $M_{V}$ P-L relation average $1.8$ in $(B\\!-\\!V)^{0}$ (but vary with period) and $2.4$ in $(V\\!-\\!I)^{0}$ (Fig.~\\ref{fig:CPL} and Eqs.~\\ref{eq:CPL:BV:lt1}$-$\\ref{eq:CPL:VI}). The slope of the constant-period lines is of course steeper than this in the $M_{B}-(B\\!-\\!V)$ plane and more shallow in the $M_{I}-(V\\!-\\!I)$ plane (Eqs.~\\ref{eq:alpha:B:BV} and \\ref{eq:alpha:I:VI}), showing why the scatter in the P-L relations becomes progressive smaller from $B$, through $V$, to $I$, as was originally shown \\citep{Sandage:58,Sandage:72} as due to the finite color width of the instability strip.  \\noindent (8) The large data base permits a high-weight mapping of several Cepheid characteristics with position in the strip. The data binned in narrow period intervals from 3 to 40 days, (i.e.\\ along lines of constant period in the CMD), show (Fig.~\\ref{fig:Bamp}) that, except near 10 days, the amplitude is largest at the blue edge of the strip, just as for the RR Lyrae stars in the low mass instability strip peculiar to them. The analogy with the RR Lyrae stars is stronger when the data are binned in narrow intervals of absolute magnitude (Fig.~\\ref{fig:BampBV}) similar to the case of RR Lyrae variables in globular clusters. Fig.~\\ref{fig:BampBV} shows that the correlation of light-curve amplitude with color (at the relevant absolute magnitude from $-2.75$ to brighter than $-4.75$) is in the sense that largest amplitude occurs at the bluest color for Cepheids fainter than $M_{V}= -4.25$ (periods smaller than 7 days), the trend reverses for periods of 10 to 15 days ($-4.25 < M_{V} < -4.75$) and reverts to the original sense for $P > 20$ days as was earlier found by \\citet{Sandage:Tammann:71} from a much smaller data sample.  \\noindent (9) The Fourier coefficient $R_{12}$ varies systematically with amplitude and in color across the strip as shown in Fig.~\\ref{fig:R21BampBV}, binned by period and in Fig.~\\ref{fig:BVRphi21}, in analogy with the RR Lyrae correlations, binned by absolute magnitude. Large $R_{21}$ means that the second term of the Fourier series, $\\sin 2x$, is large, showing that the light curve is highly peaked (approaching a saw tooth for $R_{21} = 0.5$; Sect.~\\ref{sec:InstabilityStrip:Rphi21}). Small $R_{21}$ stands for nearly symmetrical light curves, which, by Fig.~\\ref{fig:R21BampBV}, occur at small amplitude in all period ranges.  Hence, Figs.~\\ref{fig:R21BampBV} and \\ref{fig:BVRphi21} show that the light curve shapes are highly peaked toward the blue edge of the instability strip, becoming nearly sinusoidal (and of small amplitude) toward the red. {\\em $R_{12}$ varies systematically across the strip}. This is an important result concerning the physics of pulsation. Large amplitude with its near saw-tooth shape (small interval in phase between minimum and maximum light) means that the pulsation is more strongly driven into the non-linear regime at the blue edge than toward the red edge. Hence, the prediction is that the PV work diagram in the pulsation thermodynamics must have a larger non-dissipative area for Cepheids near the blue edge than near the red edge \\citep[cf.][]{Christy:66}.  \\noindent (10) By necessity, from these correlations of $R_{21}$ with color (Figs.~\\ref{fig:R21BampBV} and \\ref{fig:BVRphi21}), and from the correlations of color and period, (Fig.~\\ref{fig:logPBVI}) there must be a correlation of $R_{21}$ with period at constant absolute magnitude. This is confirmed in Fig.~\\ref{fig:logPRphi21}. The reason is obvious (by inspecting the CMD of Fig.~\\ref{fig:InstabilityStrip}), once it is known that $R_{21}$ varies systematically with color across the strip.  \\noindent (11) Application of the PLC relation to the entirety of the OGLE LMC data gives the {\\em relative\\/} distance to each individual LMC Cepheid. This permits determination of the orientation of the bar in the line of sight. The eastern edge of the bar is only $0\\fm03$ closer than the western edge (Fig.~\\ref{fig:LMCorientation}), confirming earlier conclusions by \\citet{Welch:etal:87} and \\citet{Caldwell:Laney:91}.   Finally we reiterate that the unexpected slope difference between the P-L relations particularly of long-period Cepheids in the Galaxy and LMC can be understood as the consequence of the different gradients in the luminosity-temperature diagram, yet the reason {\\em why\\/} these gradients are different remains unknown.    % ****************************************************************** %  Acknowledgments % ******************************************************************"
    },
    "0402/astro-ph0402622_arXiv.txt": {
        "abstract": "\\noindent{\\bf A luminous X-ray source is associated with a cluster (MGG-11) of young stars $\\sim 200$\\,pc from the center of the starburst galaxy M82 \\cite{Matsumoto1999,Kaaret2001}. The properties of the X-ray source are best explained by a black hole with a mass of at least 350\\,\\msun\\, \\cite{Matsumoto2001,Strohmayer2003}, which is intermediate between stellar-mass and supermassive black holes.  A nearby but somewhat more massive star cluster (MGG-9) shows no evidence of such an intermediate mass black hole \\cite{Matsumoto1999,Matsumoto2001}, raising the issue of just what physical characteristics of the clusters can account for this difference. Here we report numerical simulations of the evolution and the motions of stars within the clusters, where stars are allowed to mergers with each other. We find that for MGG-11 dynamical friction leads to the massive stars sinking rapidly to the center of the cluster to participate in a runaway collision, thereby producing a star of 800--3000\\,\\msun, which ultimately collapses to an black hole of intermediate mass. No such runaway occurs in the cluster MGG-9 because the larger cluster radius leads to a mass-segregation timescale a factor of five longer than for MGG-11.  } ",
        "introduction": " ",
        "conclusions": ""
    },
    "0402/astro-ph0402308_arXiv.txt": {
        "abstract": "We explore the near-infrared (NIR) $K$-band properties of galaxies within 93 galaxy clusters and groups using data from the Two Micron All Sky Survey (2MASS).  We use X-ray properties of these clusters to pinpoint cluster centers and estimate cluster masses. By stacking all these systems, we study the shape of the cluster luminosity function and the galaxy distribution within the clusters. We find that the galaxy profile is well described by the NFW profile with a concentration parameter $c \\sim 3$, with no evidence for cluster mass dependence of the concentration. Using this sample, whose masses span the range from $3\\times10^{13}M_\\odot$ to $2\\times10^{15}M_\\odot$, we confirm the existence of a tight correlation between total galaxy NIR luminosity and cluster binding mass, which indicates that NIR light can serve as a cluster mass indicator. From the observed galaxy profile, together with cluster mass profile measurements from the literature, we find that the \\mlr is a weakly decreasing function of cluster radius, and that it increases with cluster mass. We also derive the mean number of galaxies within halos of a given mass, the halo occupation number. We find that the mean number scales as $N\\propto M^{0.84\\pm0.04}$ for galaxies brighter than $M_K=-21$, indicating high mass clusters have fewer galaxies per unit mass than low mass clusters.  Using published observations at high redshift, we show that higher redshift clusters have higher mean occupation number than nearby systems of the same mass.  By comparing the luminosity function and radial distribution of galaxies in low mass and high mass clusters, we show that there is a marked decrease in the number density of galaxies fainter than $M_*$ as one moves to higher mass clusters; in addition, extremely luminous galaxies are more probable in high mass clusters. We explore several processes-- including tidal interactions and merging-- as a way of explaining the variation in galaxy population with cluster mass. ",
        "introduction": "\\label{sec:intro}   Understanding galaxy formation is one of the most outstanding challenges in cosmology. The development of both semianalytic \\citep[e.g.][among others]{swhite78,kauffmann93,cole00} and numerical \\citep[e.g.][]{kauffmann99,springel01b} modeling have enjoyed tremendous success, in the sense that those models are able to match several observed properties, such as the galaxy luminosity function (LF), the morphological mix, the color, the mass-to-light ratio of the galaxies \\citep{cole00}. But as the quality of the observational constraints improves, we can expect additional theoretical challenges \\citep[e.g.][]{benson03}.    The clustering properties of dark matter and galaxies pose another tough task, for both theorists and observers.  The recent advent of the so-called halo model has introduced an important new tool on this subject \\citep[e.g.][]{seljak00,peacock00,scoccimarro01}.  The essential ingredients of this model include a description of halo abundance as a function of cosmic time and halo mass (the mass function, e.g. \\citealt{press74,sheth99, jenkins01}), a model for halo structure \\citep[e.g.][]{navarro97,moore98}, and a prescription for the bias of haloes \\citep[e.g.][]{mo96,sheth99}.  The halo occupation distribution (HOD) is a powerful tool which links the physics of galaxy formation with the clustering of dark matter and galaxies \\citep{benson00,kauffmann99,berlind02,berlind03,kravtsov03}. It assumes that the evolution and clustering of haloes are determined by the underlying cosmology, and that the physics that governs galaxy formation specifies the way galaxies populate the haloes.  The calculations of various power spectra of dark matter and galaxies are a natural outcome.  The important ingredients within the HOD framework are the mean number of galaxies per halo $N$ as a function of halo mass, the probability distribution that a halo of mass $M$ contains $N$ galaxies $P(N|M)$, and the relative distribution (both spatial and velocity) of galaxies and dark matter within haloes \\citep{berlind02}.  Here we aim to provide observational constraints on the HOD, based on our study of 93 clusters and groups using data from the Two Micron All-Sky Survey (2MASS, \\citealt{jarrett00}). Using X-ray properties of these clusters to define the cluster center and estimate the cluster binding mass, we determine the mean halo occupation number $N$ as a function of mass from $\\sim 3 \\times 10^{13} M_\\odot$ to $\\sim 2 \\times 10^{15} M_\\odot$ and also investigate the galaxy distribution and luminosity function within the clusters. We discuss the bearing of the $N$--$M$ relation on the hierarchical structure formation paradigm.  Part of our analysis is built on the tools that we develop in an earlier paper \\citep[][hereafter paper I]{lin03b}, where we examined the near-infrared (NIR) galaxy luminosity--cluster binding mass correlation ($L$--$M$ relation) for a sample of 27 nearby clusters, using the second release of the 2MASS data.  Here, for a much larger sample, we will study the galaxy distribution within the clusters and the faint-end shape of the cluster luminosity function, solve the $K_s$-band luminosity functions for individual clusters, and derive the total light and galaxy number within the virial radius as a function of cluster binding mass.  Our findings provide some constraints on cluster evolution scenarios.   In \\S\\ref{sec:analysis} we briefly describe the technique developed in paper I for examination of the cluster NIR $L$--$M$ relation. Next in \\S\\ref{sec:stack} we begin our analysis with two fundamental properties of galaxy clusters: the luminosity function (LF) and the spatial distribution of the member galaxies.  With constraints on galaxy distributions both in real space and in luminosity space derived directly from the data, we proceed to calculate the $L$--$M$ relation, point out the importance of the contributions from the brightest cluster galaxies (BCGs) (\\S\\ref{sec:lm}), and examine the mean halo occupation number (\\S\\ref{sec:hod}). We investigate the possible mechanisms that are responsible for the observed behavior of the halo occupation distribution in \\S\\ref{sec:further}.  Possible systematics that may affect our results are discussed in \\S\\ref{sec:system}. Finally, in \\S\\ref{sec:summary}, we summarize our results. The Appendix provides some further tests of the robustness of our analysis.  Throughout the paper we assume the density parameters for the matter and the cosmological constant to be $\\Omega_M = 0.3$, $\\Omega_\\Lambda = 0.7$, respectively, and the Hubble parameter to be $H_0=70\\,h_{70}$~km~s$^{-1}$~Mpc$^{-1}$. ",
        "conclusions": "\\label{sec:summary}   We have analyzed a sample of 93 galaxy clusters whose virial masses range from $3\\times 10^{13} h_{70}^{-1} M_\\odot$ to $1.7\\times 10^{15} h_{70}^{-1} M_\\odot$.   We have used the peak in X-ray surface brightness to define the cluster centers, and we have used the observed correlation between emission-weighted mean temperature and cluster virial mass \\citep{finoguenov01} along with published cluster redshifts to estimate the scale of the cluster virial region.  The following bulk properties of these systems are studied, with emphasis on any mass related trends: the composite cluster luminosity function, the galaxy distribution within the clusters, the correlation between total galaxy $K$-band luminosity and the cluster binding mass, and the halo occupation number.  Our analysis shows that the luminosity function (excluding the BCG) is well fit by a Schechter function with a faint-end slope between $-1.1 \\lesssim \\alpha \\lesssim -0.84$, and characteristic magnitude and number density between $-24.34 \\le M_* -5\\log h_{70} \\le -24.02$ and $3.01\\,h_{70}^3$ Mpc$^{-3} \\le \\phi_* \\le 4.43\\,h_{70}^3$ Mpc$^{-3}$.   A NIR survey of nearby clusters to a depth of $M_K \\sim -19$ would be useful in further constraining the faint-end slope. The composite galaxy distribution ($0.02 \\le r/r_{200} \\le 2$) is found to be an NFW profile with concentration $c_g = 2.9\\pm 0.2$. This is somewhat smaller than the observed dark matter  concentration, as well as that found in numerical studies. This implies that galaxies have a more extended distribution than that of the dark matter.  By solving the LF parameters $L_*$ and $\\phi_*$ for each cluster, we obtain the total galaxy luminosity, as well as total galaxy number (\\S\\ref{sec:lmhod}) above some limiting absolute magnitude $M_{K,low}=-21$. By comparing the galaxy NIR light--cluster mass correlation at different radii ($r_{500}$ \\& $r_{200}$) we find that the \\mlr decreases slightly with radius (as would be expected given the differing distributions of galaxy and dark matter). The $L$--$M$ correlation has a slope smaller than unity ($0.69\\pm 0.04$, see Eqn \\ref{eq:lm}, Table 1), which implies that cluster \\mlr is an increasing function of cluster mass. However, excluding the BCGs significantly affects the normalization and the slope of the $L$--$M$ correlation (but of course has only a minor effect on the halo occupation number). Without the BCGs, the total galaxy light roughly scales as the total galaxy number, which increases more rapidly with cluster mass (Eqn~\\ref{eq:hod}, slope is $0.84\\pm0.04$; Table 1). The slope of the $N$--$M$ correlation agrees well with the prediction of numerical simulations. In addition, we find the distribution of the halo occupation number about the mean is Poissonian (Figure~\\ref{fig:hod}).  This flatness of the $N$--$M$ relation in our local cluster sample indicates that higher mass clusters have fewer galaxies per unit mass than low mass clusters. We examine the occupation number of higher redshift clusters using published NIR and X-ray data; in crude agreement with numerical simulations, the higher redshift clusters have more galaxies per unit mass than nearby clusters. As massive clusters accrete lower mass systems, there must be  processes at work that alter the galaxy populations, either reducing galaxy masses or perhaps even destroying some galaxies. Through further examination of our local cluster sample we find that the radial distribution of galaxies appears to be independent of cluster mass; in addition the ultra bright end of the luminosity function is simply not present within low mass clusters and groups, and the number density of galaxies fainter than about $M_*$ is higher in low mass clusters than in high mass clusters. Thus, the transformation of the galaxy population from low mass to high mass clusters appears to be a building up of a few extremely luminous galaxies and the destruction or faintening of galaxies fainter than $M_*$. We note that there does not appear to be enough light in missing galaxies in high mass clusters to account for the increases in the bright end of the luminosity function.  We investigate possible mechanisms that shape the slope of the halo occupation $N$--$M$ correlation (\\S\\ref{sec:further}).  We find that none of the mechanisms we consider can by itself offer a satisfactory explanation of the galaxy population differences between high and low mass clusters.  Among these, tidal interactions (stripping, galaxy harassment) could be quite important,  but presumably merging must also play a role in the production of very luminous galaxies. Further theoretical progress together with additional observational constraints from extensions of studies like ours will presumably lead to a clearer picture of the process of galaxy evolution within the cluster environment.  One implication of the observed $N$--$M$ relation in both our local sample and the high-$z$ sample is that galaxies are destroyed or dramatically altered through various dynamical processes in the cluster environment. One thus expects that the amount of the diffuse intracluster light would exhibit some correlation with the cluster mass \\citep[e.g.][]{malumuth84}. In a companion paper (paper III), we will study the relationship between the BCG luminosity and the cluster mass, using it to limit the amount of the intracluster diffuse light. This will provide further insights into the cluster formation scenarios and the evolution of cluster galaxies.  An improved understanding of  the redshift evolution of the HOD is critical for those who would use deep, NIR or optical cluster surveys to study cosmology.  High yield cluster surveys must employ simple mass estimators such as the galaxy number or luminosity, because it would not be feasible to conduct detailed mass measurements using X-ray imaging spectroscopy, galaxy velocity dispersions or weak lensing. Our study provides a local calibration for these mass--observable relations. However, extracting cosmology from the observed cluster redshift distribution will require knowledge of how these relations evolve. Our comparison of the local sample to the published data on the high-$z$ clusters suggests that cluster virial regions are better represented by galaxy number and light at high redshift than they are locally. Depending on how the non-cluster galaxies are evolving, this could boost optical and NIR cluster finding at high redshift.  It remains to be seen whether the recently discussed cluster survey self-calibration approach \\citep{majumdar03b,hu03b} will work with the higher scatter galaxy-based mass indicators."
    },
    "0402/astro-ph0402552_arXiv.txt": {
        "abstract": "{ We report the results of a search for the double-peaked blue-skewed infall signature in the bright-rimmed cloud core SFO 11NE SMM1. Observations of the optically thick HCO$^{+}$ and optically thin H$^{13}$CO$^{+}$ J=3--2 lines reveal that there is indeed a characteristic double-peaked line profile, but skewed to the red rather than the blue. Modelling of the dust continuum emission and line profiles show that the motions within SFO 11NE SMM1 are consistent with a collapsing central core surrounded by an expanding outer envelope. We show that the collapse is occurring at a similar rate to that expected onto a single solar-mass protostar and is unlikely to represent the  large-scale collapse of gas onto the infrared cluster seen at the heart of SFO 11NE SMM1. The outer envelope is expanding at a much greater rate than that expected for a photoevaporated flow from the cloud surface. The modelled expansion is consistent with the bulk cloud re-expansion phase predicted by radiative-driven implosion models of cometary clouds. ",
        "introduction": "Spectroscopic evidence for collapse motions in low-mass star-forming regions has now become relatively widespread (e.g. Myers et al.~\\cite{myers00}; Evans \\cite{evans99}). The classic spectroscopic signature of infall or collapse is an asymmetric  blue-skewed double-peaked optically thick line profile with self-absorption at the systemic velocity (Myers et al.~\\cite{myers00}). The interpretation of this profile can be fraught with difficulty: unrelated clouds located along the line-of-sight can mimic the infall signature (Myers et al.~\\cite{myers00}), molecular outflows add extra complexity to the gas kinematics (Hogerheijde et al.~\\cite{hvdbvl98}), chemical effects such as depletion  may mask the infall signature completely (Rawlings \\& Yates \\cite{ry01}) and  rotation of the molecular core can completely reverse the blue-skewed profile to  the red (Zhou \\cite{zhou95}).  In order to conclusively demonstrate the presence of infall it is usually necessary to map the emission of both optically thick and thin lines across the molecular cloud core and compare their profiles to collapse models (e.g.~Gregersen et al.~\\cite{gemm00}; Lee, Myers \\& Tafalla \\cite{lmt01}).  Despite these difficulties, collapse motions have been identified in several protostellar and preprotostellar cores. The objects that have been mainly investigated so far are Class 0 \\& I protostars (Gregersen et al.~\\cite{gezc97}; Mardones et al.~\\cite{mmtwb97}) and preprotostellar cores  (Lee, Myers \\& Tafalla~\\cite{lmt99}, \\cite{lmt01}). Most of these objects are located in nearby relatively isolated low-mass star-forming regions such as the Perseus, Serpens, Taurus and Ophiuchus molecular clouds (Mardones et al.~\\cite{mmtwb97}). Few star-forming regions associated with the clustered star-forming mode (Lada, Strom \\& Myers \\cite{lsm99}) have been searched for the presence of infall. Two notable exceptions include recent searches for collapse motions of gas associated with young stellar clusters themselves (Williams \\& Myers \\cite{wm99}; Williams \\& Garland \\cite{wg02}) and for infall within bright-rimmed clouds associated with large HII regions (de Vries, Narayanan \\& Snell \\cite{dvns02}).  Both types of search report limited success with collapse motions identified in two young clusters (Cepheus A and NGC 2264) and one bright-rimmed cloud (SFO 18). It is the latter type of object that we will focus on in this paper. Bright-rimmed clouds are potential regions of triggered or induced star formation (Sugitani, Fukui \\& Ogura~\\cite{sfo}; Sugitani \\& Ogura \\cite{so94}), whereby the collapse of the molecular gas comprising the cloud and the ensuing star formation process is initiated by the action of an external trigger. In bright-rimmed clouds the external  trigger is most likely the photoionisation-induced shocks that are driven into the clouds by the ionisation of their outer layers by nearby OB stars (e.g.~Elmegreen \\cite{e91}). These shocks compress the molecular gas of the clouds, triggering the collapse of the cloud and forming a dense, possibly quasistatic core at the cloud centre (Bertoldi \\cite{bertoldi}; Lefloch \\& Lazareff \\cite{ll94}, \\cite{ll95}).  However, although de Vries, Narayanan \\& Snell (\\cite{dvns02}) report the detection of an infall signature towards the bright-rimmed cloud \\object{SFO 18}, it is the exception rather than the rule in their sample of bright-rimmed clouds. The overall blue excess of the bright-rimmed cloud sample is less than that of the Class 0 \\& I protostellar cores observed by Mardones et al.~(\\cite{mmtwb97}) and Gregersen et al.~(\\cite{gemm00}). This phenomenon is puzzling given the strong velocity gradients predicted by the Lefloch \\& Lazareff (\\cite{ll94}) radiative-driven implosion model. However the  small sample size of de Vries, Narayanan \\& Snell (\\cite{dvns02}) means that their result may not be statistically significant. It is also possible that a flattened temperature gradient across the bright-rimmed clouds, caused by their external heating, could mask the classic asymmetric infall signature (de Vries, Narayanan \\& Snell \\cite{dvns02}). A wider infall survey of a larger number of bright-rimmed clouds, coupled with more realistic radiative transfer modelling of their temperature and velocity gradients is required to address the issue of infall in these clouds. Such a study would also permit the detailed investigation of the kinematics of these clouds, which is important in the context of refining and expanding the existing radiative driven implosion models to take into account star formation processes  (e.g.~Bertoldi \\cite{bertoldi}; Lefloch \\& Lazareff \\cite{ll94}, \\cite{ll95}).  We are currently undertaking such a survey of a statistically significant sample of bright-rimmed clouds. As a fore-runner to the survey we present the results of a search for infall toward the bright-rimmed cloud \\object{SFO 11NE}. \\object{SFO 11NE} is located at the edge of the HII region IC 1848, approximately 4\\arcmin\\ NE of the bright-rimmed cloud SFO 11  (Thompson et al.~\\cite{thompson03a}; Sugitani, Ogura \\& Pickles \\cite{osp02}). SCUBA sub-mm continuum and JCMT CO observations reveal a dense core of molecular gas and dust located at the head of the cometary cloud (\\object SFO 11NE SMM1), whilst 2MASS images show  that the dense dust core harbours a cluster of candidate young stellar objects and protostars (Thompson et al.~\\cite{thompson03a}). The pressure balance and morphology of the cloud cloud suggest that the cloud is likely to be in the early collapse phase described by Lefloch \\& Lazareff (\\cite{ll94}) and is thus a good candidate in which to search for infall. ",
        "conclusions": "We observed the bright-rimmed cloud core SFO 11NE SMM1 in sub-millimetre wave lines of HCO$^{+}$ and H$^{13}$CO$^{+}$ to search for the characteristic optically thick blue-peaked infall signature (e.g. Myers et al.~\\cite{myers00}). The HCO$^{+}$ emission toward SFO 11NE SMM1 is indeed double-peaked, but with an excess toward the red end of the spectrum rather than the blue. The H$^{13}$CO$^{+}$ spectra show a single peak at the systemic velocity and this suggests that the double-peaked optically thick profile originates from a single gas core rather than a line-of-sight coincidence of two different velocity components.  The line data were modelled with the  radially symmetric radiative transfer model RATRAN (Hogerheijde \\& van der Tak \\cite{hvdt00}). SCUBA and IRAS HIRES data from Thompson et al.~(\\cite{thompson03a}) were used to construct a detailed physical model of the temperature and density structure of the cloud core against which the line emission could be modelled. The resulting physical model of the source is a cloud core of radius 0.26 pc, with a density at the outer radius of $4.5 \\times 10^{3}$ cm$^{-3}$ increasing toward the centre of the core as r$^{-3/2}$. The SED and radial flux profile of the core suggest that the temperature of the core consists of two components: an inner warm region at 30 K surrounded by a colder envelope of 18 K.  The HCO$^{+}$ and H$^{13}$CO$^{+}$ line profiles were found to be consistent with a combination of a central core collapsing at a constant velocity of 0.1 km\\,s$^{-1}$, surrounded by an expanding envelope moving outwards at a constant velocity of 0.3 km\\,s$^{-1}$.  We derive the mass infall and expansion rates to first order  by a simple consideration of the mass, velocity and radius of the collapsing and expanding regions. The mass infall rate is consistent with that estimated for a single solar-mass protostar (e.g.~Zhou \\cite{zhou95}) and it is possible that we are observing collapse onto a single member of the infrared cluster seen at the heart of SFO 11NE SMM1 (Thompson et al.~\\cite{thompson03a}).  The expansion rate is an order of magnitude larger than can be sustained by a photoevaporative flow from the outer layers of the cloud (Lefloch \\& Lazareff \\cite{ll94}). The expansion is consistent with either the bulk re-expansion of the cloud predicted by the RDI models of Lefloch \\& Lazareff (\\cite{ll94}) or multiple unresolved bipolar molecular outflows. Given both the lack of supportive evidence for outflows and the relatively high mass of expanding gas predicted by the radiative transfer modelling we conclude that the likeliest scenario for the expanding gas is a bulk re-expansion of the bright-rimmed cloud. If so, SFO 11NE SMM1 is at the beginning of its re-expansion phase following the maximum compression of the cloud by photoionisation-induced shocks.  Although our model of the cloud core as undergoing central collapse and outer expansion is consistent with the data, we stress that this is just one possible interpretation. Observations of higher sensitivity and  spatial resolution are required to investigate the kinematics of SFO 11NE SMM1 at smaller spatial scales and provide sufficient data for fully two-dimensional radiative transfer  modelling of the line profiles (e.g.~Hogerheijde \\& van der Tak \\cite{hvdt00}; Phillips \\& Little \\cite{pl00}) . In particular, the possibilities that the cloud core is rotating about its north-south axis and that the observed expansion may be in part due to bipolar molecular outflows must be investigated. Our forthcoming survey of a statistically significant number of BRCs will also reveal whether infall motions are as common within these clouds as more isolated low-mass star-forming regions."
    },
    "0402/astro-ph0402078_arXiv.txt": {
        "abstract": "{ We present new mid-infrared observations of objects in the vicinity of the O-star $\\sigma$\\,Orionis, obtained with TIMMI-2 at ESO. By constraining their near- and mid-infrared spectral energy distributions, we established the nature of previously known IRAS sources and identified new mid-infrared sources as young stellar objects with circumstellar disks, likely massive members of the $\\sigma$\\,Ori cluster. For two of these objects we have obtained spectroscopy in the 8--13\\,$\\mu$m range in order to investigate the chemistry of the dust grains. TX\\,Ori exhibits a typical silicate emission feature at 10\\,$\\mu$m, with a feature at about 11.2\\,$\\mu$m that we identify as due to crystalline olivine. The IRAS\\,05358$-$0238 spectrum is very unusual, with a weak silicate feature and structure in the range 10--12\\,$\\mu$m that may be explained as due to self-absorbed forsterite. We also provide the first evidence for the presence of circumstellar disks in the jet sources Haro\\,5-39/HH\\,447, V510\\,Ori/HH\\,444 and V603\\,Ori/HH\\,445. ",
        "introduction": "Disk-like structures seem to be ubiquitous during the formation and early evolution of low-mass stars. Disks are also the birthplace of planetary bodies. A particularly challenging problem is how to reconcile the relatively long timescales for planet formation ($>$\\,10\\,Myr, Bodenheimer et al.\\ \\cite{bodenheimer00}) and the rather quick destruction of circumstellar disks (e.g.\\ Haisch et al.\\ \\cite{haisch01a}). Amongst the many (yet) unanswered disk-related questions are: what are the timescales of disk dissipation; how does the chemical and physical evolution of dust proceed from small interstellar dust grains through pebble-sized particles to larger bodies? And how does the local physical environment (in particular in OB associations) influence these processes?  Circumstellar disks have traditionally been identified from near-infrared (near-IR) colours ($JHK$). Recently, several L-band surveys of young stellar populations proved that the K-band excess is a rather incomplete and unreliable disk indicator. Furthermore, theoretical work on the IR signatures of circumstellar disks suggests that L-band observations can detect disks even for very low disk masses (Wood et al.\\ \\cite{wood02}). Thus L-band surveys are very efficient in detecting circumstellar material and are thought to be largely complete for very young clusters --- Haisch et al\\ (\\cite{haisch01b}) found this to be case for the embedded cluster NGC\\,2024. However, for older clusters, detection in the L-band might become more difficult if disk evolution leads to the removal of the hotter circumstellar dust component. Furthermore these surveys still do not provide enough information on the geometry of the system (e.g.\\ protostellar object versus young stellar object with disk) and they provide little insight into the properties of the circumstellar material.  Mid-IR observations can unequivocally identify circumstellar material, constrain the spectral energy distribution and the physical parameters of the observed system and, through spectroscopy, allow the identification of the chemical and mineral species present in the dust grains. However such observations are technically challenging and relatively few clusters and associations have been surveyed in the N-band (e.g.\\ Taurus-Auriga, Kenyon \\& Hartmann \\cite{kenyon95}; $\\rho$\\,Ophiuchi cloud, Green et al.\\ \\cite{green94}; NGC\\,2024, Haisch et al.\\ \\cite{haisch01b}; NGC\\,3603, N\\\"{u}rnberger \\& Stanke \\cite{nurnberger03}), and mid-IR spectroscopy (around 10\\,$\\mu$m) has mostly concentrated on the more massive Herbig Ae/Be objects (Bouwman et al.\\ \\cite{bouwman01}) and objects in the Taurus-Auriga and $\\rho$\\,Ophiuchi complexes (Hanner et al.\\ \\cite{hanner95,hanner98}).  $\\sigma$\\,Orionis is a Trapezium-like system with an O9.5\\,V primary. The population of low-mass stars spatially clustered around this system was discovered as bright X-ray sources in ROSAT images (Wolk \\cite{wolk96}; Walter et al.\\, \\cite{walter97}). A recent L$'$-band survey of low-mass $\\sigma$\\,Orionis cluster members has revealed that $\\sim$\\,46\\% of these objects have circumstellar disks, at a cluster age of 3$-$5\\,Myr (Oliveira et al.\\ \\cite{oliveira04}). A mid-infrared source has been discovered very close to $\\sigma$ Orionis itself, apparently a proto-planetary disk being dispersed by the intense ultraviolet radiation from this massive star (van Loon \\& Oliveira \\cite{loon03}). A few IRAS sources were known in the vicinity of $\\sigma$\\,Ori. In this paper we describe new mid-IR imaging observations within an area around $\\sigma$\\,Ori, aimed at revealing the nature of the known mid-IR sources and detecting mid-IR emission from other dusty pre-main-sequence (PMS) stars. For a few of these objects we obtained 8$-$13\\,$\\mu$m spectra in order to determine the composition of the circumstellar dust. ",
        "conclusions": "We have performed mid-IR observations of the components of the $\\sigma$\\,Ori multiple system and of several suspected young stars in the vicinity of this system. We used N1-band observations to unequivocally ascertain whether circumstellar material is present around these objects. As expected, the early-type members of the multiple system were all found to be devoid of circumstellar dust (or free-free emission). From the suspected 10 young late-type objects, one object shows IR magnitudes consistent with the stellar photosphere and 7 objects clearly show evidence for circumstellar dust material, of which one object (IRAS\\,05358$-$0238) we classify as Class\\,I --- i.e.\\, it exhibits a very substantial excess at these wavelengths. For 2 other objects we could only obtain upper limits for their N1-band brightness.  A comparison of $KL'N1$ colours with Taurus-Auriga observations and model expectations suggests that at least 4 objects (not including IRAS\\,05358$-$0238, for which we do not have an L$^\\prime$-band measurement) possess rather massive circumstellar disks and seem to be actively accreting from them. Only one of these objects has been confirmed as a member of the $\\sigma$\\,Ori cluster; however the detection of circumstellar dusty material hints at a young age, and it seems unlikely that these are all interlopers coming from other sites of more recent star formation. If these objects are indeed cluster members, it would imply that, at an age of 3$-$5\\,Myr (e.g.\\, B\\'{e}jar et al.\\ \\cite{bejar99}; Oliveira et al.\\ \\cite{oliveira02}; Oliveira et al.\\ \\cite{oliveira04}), the more massive late-type cluster members still have massive accretion disks. The presence of accretion disks seems to extend to lower masses (Zapatero Osorio et al.\\ \\cite{osorio02}; Oliveira et al.\\ \\cite{oliveira04}). Either these objects are younger than the bulk of (non-accreting) cluster members or accretion disks (can) survive relatively long.  Our detections of mid-IR excess emission provide the first evidence for the presence of circumstellar disks in the irradiated-jet sources Haro\\,5-39/HH\\,447, V510\\,Ori/HH\\,444 and V603\\,Ori/HH\\,445. This supports the belief that accretion disks feed and help collimate the fast polar outflows responsible for the Herbig-Haro structures.  For the brightest (in the N1-band) objects IRAS\\,05358$-$0238 and TX\\,Ori, we also performed imaging in the Q1-band and spectroscopy in the N-band. This allowed us to probe the physical and chemical conditions of their circumstellar environments. The analysis of the SEDs of these objects provided us with typical circumstellar dust temperatures.  The mid-IR spectrum of TX\\,Ori reveals a typical silicate emission feature, but with an extra component at 11.2\\,$\\mu$m that can be reproduced by addition of optically thin emission from crystalline silicate (forsterite). This can be regarded as evidence of dust processing. We are unable to establish whether dust particle coagulation has occurred as well.  The spectrum of IRAS\\,05358$-$0238 is extremely unusual. The ratio of IR to stellar luminosity approaches unity, indicating the presence of a substantial amount of dust. However, the spectrum does not show the typical silicate feature; in fact, we are only able to describe it as being dominated by forsterite in self-absorption, something which has never been observed before. The status of this object remains uncertain, though, without evidence for (nor against) either its association with the $\\sigma$\\,Ori cluster or its youth.  \\appendix"
    },
    "0402/astro-ph0402287_arXiv.txt": {
        "abstract": "{ We discuss the effects of rotation on the evolution of accreting carbon-oxygen white dwarfs, with the emphasis on possible consequences in Type Ia supernova (SN Ia) progenitors. Starting with a slowly rotating white dwarf, we consider the accretion of matter and angular momentum from a quasi-Keplerian accretion disk. Numerical simulations with initial white dwarf masses of 0.8, 0.9 and  1.0 \\Msun{} and accretion of carbon-oxygen rich matter at rates of $3\\dots10\\times10^{-7}$ \\msyr{} are performed. The models are evolved either up to a ratio of rotational to potential energy of $T/W=0.18$ --- as angular momentum loss through gravitational wave radiation will become important for $T/W < 0.18$ --- or to central carbon ignition. The role of the various rotationally induced hydrodynamic instabilities for the transport of angular momentum inside the white dwarf is investigated. We find that the dynamical shear instability is the most important one in the highly degenerate core, while Eddington Sweet circulation, Goldreich-Schubert-Fricke instability and secular shear instability are most relevant in the non-degenerate envelope. Our results imply that accreting white dwarfs rotate differentially throughout, with a shear rate close to the threshold value for the onset of the dynamical shear instability. As the latter depends on the temperature of the white dwarf, the thermal evolution of the white dwarf core is found to be relevant for the angular momentum redistribution. As found previously, significant rotation is shown to lead to carbon ignition masses well above 1.4~\\Msun. Our models suggest a wide range of white dwarf explosion masses, which could be responsible for some aspects of the diversity  observed in SNe~Ia. We analyze the potential role of the bar-mode and the $r$-mode instability in rapidly rotating white dwarfs, which may impose angular momentum loss by gravitational wave radiation. We discuss the consequences of the resulting spin-down for the fate of the white dwarf, and the possibility to detect the emitted gravitational waves at frequencies of $0.1 \\dots 1.0$ Hz in nearby galaxies with LISA. Possible implications of fast and differentially rotating white dwarf cores for the flame propagation in exploding white dwarfs are also briefly discussed. ",
        "introduction": "Type Ia Supernovae (SNe Ia) have a particular importance in astrophysics. Observations of SNe Ia at low redshift showed a clear correlation between the peak brightness and the width of the light curve (Phillips~\\cite{Phillips93}; Phillips' relation), which allowed to use SNe Ia as distance indicators for galaxies even beyond $z=1$. This made SNe Ia an indispensable tool for cosmology, in particular to determine the cosmological parameters (e.g. Hamuy et al.~\\cite{Hamuy96}; Branch~\\cite{Branch98}; Leibundgut~\\cite{Leibundgut00},~\\cite{Leibundgut01}). The new cosmology with a non-zero cosmological constant has been initiated by the observational evidence deduced from observations of SNe Ia at high redshift (Perlmutter et al.~\\cite{Perlmutter99}; Riess et al.~\\cite{Riess00}).  Recent analyses of SNe Ia have revealed, however, that SNe Ia are not perfectly homogeneous but show some diversity in their light curves and spectra (e.g. Branch~\\cite{Branch01}; Nomoto et al.~\\cite{Nomoto03}; Li et al.~\\cite{Li03}). This leaves concerns about applying the Phillips' relation to very distant SNe Ia. An understanding of the origin of the diversity observed in SNe Ia is thus a crucial task for stellar evolution theory, which requires to identify the detailed evolutionary paths of SNe Ia progenitors.  Unlike core collapse supernovae, Type Ia supernovae (SNe Ia) are believed to occur exclusively in binary systems (e.g. Livio~\\cite{Livio01}). Although it is still unclear which kinds of binary systems lead to SNe Ia, non-degenerate stars such as main sequence stars, red giants or helium stars are often assumed as white dwarf companion (e.g. Li \\& van den Heuvel~\\cite{Li97}; Hachisu et al.~\\cite{Hachisu99}; Langer et al.~\\cite{Langer00}; Yoon \\& Langer~\\cite{Yoon03}). The white dwarf is then assumed to grow to the Chandrasekhar limit by mass accretion from its companion, with accretion rates which allow steady shell hydrogen and helium burning (\\Mdot{} $\\gsim 10^{-7}$ \\msyr{}). An understanding of the physics of mass accretion is therefore indispensable to investigate the evolution of accreting white dwarfs.  Although the mass accretion process in white dwarfs has been discussed by many authors (e.g. Iben~\\cite{Iben82}; Nomoto~\\cite{Nomoto82}; Fujimoto \\& Sugimoto~\\cite{Fujimoto82}; Saio \\& Nomoto~\\cite{Saio85},~\\cite{Saio98}; Kawai et al.~\\cite{Kawai88}; Cassisi et al.~\\cite{Cassisi98}; Piersanti et al.~\\cite{Piersanti00}; Langer et al.~\\cite{Langer02}), little attention was so far devoted to the effects of angular momentum accretion and the ensuing white dwarf rotation (see Sect.~\\ref{sect:previous}). As the evolution of stars can generally be affected by rotation (e.g. Heger \\& Langer~\\cite{Heger00b}; Maeder \\& Meynet~\\cite{Maeder00}), this might be particularly so in accreting white dwarfs: Since the transfered matter from the white dwarf companions may form a Keplerian disk which carries a large amount of angular momentum, the resultant accretion of angular momentum will lead to the spin-up of the white dwarf (e.g. Durisen~\\cite{Durisen77}; Ritter~\\cite{Ritter85}; Narayan \\& Popham~\\cite{Narayan89}; Langer et al.~\\cite{Langer00},~\\cite{Langer02},~\\cite{Langer03}). The observation that white dwarfs in cataclysmic variables rotate much faster than isolated ones (Sion~\\cite{Sion99}) provides evidence for accreting white dwarfs indeed being spun up. Rapidly rotating progenitors may also lead to aspherical explosions, which may give rise to the observed polarization of SNe Ia (Howell et al.~\\cite{Howell01}; Wang et al.~\\cite{Wang03}).  Here, we make an effort to investigate in detail the possibility of angular momentum accretion, and the role of the various rotationally induced hydrodynamic instabilities in transporting angular momentum into the white dwarf core, and in establishing the the pre-explosion angular momentum profile. The remainder of this paper is organized as follows. We evaluate the possible mechanisms for angular momentum transport in accreting white dwarfs in Sect.~\\ref{sect:angmom}. Our approach to the problem, including the numerical methods and physical assumptions, is discussed in Sect.~\\ref{sect:simulation}, where previous papers on rotating white dwarf models are also reviewed (Sect.~\\ref{sect:previous}). Numerical results are presented in Sect.~\\ref{sect:results}, with the emphasis on the process of angular momentum transport in the white dwarf interior. Pre-explosion conditions of accreting white dwarfs and their implications for the diversity of SNe Ia are discussed in Sect.~\\ref{sect:final}, ~\\ref{sect:diversity} and ~\\ref{sect:explosion}. The possibility of detecting gravitational waves from SNe Ia progenitors is examined in Sect.~\\ref{sect:gwr}. Our conclusions are summarized in Sect.~\\ref{sect:conclusion}. ",
        "conclusions": "\\label{sect:conclusion}  We summarize the results of this paper as follows.  1. The role of the Eddington sweet circulation, the GSF instability and the shear instability for the transport of angular momentum in non-magnetized white dwarfs has been investigated (Sect.~\\ref{sect:angmom}). Although Eddington sweet circulation and the GSF instability are important for the redistribution of angular momentum in the non-degenerate envelope, their importance is small in the degenerate core compared to the shear instability. The secular shear instability can not operate in the strongly degenerate core, since the thermal diffusion time becomes longer than the turbulent viscous time for densities higher than a critical value (i.e., $\\rho \\gsim 10^6 \\dots 10^7 ~{\\rm g/cm^3}$, Fig.~\\ref{fig:rho_ssi}). On the other hand, the criterion for the dynamical shear instability is significantly relaxed for higher density because the buoyancy force becomes weaker with stronger degeneracy (Fig.~\\ref{fig:sigma}). As a result, the degenerate inner core is dominated by the dynamical shear instability in accreting white dwarfs as shown in Sect.~\\ref{sect:spin}.  2. We have followed the redistribution of the angular momentum in accreting white dwarfs by the above mentioned processes (Sect.~\\ref{sect:spin}). We find that accreting white dwarfs do not rotate rigidly, but differentially throughout their evolution, for the considered accretion rates (\\Mdot{} = $3\\dots10 \\times 10^{-7}$ \\msyr). In the degenerate core, once the shear factor decreased to the threshold value for the onset of the dynamical shear instability, the time scale for further angular momentum transport becomes larger than the accretion time scale. Accordingly, strong differential rotation is retained in the inner core with a shear strength near the threshold value for the dynamical shear instability (Fig.~\\ref{fig:sigma_6a}).   3. Accreting white dwarfs, as they rotate differentially, may not reach central carbon ignition even when they grow beyond the canonical Chandrasekhar limit of $\\sim$1.4~\\Msun. This is in accordance with previously obtained results based on other methods (Ostriker \\& Bodenheimer~\\cite{Ostriker68a},~\\cite{Ostriker73}; Durisen~\\cite{Durisen75b}; Durisen \\& Imamura~\\cite{Durisen81}). A secular instability to gravitational wave radiation through the $r$-mode or the bar-mode may be important for determining the final fate of accreting white dwarfs. The masses of exploding white dwarfs are expected to vary in the range from the canonical Chandrasekhar mass ($\\sim$1.4 \\Msun) to the maximum possible mass that the white dwarf can achieve by mass accretion in a binary system ($\\sim 2.0$~\\Msun, Langer et al.~\\cite{Langer00}). This may have consequences for the diversity in the brightness and polarization of SNe~Ia.  4. Fast rotation in the white dwarf core may change the supernova explosion since it can affect the evolution of the turbulent nuclear burning flames by providing a large amount of turbulent kinetic energy and/or by enhancing the burning surface significantly. It needs to be clarified whether this may be a plausible mechanism to induce the transition from deflagration to detonation.   5. White dwarfs which accreted enough angular momentum may be detectable sources of gravitational waves in the near future. Our models show that these will emit gravitational waves with frequencies of $0.1 - 1.0$ Hz. Space-based interferometric gravitational wave detectors such as LISA could observe such signals from rapidly rotating SNe Ia progenitors in nearby galaxies."
    },
    "0402/astro-ph0402293_arXiv.txt": {
        "abstract": "{On 2003 September 17 INTEGRAL discovered a bright transient source 3$\\degr$ from the Galactic Center, \\src. The field containing the transient was observed by XMM-Newton on 2003 March 17 and September 11 and 17.  A bright source, at a position consistent with the INTEGRAL location, was detected by the European Photon Imaging Camera (EPIC) during both September observations with mean 0.5--10\\,keV unabsorbed luminosities of 1.1$\\times$10$^{35}$ and 5.7$\\times$10$^{35}$\\,erg\\,s$^{-1}$ for an (assumed) distance of 8\\,kpc.  The source was not detected in 2003 March with a 0.5--10\\,keV luminosity of $<$3.8$\\times$10$^{32}$\\,erg\\,s$^{-1}$.  The September 11 and 17 EPIC spectra can be represented by a power-law model with photon indices of 2.25$\\pm$0.15 and 1.42$\\pm$0.17, respectively. Thus, the 0.5--10\\,keV spectrum hardens with increasing intensity. The low-energy absorption during both September observations is comparable to the interstellar value.  The X-ray lightcurves for both September observations show energy dependent flaring which may be modeled by changes in either low-energy absorption {\\it or} power-law index. ",
        "introduction": "\\label{sect:intro}  About a dozen new hard X-ray transients have been discovered in the last year during INTEGRAL (Winkler et al.~\\cite{w:03}) observations of the galactic center region with the soft gamma-ray imager IBIS/ISGRI (Lebrun et al.~\\cite{le:03}).  Their unusual spectral hardness has led to suggestions that these sources comprise a group of highly absorbed galactic binaries (Revnivtsev et al.~\\cite{re:03}). These are being preferentially detected due to the good sensitivity and large field of view of IBIS/ISGRI above 15\\,keV. XMM-Newton observations of the first of these, IGR\\,J16318$-$4848, revealed intense Fe~K$\\alpha$ and K$\\beta$ and Ni~K$\\alpha$ emission lines as well as strong low-energy absorption (Matt \\& Guainazzi~\\cite{mg:03}; Walter et al.~\\cite{wa:03}).  XMM-Newton observations of IGR\\,J16320$-$4851 revealed a featureless hard continuum (Rodriguez et al.~\\cite{r:03b}). IGR\\,16358$-$4726 was observed by {\\it Chandra} and showed a hard power-law spectrum with a 5880\\,s periodic intensity modulation (Patel et al.~\\cite{p:03}).   \\begin{table*}  \\caption[]{XMM-Newton observation log. The EPIC modes are Full Frame (FF) and Timing (TI).  In the 2003 March 17 observation \\src\\ was outside the OM field of view. The effective wavelengths of the UVW1, UVW2 and UVM2 filters are 2945\\,\\AA, 2180\\,\\AA, and 2340\\,\\AA, respectively.  } \\begin{flushleft} \\begin{tabular}{lccccccccc} \\hline \\hline\\noalign{\\smallskip} Obs & Start Time & End Time & \\mc{2}{c}{Exposure (ks)} & \\mc{3}{c}{EPIC mode} & \\mc{2}{c}{OM} \\\\ &     (UTC)  &   (UTC)  & pn &MOS                  &  pn & MOS1 & MOS2    & Exp (ks) & Filter\\\\ \\noalign{\\smallskip\\hrule\\smallskip} 1 & 2003~Mar~17~20:53 & Mar 18~00:33 & 10.2 & 11.9 & FF  & FF & FF & \\dots & \\dots   \\\\ 2 & 2003~Sep~11~17:54 & Sep~11~22:09 & 9.3  & 11.0 & FF  & FF & FF & 2.2, 3.2 & B, UVW1  \\\\ 3 & 2003~Sep~17~17:13 & Sep~17~19:59 & 2.5  & 8.3  & TI  & TI & FF & 1.8, 1.6 & UVM2, UVW2 \\\\ \\noalign{\\smallskip\\hrule\\smallskip} \\end{tabular} \\end{flushleft} \\label{tab:obs}  \\end{table*}   We report on a new bright transient source, \\src, discovered using IBIS/ISGRI on 2003 September 17 at 01:10 UTC during an observation of the Galactic Center region (Sunyaev et al.~\\cite{s:03}).  The source intensity was about 160\\,mCrab, 60\\,mCrab, and $<$15\\,mCrab (at 3$\\sigma$ confidence) in the 18--25\\,keV, 25--50\\,keV and 50--100\\,keV energy ranges. \\src\\ was bright for around 2 hours and then faded below the IBIS/ISGRI detection threshold. The source was again detected in outburst by IBIS/ISGRI later the same day between 06 and 14~hrs~UTC (Grebenev et al.~\\cite{g:03}).  By chance, XMM-Newton observed the region of sky containing \\src\\ only 5 days before the INTEGRAL discovery. A source was clearly detected at a position consistent with that reported from INTEGRAL, with a mean 2--10\\,keV intensity of 4.5$\\times$10$^{-12}$\\,erg\\,cm$^{-2}$~s$^{-1}$ (Gonz\\'alez-Riestra et al.~\\cite{gr:03}).  Rodriguez et al.~(\\cite{r:03a}) reported on a possible optical/infrared counterpart in the USNO~B1.0 catalog (B = 13.9--14.5$\\pm$0.3~mag) and 2MASS all-sky quick-look image archive.  They also noted that there are 3 fainter candidates within the preliminary 10\\arcsec\\ XMM-Newton uncertainty region in the 2MASS image.  The field containing \\src\\ was observed three times by XMM-Newton. The first two observations (2003 March 17 and September 11) were part of a program to study the nova V4643\\,Sgr (see also Gonz\\'alez-Riestra et al.~\\cite{gr:03}), and the third was a Target of Opportunity observation triggered by the INTEGRAL discovery. Here, we present results from all three XMM-Newton observations.  \\begin{figure*} \\begin{center} \\includegraphics[height=17cm,angle=90]{x-ray_image.ps} \\caption[]{2003 September 11 EPIC-pn (left) and OM B filter (right) images of the area around the 2\\arcmin\\ INTEGRAL uncertainty region for IGR J17544-2619 of Sunyaev et al. (\\cite{s:03}, big circle). The contours shown in the optical image correspond to the EPIC-pn image (0.005, 0.010, 0.015, 0.020 and 0.025 counts~s$^{-1}$), which clearly excludes 1RXS\\,J175428.3-262035 (small circle).} \\label{fig:EPIC_image} \\end{center} \\end{figure*}   The XMM-Newton Observatory (Jansen et al.~\\cite{j:01}) includes three 1500\\,cm$^2$ X-ray telescopes each with an European Photon Imaging Camera (EPIC) at the focus. Reflection Grating Spectrometers (RGS, den Herder et al.~\\cite{dh:01}) are located behind two of the telescopes. In addition, a coaligned optical/UV Monitor (OM, Mason et al.~\\cite{m:01}) is included.  Two of the EPIC imaging spectrometers use MOS CCDs (Turner et al.~\\cite{t:01}) and one uses a pn CCD (Str\\\"uder et al.~\\cite{s:01}). ",
        "conclusions": "\\label{sect:discussion}  The mean 0.5--10\\,keV unabsorbed luminosity of \\src\\ during the 2003 September 11 and 17 observations is 1.1$\\times$10$^{35}$\\,erg\\,s$^{-1}$ and 5.7$\\times$10$^{35}$\\,erg\\,s$^{-1}$ for an (assumed) distance of 8\\,kpc. The peak reached during the flare on 2003 September 17 is about 8.5$\\times$10$^{35}$\\,erg\\,s$^{-1}$ Such luminosities are only reached in systems where the compact object is a neutron star or black hole.  The state observed in 2003 March corresponds to an 0.5--10\\,keV luminosity of $<$4$\\times$10$^{32}$\\,erg\\,s$^{-1}$. This upper limit is consistent with the luminosities observed for quiescent X-ray binary transients containing either a neutron star (e.g., Campana \\& Stella~\\cite{cs:03}) or a black hole (e.g., Tomsick et al.~\\cite{t:03}).  The total dynamic flux range (quiescence to peak flaring) seen during the three XMM-Newton observations is a factor $\\approxgt$2000.  This is also typical for such transient X-ray binaries.  The (maximum) observed fluxes from the INTEGRAL observations (Sunyaev et al.~\\cite{s:03}) may constrain the spectral model in the soft and hard X-ray bands, as well as the level of X-ray activity. Given the quoted fluxes, and assuming that the hard X-ray spectrum consists of a single power-law, the photon index was $\\sim$4 during the INTEGRAL observations. Assuming the interstellar \\nh, this would give an extrapolated absorbed 0.5--10\\,keV flux of $\\sim$6.4\\,Crab. If during the INTEGRAL observations the spectral index was $\\sim$2.2, as measured in the 0.5--10\\,keV energy range on 2003 September 11, a high-energy cut-off at $\\sim$14\\,keV is required in the spectrum in order to explain the (maximum) 18--25\\,keV and 25--60\\,keV fluxes and the 50--100\\,keV flux upper limit. Extrapolating this spectrum gives an absorbed flux of $\\sim$0.45\\, Crab (0.5--10\\,keV).  Note that this value would correspond to an unabsorbed luminosity of about 3$\\times$10$^{38}$\\,erg\\,s$^{-1}$ (at 8\\,kpc), similar to that reached by classical X-ray binary transients (e.g., Chen et al.\\ 1997). Unfortunately, there are no closeby (i.e., within hours) {\\it RXTE} All-Sky Monitor (ASM) measurements of \\src\\ (R.~Remillard, private communication), to verify whether the source was active around the time of the INTEGRAL observations. The {\\it RXTE}/ASM measurements within days of the INTEGRAL detections give typical upper limits of $\\sim$50\\,mCrab.  But since the source is highly variable in both soft and hard X-rays, these do not provide stringent constraints either.  We note that the quiescent state and the low-level flaring seen with XMM-Newton and the high-level activity seen by INTEGRAL is reminiscent of SAX\\,J1819.3$-$2525 (V4641\\,Sgr).  This system also shows low-level activity around 10$^{36}$\\,erg\\,s$^{-1}$ (in 't Zand et al.\\ 2000), strong and short high-level activity (e.g., Revnivtsev et al.~\\cite{re:02}), while in quiescence it reaches $\\simeq$4$\\times$10$^{31}$\\,erg\\,s$^{-1}$ (0.3--8\\,keV; Tomsick et al.~\\cite{t:03}).  The compact object in SAX\\,J1819.3$-$2525 is most probably a black hole (Orosz et al.~\\cite{o:01}).  Our observed OM magnitudes, combined with the optical/infrared magnitudes reported by Rodriguez et al. (\\cite{r:03a}), and assuming an absorption of 2$\\times10^{22}$ \\hcm\\ and a distance of 8 kpc, are consistent with an early O-type companion. However, a foreground object cannot be ruled out; we note the possible presence of fainter optical candidates in the XMM-Newton error circle (Rodriguez et al. \\cite{r:03a}).  Future observations will hopefully shed more light on \\src. In particular, monitoring campaigns by INTEGRAL and multi-wavelength observations may allow the nature of the compact object to be elucidated."
    },
    "0402/astro-ph0402546_arXiv.txt": {
        "abstract": "{ We present our proposal for a small X-ray mission DIOS (Diffuse Intergalactic Oxygen Surveyor) to perform survey observations of warm-hot intergalactic medium using OVII and OVIII emission lines.  This will be proposed to a small satellite program planned by ISAS/JAXA in Japan for a launch in 2008. The instrument consists of an array of TES microcalorimeters with an energy resolution 2 eV, cooled by mechanical coolers. The X-ray telescope will employ 4-stage reflection mirrors with a focal length as short as 70 cm and an angular resolution $2'$.  In addition to DIOS, we briefly describe the NeXT (New X-ray Telescope) mission, which is a larger Japanese X-ray observatory to be launched in 2010 and plans to explore non-thermal processes in the universe. ",
        "introduction": "Following the X-ray spectroscopic mission Astro-E2 to be launched in 2005 \\citep[see][]{furusho04,ohashi01}, Japanese X-ray astronomy groups are considering several satellite missions in the future. Here, we will describe a dedicated small mission to explore the structure of diffuse intergalactic medium, called DIOS (Diffuse Intergalactic Oxygen Surveyor), to be launched in 2008. We also briefly describe a larger X-ray mission NeXT (New X-ray Telescope mission) to appear in 2010. These missions will bring a substantial advance in our understanding of the hot-gas distribution in the universe, as well as in our own galaxy. ",
        "conclusions": ""
    },
    "0402/astro-ph0402636_arXiv.txt": {
        "abstract": "Our Milky Way Galaxy is a typical large spiral galaxy, representative of the most common morphological type in the local Universe.  We can determine the properties of individual stars in unusual detail, and use the characteristics of the stellar populations of the Galaxy as templates in understanding more distant galaxies.  The star formation history and merging history of the Galaxy is written in its stellar populations; these reveal that the Galaxy has evolved rather quietly over the last $\\sim 10$~Gyr. More detailed simulations of galaxy formation are needed, but this result apparently makes our Galaxy unusual if $\\Lambda$CDM is indeed the correct cosmological paradigm for structure formation. While our Milky Way is only one galaxy, a theory in which its properties are very anomalous most probably needs to be revised. Happily, observational capabilities of next-generation facilities should, in the the forseeable future, allow the aquisition of detailed observations for all galaxies in the Local Group. ",
        "introduction": "The origins and evolution of galaxies, such as our own Milky Way, and of their associated dark matter haloes are among the major outstanding questions of astrophysics.  Detailed study of the zero-redshift Universe provides complementary constraints on models of galaxy formation to those obtained from direct study of high-redshift objects.  Stars of mass similar to that of the Sun live for essentially the present age of the Universe and nearby low-mass stars can be used to trace conditions in the high-redshift Universe when they formed, perhaps even the `First Light' that ended the Cosmological Dark Ages.  While these stars may well not have formed in the galaxy in which they now reside (especially if the CDM paradigm is valid), several important observable quantities are largely conserved over a star's lifetime -- these include surface chemical elemental abundances (modulo effects associated with mass transfer in binaries) and orbital angular momentum (modulo the effects of torques and rapidly changing gravitational potentials).  Excavating the fossil record of galaxy evolution from old stars nearby allows us to do Cosmology locally, and is possible to some extent throughout the Local Group, with the most detailed information available from the Milky Way Galaxy.  I here discuss our knowledge of the stellar populations of the Milky Way and the implications for models of galaxy formation. Complementary results for M31 are presented by Brown (this volume). ",
        "conclusions": "The properties of the stellar populations of the Milky Way contain much information about the star formation history and mass assembly history of the Galaxy.  The Milky Way has merged with, is merging with, and will merge with, companion galaxies, which contribute stars, gas and dark matter.  Debris from the Sagittarius dwarf galaxy dominates recent accretion into the outer Galaxy, while the data are consistent with little stellar accretion into the inner Galaxy, including the disk. Predominantly gaseous accretion is relatively unconstrained, and is favoured by models of chemical evolution (cf.~Tosi's contribution). Planned and ongoing large spectroscopic surveys will tightly constrain the existence and orgins of stellar phase-space substructure.  The relatively quiescent merging history of the Milky Way that is implied by the mean properties of the stellar components is rather atypical in $\\Lambda$CDM cosmologies.  What about the rest of the Local Group?"
    },
    "0402/astro-ph0402400_arXiv.txt": {
        "abstract": "We present the results of a search for variable stars in the Local Group dwarf galaxy Phoenix. Nineteen Cepheids, six candidate long-period variables, one candidate eclipsing binary and a large number of candidate RR Lyrae stars have been identified. Periods and light curves have been obtained for all the Cepheid variables. Their distribution in the period--luminosity diagram reveals that both Anomalous Cepheids (AC) and short-period Classical Cepheids \\hbox{(s-pCC)} are found in our sample. This is the first time that both types of variable star are identified in the same system even though they likely coexist, but have gone unnoticed so far, in other low-metallicity galaxies like Leo~A and Sextans~A. We argue that the conditions for the existence of both types of variable star in the same galaxy are a low metallicity at all ages, and the presence of both young and intermediate-age (or old, depending on the nature of AC) stars. The RR Lyrae candidates trace, together with the well developed horizontal branch, the existence of an important old population in Phoenix. The different spatial distributions of s-pCC, AC and RR Lyrae variables in the Phoenix field are consistent with the stellar population gradients found in Phoenix, in the sense that the younger population is concentrated in the central part of the galaxy. The gradients in the distribution of the young population within the central part of Phoenix, which seem to indicate a propagation of the recent star formation, are also reflected in the spatial distribution of the s-pCC. ",
        "introduction": "\\label{intro}  Variable-star studies are experiencing a spectacular resurgence thanks to the microlensing projects which, as a by-product, are supplying huge samples of variable stars, mainly in the LMC, SMC and the Galactic Bulge. On another hand, HST and many ground-based sites with excellent seeing are allowing us to probe deep into the stellar populations of Local Group dwarf irregular (dIrr) galaxies, thus providing information on variable stars only observable before in the dwarf-galaxy satellites of the Milky Way. These enlarged samples of variable stars of different types provide the necessary information to deepen our understanding of the characteristics of each type, the relationships between them, and the physical mechanisms involved in their light variation, which yield important tests of stellar evolution models. At the same time, the gathering of data in different galaxy environments makes it possible to relate the types of variable stars to the stellar populations in their host galaxies.  These studies have provided, in particular, relatively large samples of short-period Classical Cepheids (s-pCC) in a number of dIrr galaxies: the SMC (Bauer et al.\\ 1999; Udalski et al.\\ 1999), IC1613 (Dolphin et al.\\ 2001), Leo~A (Dolphin et al.\\ 2002) and Sextans~A (Dolphin et al.\\ 2003). As pointed out by Gallart et al.\\ (1999) and Dolphin et al.\\ (2002), short-period (P~$\\simeq$~1 day) variables, with Cepheid-like light curves and luminosities about 1 magnitude above the horizontal branch can be produced by young\\footnote{Throughout the paper, we will consider as young populations those stars younger than 1 Gyr, while intermediate-age populations will be defined broadly as those having ages of 1--10 Gyr} stars in the phase of core He burning which, in the case of low-metallicity stars, experience blue loops extended enough to the blue to cross the instability strip. The so-called Anomalous Cepheids (AC) and s-pCC lie in similar positions in both the color--magnitude and the period--luminosity (PL) diagrams, and they are found in basically every dwarf Spheroidal (dSph) galaxy that has been surveyed for them, as well as in one globular cluster (Pritzl et al.\\ 2002; and references therein). The term ``Anomalous Cepheid'' was first introduced by Norris \\& Zinn (1975) because they are more luminous at a given period than the Population II Cepheids found in globular clusters. AC are $\\simeq\\,$0.5--2 magnitudes brighter than RR Lyrae stars, and their periods range from 0.3 to 2 days. Concerning their evolutionary status, there is general agreement that they are metal-poor stars with mass $\\simeq 1.5 M_\\odot$, occupying the instability strip during their horizontal-branch phase of evolution (Demarque \\& Hirshfeld 1975; Hirshfeld 1980; Bono et al.\\ 1997). They may be either i)~evolved, single, intermediate-age ($\\le 5$ Gyr) stars, or ii)~the evolved products of mass transfer in old ($\\simeq 12$ Gyr) binary systems.  The predicted behaviour in both cases is very similar (Bono et al.\\ 1997), and it has therefore been considered impossible to distinguish between the two possible origins of AC from their location in the PL and color-magnitude diagrams alone (Nemec, Nemec \\& Lutz 1994 and Dolphin et al.\\ 2002).   A reanalysis of all the data accumulated to date on AC by Bono et al.\\ (1997) and Pritzl et al. (2002) provided a slightly different locus for the AC in the PL diagram that, as we will show in the present paper, may allow one to discriminate these two types of variable star. Since they respectively trace stellar populations of different ages and of a very specific metallicity range, this distinction is relevant to the use of variable stars as bright tracers of stellar populations of different ages (in this case, AC as tracers of a fainter intermediate-age or old population), and as anchors of the age--metallicity relation of individual dwarf galaxies.  Prior to this study, AC and s-pCC have not been identified in the same galaxy.  As mentioned above, AC are routinely found in dSph galaxies while s-pCC have been found in dIrr galaxies.  Since both types of variable star have very similar characteristics, it may be that the classification of some of them has been prejudiced by the type of galaxy in which they were found. A suitable type of object in which to search for {\\it both\\/} types of variable are the so-called transition-type dSph/dIrr galaxies (Mateo 1998).  The closest such system is Phoenix. We had started a search for variable stars in this galaxy, and along the way we realized that it was, in fact, an ideal system to address this problem.  Phoenix is a low-mass ($M_{tot}=3.3\\times10^6 M_{\\odot}$), low-metallicity ([Fe/H]$\\,$=$\\,$--1.4) system at a distance of about 400 Kpc [$(m-M)_\\circ=23.0 \\pm 0.1$] from the Milky Way (Mateo 1998; Mart\\'\\i nez-Delgado, Gallart \\& Aparicio 1999, MGA99 hereinafter; Held, Saviane \\& Momany 1999; and references therein).  The fact that it has a small young stellar population (up to about 100 Myr old), embedded in a substantial population of old and intermediate-age stars (MGA99; Held et al.\\ 1999) and little or no gas (St-Germain et al.\\ 1999; Gallart et al. 2001) has motivated its classification as a transition type dSph/dIrr galaxy. No studies of variable stars in Phoenix have been published to date, although some suspected variable stars were reported by van~de~Rydt, Demers \\& Kunkel (1991) and MGA99. Both its low metallicity and the existence of stars of all ages made it a good candidate to search for both AC and s-pCC. In this paper we report such a discovery. We also discuss the possible existence of both types of variable star in other dwarf galaxies and the information that they may provide on the age--metallicity relations of their host galaxies. We also report, for the first time, the discovery of a large number of RR Lyrae variable star candidates in Phoenix which trace, together with the well developed horizontal branch, the existence and distribution of an old population in the galaxy. ",
        "conclusions": "A search for variable stars has been conducted in the Phoenix dwarf galaxy. Nineteen Cepheids (either AC or s-pCC), six candidate long-period variables, one candidate eclipsing binary and a large number of candidate RR Lyrae stars have been identified. Periods and light curves have been obtained for all the Cepheid variables. Their distribution in the PL diagram reveals that both AC and s-pCC are present in our sample. This is the first time that both types of variable star, which belong to metal-poor old/intermediate-age and metal-poor young populations respectively, have been identified in the same system. We show, however, that they also likely coexist in Leo~A and Sextans~A, even thought the fact originally went unnoticed. We note that, thanks to the very specific conditions of age and metallicity required for the occurrence of these variables, they can provide important hints on the age--metallicity relation of the host galaxy. For example, in the case of Phoenix they imply, according to current stellar-evolution models, a very low metallicity ($Z=0.0001$) for intermediate-age ($\\simeq$ 1--2 Gyr) stars, while a slightly larger metallicity $Z=0.0004-0.001$ is possible for the young ($< 1\\,$Gyr) population in the galaxy.  The RR Lyrae candidates, together with the well developed horizontal branch, trace the existence of an important old population in Phoenix. The different spatial distributions of AC, s-pCC, and RR Lyrae variables is consistent with the stellar population gradients found in Phoenix (MGA99), in the sense that the younger populations are progressively more concentrated toward the central part of the galaxy. The gradient in the mean age of the youngest populations in the center of Phoenix, which seems to indicate a propagation of the recent star formation as suggested by MGA99, is also reflected in the spatial distribution of the s-pCC."
    },
    "0402/astro-ph0402366_arXiv.txt": {
        "abstract": "{ Using a purely analytic approach to gaseous and dark matter halos, we study the cross-correlation between the Sunyaev-Zel'dovich (SZ) sky and galaxy survey under flat sky approximation, in an attempt to acquire the redshift information of the SZ map. The problem can be greatly simplified when it is noticed that the signals of the SZ-galaxy correlation arise only from hot gas and galaxies inside the same massive halos (i.e. clusters), and field galaxies make almost no contribution to the cross-correlation. Under the assumption that both the hot gas and galaxies trace the common gravitational potential of dark halos, we calculate the expected cross SZ-galaxy power spectra for the WMAP/Planck SZ maps and the SDSS galaxy sample at small scales $100<l<1000$. It turns out, however, that it is not presently feasible to measure such small angular cross power spectra because of the high noise levels at $l>400$ with the WMAP/Planck experiments. Future SZ observations with better angular resolutions and sufficiently wide sky coverages will be needed if this technique is applied for the statistical measurement of redshift distribution of the SZ sources. ",
        "introduction": "Most of the baryons in the local universe exist in the form of warm-hot intergalactic medium with temperature of $T\\sim 10^5-10^7$K as a result of  gravitationally driven shocks and adiabatical compression when they fall into large-scale structures and collapsed dark matter halos (e.g. Cen \\& Ostriker 1999; Dav\\'e et al. 2001). In the latter case, the very hot baryons gravitationally bound in massive halos such as groups and clusters usually manifest themselves by strong diffuse X-ray sources in terms of bremmstrahlung emission, which are directly detectable with current X-ray instruments. Moreover, the energetic electrons of the hot gas also interact the passing cosmic microwave background (CMB) photons through the so-called Sunyaev-Zel'dovich (SZ) effect,  giving rise to a subtle change in the CMB spectrum. Current SZ measurements of known clusters at high signal-to-noise are now routine (e.g. Carlstrom et al. 2002) and recent detection of the excess power relative to primordial CMB anisotropy at arcminute scales has been successfully attributed to the statistical signals of the thermal SZ effect (Mason et al. 2003; Bond et al. 2003; Komatsu \\& Seljak 2002).   However, both X-ray and SZ measurements contain no information about the redshift distributions of the hot baryons and their host groups/clusters. Spectroscopic follow-up observations should be made to complement our knowledge of the location and evolution of the hot baryons and their host halos. It has been realized that a more practical and powerful approach to extracting the distance information from the X-ray and SZ maps is perhaps to cross-correlate the X-ray and SZ maps with the existing galaxy catalogs or ongoing deep galaxy surveys (e.g. Seljak, Burwell \\& Pen 2000; Zhang \\& Pen 2001;2003; Zhang, Pen \\& Wang 2002). In particular, the problem can be greatly simplified if we notice that the signals of the cross-correlation between the SZ (or X-ray) map and galaxy survey arise primarily from cluster galaxies, and field galaxies make almost no contribution to the SZ(or X-ray)-galaxy correlation. This permits a straightforward calculation of the power spectrum of the SZ(or X-ray)-galaxy correlation from an analytic model of dark halo abundance, i.e., the Press \\& Schechter (1974; PS) formalism, along with a reasonable prescription of galaxy and gas distributions inside a given halo of mass $M$ at redshift $z$. For the latter, one can adopt either the simplest self-similar NFW-like profile suggested by  numerical simulations (Navarro, Frenk \\& White 1995; NFW), or the empirically motivated density profiles such as the King model or $\\beta$-model, in combination with the halo occupation distribution (see Cooray \\& Sheth 2003 for a recent review). In this paper, we explore the power spectrum of the SZ-galaxy cross-correlation based purely on the halo approach. Similar technique has been recently applied by Zhang \\& Pen (2003) to the study of the cross-correlation of the soft X-ray background and galaxy survey and by Wu \\& Xue (2003) to the study of the auto-correlation of the soft X-ray background. In their early work, Zhang \\& Pen (2001) also investigated the SZ-galaxy cross-correlation using the so-called continuum field model, in which the gas distribution in dark halos is given by a convolution of the dark matter distribution with a Gaussian window function, while the matter fluctuation in the highly non-linear regime (e.g. clusters) is calculated in terms of the hyper-extended perturbation theory.   For the purpose of actual applications, we will calculate the expected cross-correlation between the SZ maps observed by WMAP/Planck and the galaxy survey by the Sloan Digital Sky Survey (SDSS), and assess the feasibility of extracting the redshift information of the WMAP/Planck SZ maps using this statistical approach. Throughout this paper we adopt a flat cosmological model ($\\Lambda$CDM) with the best fit parameters determined by WMAP: $\\Omega_{\\Lambda}=0.73$, $\\Omega_M=0.27$, $h=0.71$, $\\Omega_b h^2=0.0224$, $\\sigma_8=0.84$ and $n_s=0.93$ ",
        "conclusions": "Cross-correlation between the SZ map and galaxy survey will reveal, in a statistical manner, important information on the redshift distribution of hot baryons in the universe, which requires no spectroscopic follow-up observation of individual sources. Recall that the major advantage of measurement of the SZ power spectrum over the SZ cluster survey is that one can acquire weak SZ signals at high statistical significance level without the need for resolving individual clusters. Indeed, time-consuming spectroscopic follow-up observations of the SZ sources may eventually throw a shadow on the effective applications of the SZ power spectrum if redshift distribution is ultimately concerned.  Cross-correlation between the SZ map and galaxy survey provides a simple approach which is just suited for the problem. In particular, theoretical prediction of the SZ-galaxy power spectrum can be greatly simplified if one notices that the signals of the SZ-galaxy cross-correlation are primarily produced by the hot gas and galaxies within the same clusters. Namely, field galaxies make almost no contribution to the SZ-galaxy cross-correlation.  Assuming that both the intracluster gas and cluster galaxies follow the same dark matter distribution with a functional form of the NFW-like profile and the abundance of dark halos is described by the PS mass function, we have predicted the SZ-galaxy angular power spectrum. As it is expected, the power spectrum indeed shows a moderately strong correlation indicated by an overall correlation coefficient of $r_{\\rm corr}\\approx 0.2$ at $100<1<1000$. However,  applying our algorithm to the SDSS galaxy survey and MAP/Planck SZ maps yields an unpleasant result: It is unlikely that one can acquire meaningful information about the redshift distribution of the MAP/Planck SZ maps because the SZ signals are significantly below the noise levels in our interested range of multipoles, $l>1000$, although the SDSS galaxy sample should be just suited for such a purpose. One possibility is to work at large angular scales $l<100$, which is related to the SZ effect produced by nearby superclusters and clustering of galaxies (Hern\\'andez-Monteagudo \\& Rubin\\~o-Mart\\'in 2003; Afshordi et al. 2004). Another application is to utilize the CMB data at very small angular scales beyond $l=2000$ to be obtained by many ongoing/upcoming  experiments such as ACBAR, AMiBA, BIMA, CBI, etc. In this case, one needs instead to deal with the small-sky coverage problem. Indeed, a similar work should be done on the feasibility of cross-correlating these small-sky coverage CMB data with  existing galaxy catalogs. We conclude that the cross-correlation between the SZ map and galaxy survey can in principle yield valuable information about the redshift distribution of the host baryons in the universe. However, the actual applications of this technique to real observations may not be possible until high angular resolution and sensitivity SZ power spectrum over a wide sky coverage is achieved."
    },
    "0402/astro-ph0402199_arXiv.txt": {
        "abstract": "We describe a new code which can accurately calculate the relativistic effects which distort the emission from an accretion disc around a black hole. We compare our results for a disk which extends from the innermost stable orbit to $20r_{g}$ in both Schwarzschild and maximal ($a=0.998$) Kerr spacetimes with the two line profile codes which are on general release in the XSPEC spectral fitting package. These models generally give a very good description of the relativistic smearing of the line for this range of radii. However, these models have some limitations. In particular we show that the assumed form of the {\\em angular} emissivity law (limb darkening or brightening) can make significant changes to the derived line profile where lightbending is important. This is {\\em always} the case for extreme Kerr spacetimes or high inclination systems, where the observed line is produced from a very large range of different emitted angles. In these situations the assumed angular emissivity can affect the derived {\\em radial} emissivity. The line profile is not simply determined by the well defined (but numerically difficult) physical effects of strong gravity, but is also dependent on the poorly known astrophysics of the disc emission. ",
        "introduction": "Material in an accretion disk around a black hole is orbiting at high velocity, close to the speed of light, in a strong gravitational potential. Hence its emission is distorted by doppler shifts, length contraction, time dilation, gravitational redshift and lightbending. The combined impact of these special and general relativistic effects was first calculated in the now seminal paper of \\cite{C75}, where he used  a \\emph{transfer function} to describe the relativistic effects. The observed spectrum from an accretion disc around a Kerr black hole is the  convolution of  this with the intrinsic disc continuum emission.  While such models have been used to try to determine the gravitational potential from the observed accretion disk spectra (e.g. \\citealt{LN89,EMH91,E93,M00,GME01}), these attempts suffer from our limited knowledge of the spectral shape of the intrinsic accretion disk emission (see e.g. the review by \\citealt{B02}). It is much easier to determine the relativistic effects from a {\\em sharp} spectral feature, such as the iron fluorescence line expected from X-ray illumination of an accretion disc \\citep{F89}. An originally narrow atomic transition is transformed into broad, skewed profile whose shape is given {\\em directly} by the transfer function.  Observationally, evidence for a relativistically smeared iron line first came from the ASCA observation of the active galactic nuclei (AGN) MCG-6-30-15 \\citep{T95}. Further observations showed evidence for the line profile being so broad as to require a maximally spinning black hole \\citep{I96}. More recent data from XMM are interpreted as showing that the line is even wider than expected from an extreme Kerr disk, requiring direct extraction of the spin energy from the central black hole as well as the immense gravitational potential \\citep{W01}.  Such results are incredibly exciting, but X-ray spectral fitting is not entirely unambiguous. There is a complex reflected continuum as well as the line (\\citealt{NKK00,BRF01}). For an ionised disk (as inferred for MCG-6-30-15) the current models in general use ({\\tt pexriv} in the {\\sc XSPEC} spectral fitting package) are probably highly incomplete \\citep{RFY99}. Complex ionised absorption also affects AGN spectra (e.g. \\citealt{K02}) and the illuminating continuum itself can have complex curvature rather than being a simple power law.  However, in MCG-6-30-15 these issues have been examined in detail, and the results on the dramatic line width appear robust (\\citealt{VF03,R04}).  Thus there is a clear requirement that the extreme relativistic effects are well modelled. There are two models which are currently widely available to the observational community, within the \\texttt{XSPEC} spectral fitting package, \\texttt{diskline} (based on \\citealt{F89}) and \\texttt{laor} \\citep{L91}. The analytic \\texttt{diskline} code models the line profile from an accretion disc around a Schwarzschild black hole (so of course cannot be used to describe the effects in a Kerr geometry). Also, it does not include the effects of lightbending (although \\citealt{F89} outline a scheme for incorporating this) and hence does not accurately calculate all the relativistic effects for $r< 20r_{g}$ (where $r_g=GM/c^2$).  By contrast, the \\texttt{laor} model numerically calculates the line profile including lightbending for an extreme Kerr black hole, but uses a rather small set of tabulated transfer functions which limit its resolution and accuracy (see Section 3.3).  While there are other relativistic codes in the literature which do not suffer from these limitations, these are not generally readily and/or easily available for observers to use. There is a clear need for a fast, accurate, high resolution code which can be used to fit data from the next generation of satellites. In this paper we describe our new code for computing the relativistic iron line profile in both the Schwarzschild and Kerr metrics. We compare this with the \\texttt{diskline} and \\texttt{laor} models in \\texttt{XSPEC} for discs which extend down to the last stable orbit in their respective spacetimes, and highlight both the strengths and limitations of these previous models. ",
        "conclusions": "Recent observational studies have provided evidence for highly broadened fluorescent iron K$\\alpha$ lines. While there are a variety of line profiles seen (e.g. \\citealt{LZ01}), there are some objects where the line implies that there is material down to the last stable orbit in a maximally spinning Kerr spacetime (most notably MCG-6-30-15: \\citealt{W01}). However, the strong gravity codes generally used to model these effects are now over a decade old. Increased computer power means that it is now possible to improve on these models. We describe our new code to calculate these effects, which uses uses fully adaptive gridding to map the image of the disc at the observer using the analytic solutions of the light travel paths. This is a very general approach, so the code can easily be modified to incorporate different emission geometries.  We compare the results of our new code with those from {\\tt diskline} and {\\tt laor} (publically available in the {\\sc XSPEC} spectral fitting package) for Schwarzchild and extreme Kerr spacetimes.  These previous models are accurate to $\\sim 10\\%$ with realistic ($\\propto r^{-3}$) radial emissivities.  However, they make specific assumptions regarding the angular dependence of the emitted flux, which may or may not be valid. Lightbending is {\\em always} important for a disc which extends down below 20$r_g$, in that the image of the disc at the observer {\\em always} consists of a range of different emission angles. This can produce significant changes to the derived line profile, especially in extreme Kerr spacetimes. Whilst calculating strong gravitational effects is a difficult numerical problem, the underlying physics is well known. By contrast, the {\\em angular} emissivity is an astrophysical problem, and is not at all well known as it depends on the ionization state of the disc as a function both of height and radius. Before we can use the line profiles to provide a sensitive test General Relativity and probe the underlying physics, we will need to have a much better understanding of the astrophysics of accretion.  This code will be publically released for inclusion as a convolution model in the \\texttt{XSPEC} spectral fitting package. This will include arbitrary spin and  inner and  outer disc radii as well as allowing both angular and radial emissivities to be specified.  After this paper was submitted we learnt of the independent work by \\citep{DKY04} which also develops a new strong gravity code. Their results match very closely with those presented here."
    },
    "0402/astro-ph0402220_arXiv.txt": {
        "abstract": "We present the X-ray source catalogues for the {\\em XMM\\/} surveys of the 3-h and 14-h Canada-France Redshift Survey fields ($0.5-10~keV$ flux range $\\sim2\\times 10^{-15} - 10^{-13}~erg~cm^{-2}~s^{-1}$).  We use a subset of the {\\em XMM\\/} sources, which have {\\em Chandra\\/} positions, to determine the best method of obtaining optical identifications of sources with only {\\em XMM\\/} positions.  We find optical identifications for $79~per~cent$ of the {\\em XMM\\/} sources for which there are deep optical images.  The sources without optical identifications are likely to be optically fainter and have higher redshifts than the sources with identifications.  We have estimated `photometric redshifts' for the identified sources, calibrating our method using $\\sim 200$ galaxies in the fields with spectroscopic redshifts.  We find that the redshift distribution has a strong peak at $z\\sim0.7$.  The host galaxies of AGN identified in this work cover a wide range of optical properties with every galaxy type being represented, and no obvious preference for one type over another.  Redder types tend to be more luminous than blue types, particularly at lower redshifts.  The host galaxies also span a wide range of optical luminosity, in contrast to the narrow range found for the starburst galaxies detected in $\\mu Jy$ radio surveys.  We find a strong correlation between optical and X-ray luminosity similar to the Magorrian relation, although selection effects cannot be ruled out. ",
        "introduction": "Deep exposures with the most recent and powerful X-ray observatories, {\\em XMM-Newton\\/} and {\\em Chandra\\/} (e.g. Barger et al. 2003; Giacconi et al. 2002; Mainieri et al. 2002; M$^{c}$Hardy et al. 2003; Page et al. 2003), have built on the deepest {\\em ROSAT\\/} X-ray surveys (e.g. M$^{c}$Hardy et al. 1998; Hasinger et al. 1998) by going deeper and to higher X-ray energies with better positional accuracy. This has opened up the study of faint X-ray sources such as high redshift AGN, and has also revealed X-ray emission from otherwise normal galaxies at more modest redshifts (Hornschemeier et al. 2003). These surveys have now resolved the majority of the cosmic X-ray background (XRB) in the soft ($0.5-2~keV$) X-ray band with a small fraction left unaccounted for in the hard ($2-10~keV$) band (Moretti et al. 2003).  The nature of the XRB at these X-ray energies is well on the way to being understood but the peak in the XRB lies at a much higher energy ($\\sim 30~keV$).  This indicates that a population of very faint sources, with very hard spectra, make up the remaining fraction of the XRB in the hard band, and would also be expected to contribute a much greater fraction to the XRB nearer its peak (Moretti et al. 2003). Such hard sources are most likely a result of extremely high obscuration, which progressively wipes out X-ray emission from low to high energy, turning an intrinsically soft spectrum into a much harder observed one.  The radiation absorbed during this process must be re-emitted at longer wavelengths and the possibility of the Far-IR/Sub-mm background being somehow connected with the XRB is discussed in many papers (e.g. Almaini, Lawrence \\& Boyle 1999).  However, current X-ray/Sub-mm surveys suggest that the two backgrounds are only loosely related (e.g. Waskett et al. 2003; Alexander et al. 2003; Severgnini et al. 2000).  Future instrumentation with higher energy limits are likely required to fully explain the XRB and the nature of the sources that dominate its peak.  At present though, the emphasis must be turned to those sources that we can observe easily with the current instrumentation.  QSOs and type-I AGN dominate the softest X-ray energies with an increasing contribution from more obscured type-II AGN becoming important at higher energies (e.g. Gilli, Salvati \\& Hasinger 2001).  Identifying the optical counterparts to these sources is crucial for a full understanding of their properties and a great deal of effort has been expended in obtaining this information (e.g. Barger et al. 2003; M$^{c}$Hardy et al. 2003).  For example, one of the most useful quantities that can be derived from a source list is the luminosity function.  This reveals much about the nature of a population and determining its evolution with redshift can shed light on how the population as a whole changes over time.  The X-ray luminosity function (XLF) has begun to be investigated in depth by several groups (Cowie et al. 2003; Steffen et al. 2003; Ueda et al. 2003).  Both Ueda et al. (2003) and Steffen et al. (2003) find that the evolution of the XLF is a function of luminosity.  The population of X-ray sources with $L_{X}(2-10~keV)>3\\times10^{43}~erg~s^{-1}$ is dominated by type-I AGN, and the number-density of these sources increases with redshift out to $z\\sim2-3$.  At lower X-ray luminosities however, the fraction of type-II AGN increases rapidly with decreasing X-ray luminosity. The number-density of these sources appears to peak at $z<1$.  Although {\\em Chandra\\/} is better suited for identifying X-ray sources with optical counterparts ({\\em XMM\\/} has a resolution of $\\sim 6\\arcsec$ full width half maximum (FWHM) cf. $\\sim 0.5\\arcsec$ for {\\em Chandra\\/}), {\\em XMM\\/} has greater sensitivity and a larger field of view (FoV), making it better for large area surveys.  In this paper we report the results of a medium-deep {\\em XMM\\/} survey composed of two separate exposures ($\\sim 0.4$ square degrees).  We quantify the ability of such a survey to identify X-ray sources with optical counterparts by comparing the IDs for a subset of the {\\em XMM\\/} sources with the IDs obtained using {\\em Chandra\\/} positions for the same sources.  We estimate redshifts for our identified sources using photometric redshift codes.  These allow a quick, and reasonably reliable, way of obtaining redshifts for objects with multi-band photometry.  Although not as accurate as spectroscopy these techniques are becoming widely used as a short-cut for large surveys, where statistical properties are fairly insensitive to the accuracy of individual redshift measurements (Csabai et al. 2003; Fontana et al. 2000; Kashikawa et al. 2003).  These methods can also be used on objects fainter than the spectroscopic limit, where many X-ray source counterparts reside (Alexander et al. 2001).  We test two photometric redshift estimation codes on our X-ray source IDs and obtain a robust redshift distribution for those sources that could be identified reliably, while placing limits on the properties of those that could not.  Ultimately we will use our identified AGN, and their redshifts, to construct the XLF for different populations, and calculate its evolution with redshift.  The results of this study will be reported in paper-III, the next in this series.  We assume an $H_{0}$ of $75~km~s^{-1}~Mpc^{-1}$ and a concordance Universe with $\\Omega_{M}=0.3$ and $\\Omega_{\\Lambda}=0.7$. ",
        "conclusions": "We have presented source catalogues for a survey, using the {\\em XMM-Newton\\/} X-ray telescope, of $\\sim0.4$ square degrees of sky.  We show that reliable identifications can be obtained for $\\sim75~per~cent$ of the {\\em XMM\\/} sources using {\\em XMM\\/} positions alone.  Those sources that cannot be identified using {\\em XMM\\/} positions alone are optically fainter ($I_{AB}>22$) than most of the identified ones, and are likely to be AGN at generally higher redshifts.  We have obtained the following results:  \\begin{itemize}  \\item The flux ratio $f_{X}/f_{opt}$ of the sources in our survey show that they are predominantly AGN.  \\item The optical properties of the AGN span a large range of absolute magnitudes, in contrast to the result found for the starburst galaxies detected in $\\mu Jy$ radio surveys, which tend to have a very narrow range of absolute magnitudes (Chapman et al. 2003).  \\item AGN are found in host galaxies spanning the full range of Hubble types, with no clear preference.  \\item For the identified X-ray sources with good redshifts there is a strong correlation between optical and X-ray luminosity, reminiscent of the Magorrian relation between black-hole mass and bulge mass. However, this may be due to selection effects.  \\item The redshift distribution of the AGN shows a clear peak at $z\\sim0.7$.  \\end{itemize}  The last result supports other recent studies (Barger et al. 2003) that show the peak formation of super-massive black holes occurred at relatively recent times ($z<1$).  Medium-deep X-ray surveys such as ours, which resolve a large fraction of the XRB but are still dominated by AGN, are able to probe this epoch effectively.  We will use the results from this paper to calculate the X-ray luminosity function, and determine its evolution with redshift, in the next paper in this series."
    },
    "0402/astro-ph0402084_arXiv.txt": {
        "abstract": "We have performed numerical simulations of a radially perturbed ``accretion'' torus around a black hole or neutron star and find that the torus performs radial and vertical motions at the appropriate epicyclic frequencies. We find clear evidence that vertical motions are excited in a nonlinear resonance when the applied perturbation is periodic in time.  The strongest resonant response occurs when the frequency difference of the two oscillations is equal to one-half the forcing frequency, precisely as recently observed in the accreting pulsar, SAX J1808.4-3658, where the observed kHz QPO peak separation is half the spin frequency of 401 Hz. ",
        "introduction": "Millisecond oscillations in the X--ray light curves of systems known to contain neutron stars or possibly black holes have been observed for several years with the Rossi X--Ray Timing Explorer (see \\cite{vdk00} for a review). Recently, it has been pointed out that the centroid frequencies of the corresponding kilohertz quasi periodic oscillations (kHz QPOs) are in rational ratios of small integers, as 3:2 \\citep{ak01,rmmo02,Mc03,abbk03,akklr03}.  This supports suggestions that a resonance of some kind is responsible for the observed properties \\citep{ka01,ka03,t02}.  In addition, for at least one system in which coherent pulsations have been detected, indicating the spin frequency ($\\nu_{s}=401$~Hz in SAX J1808.4-3658), the separation in frequency between the kHz peaks has been recently reported to be consistent with $\\nu_{s}/2$ \\citep{wetal03}, implying that the pulsar is exciting motions in the accretion disk in a nonlinear fashion \\citep{kakls03}.  In this Letter we show that the kHz oscillations detected in SAX J1808.4-3658 can be attributed to forcing of epicyclic motions in the accretion disk by the 2.5 ms pulsar, which induces resonance at selected frequencies. The coupling between the pulsar and the disk could be due to the magnetic field, or to some structure on the surface of the star.  In other, similar systems, a frequency separation equal to the stellar spin frequency is also possible. ",
        "conclusions": "\\label{disc}  The appearance of twin kHZ QPOs in the millisecond pulsar SAX J1808.4-3658, with a separation consistent with half the known spin frequency of the pulsar, strongly indicates that a nonlinear resonance is at work, coupling the spin to vibrational modes in the disk.  Using a simple hydrodynamical model, we identify these modes with the radial and epicyclic oscillations of fluid elements slightly displaced from exact circular orbits. In this respect, the model is crucially dependent on the effects of strong gravity, to break the degeneracy between the orbital and epicyclic frequencies present in the Newtonian regime.  Under the unique assumption that the pulsar provides a periodic driving radial force to a slender torus in orbit (which we consider as a stand--in for a density enhancement in the accretion disk; see \\S~\\ref{response}), we show that the response of the torus is greatest when the spin frequency is twice the difference between the vertical and radial epicyclic frequencies. We also show that there are other possibilities for resonant motion, when the above numbers are in a 1:1 or 3:2 correspondence.  For the case of SAX J1808.4-3658, the frequency ratio 700:500=1.4, and the actual values of the frequencies observed would allow us to fix the mass of the pulsar at 1.38 solar masses in the Schwarzschild metric. However, the actual value will be different, as the epicyclic frequencies for a neutron star rotating at 401 Hz depart from the Schwarzschild values \\citep{kakls03}.  Two further points deserve comment in the context of kHZ QPOs in systems with a millisecond pulsar whose spin frequency is known. First, there is an apparent dichotomy in the values of the QPO peak separation with respect to the spin frequency.  For the ``fast'' rotators, like SAX J1808.4-3658, the separation is half the spin frequency, $\\Delta \\nu=\\nu_{s}/2$, while for ``slow'' rotators, like XTE J1807-294 (where $\\nu_{s}=190$~Hz), the separation is consistent with the spin frequency, $\\Delta \\nu=\\nu_{s}$. This could be explained within the current framework, since for fast rotators, the resonant point corresponding to $\\Delta \\nu = \\nu_{s}$ is so close to the neutron star that is it unlikely to be seen (one could say that there is no room around the pulsar for this mode to occur).  For slow rotators, both resonances, at $\\nu_s$ and $\\nu_s/2$, could in principle be observed, since they would manifest themselves at greater distances from the neutron star. However, if the modulations in flux generated in the innermost regions of the accretion disk have a dominant effect on the overall X--ray light curve, one would then expect to see the strongest signal at frequency $\\nu_s$.  Second, the twin kHz QPOs then need not necessarily occur always at the same frequencies, as the torus shifts its radial position, but their separation should nevertheless remain roughly constant (and equal to $\\nu_{s}$ or $\\nu_{s}/2$, depending on the system) under this excitation mechanism, assuming mode locking to occur.  Further timing observations of millisecond pulsars will no doubt shed light on these matters.  More generally, our simulation shows directly that a radial perturbation present in the torus can excite forced oscillation at (other) eigenfrequencies. This may have implications for mode coupling in black hole accretion disks."
    },
    "0402/astro-ph0402421_arXiv.txt": {
        "abstract": "We present an analytic description of turbulent, magnetohydrodynamic (MHD) disk accretion around black holes that specifically addresses the relationship between radial and vertical, mean-field transport of mass, momentum and energy, thereby complementing and extending numerical simulations. The azimuthal-vertical component of the magnetic stress is fundamental to an understanding of disk--corona--outflow coupling: when it is important for driving the angular momentum transport and mass accretion in the disk, it also has an important influence on the disk--corona--outflow energy budget. The Poynting flux derived from the product of this term with the Keplerian velocity also dominates the Poynting flux into the corona. The ratio of the coronal Alfv\\'{e}n velocity to the Keplerian velocity is an important parameter in disk-corona-outflow physics. If this parameter is greater than unity then energetically significant winds and Poynting flux into the corona occur. However, significant effects could also occur when this parameter is much less than unity. A limiting solution describing the case of angular momentum transport solely by the vertical-azimuthal stress has the property that all of the accretion power is channeled into a wind, some of which would be dissipated in the corona. More realistic solutions in which there is both radial and vertical transport of angular momentum would have different fractions of the accretion power emitted by the disk and corona respectively. These results have important implications for existing accretion disk theory and for our interpretation of high-energy emission and nuclear outflows from the central engines of Active Galactic Nuclei and Galactic Black Hole Candidates. ",
        "introduction": "Since the foundations were laid for a standard theory of disk accretion \\citep{pringrees72,SS73,novthorn73}, two fundamental problems were immediately recognized  \\citep{liangprice77,bisblin77,pacz78}: (1) The efficient transport of angular momentum to large radii cannot be attributed to conventional kinematic viscosity, and (2) The observed high-energy spectra and luminosities and the ubiquity of outflow phenomena from accreting black holes implies efficient vertical transport of energy from a relatively cool, dense disk to a hot, tenuous and unbound corona. While it has been widely accepted that magnetic fields provide the most plausible means of efficiently transporting both angular momentum and energy, the precise nature of this transport has remained unclear until recent numerical simulations demonstrated, unambiguously, that accretion disks work because of magnetohydrodynamic (MHD) turbulence \\citep*{balbhaw98,hawley99}. The turbulence is generated by the magnetorotational instability \\citep[MRI -- ][ and references therein]{balbhaw91}, which is driven by the free energy available from the differential rotation of the bulk flow. Notwithstanding these groundbreaking results, however, the nature of vertical energy transport from an accretion disk to a corona and/or outflow still remains an outstanding and contentious issue.  In the context of Active Galactic Nuclei (AGN), much theoretical effort has recently focussed on the physics of accretion disk coronae \\citep*{dimatt97a,dimatt97b,merlfab01a,merlfab01b,merlfab02,liu02}, commensurate with the dramatic increase in both the quality and quantity of high-energy observational data. However, there has been little improvement in coupled disk--corona models since the first phenomenological descriptions of \\citet{haardt91,haardt93}. Current models \\citep[e.g.][]{merlfab02,liu02} simply replace the  fraction of accretion power transferred from the disk to the corona with a Poynting flux quantity estimated from a mean-field buoyant velocity and an equipartition, mean-field magnetic energy density. While numerical models \\citep[e.g.][]{millstone00} do indeed show that turbulent fluctuations in a vertically stratified disk are capable of driving the magneto-gravitational modes of the Parker instability \\citep{parker55}, whether magnetic buoyancy can supply the corona with sufficient power to explain the observed high-energy emission  is questionable. Numerical simulations indicate that magnetic buoyancy is an ineffective saturation mechanism for the MRI \\citep{brandenburg95,stone96,millstone00}, while theoretical models for disk coronae require implausibly ideal buoyancy conditions and limiting accretion conditions  \\citep[e.g.][]{merloni03}. Realistically, the growth of the unstable buoyant-Parker modes, which is essentially a wave-fluctuation resonance  interaction, must compete against particle-fluctuation interactions, which correspond to dissipation of the turbulence and internal heating of the disk.  The production and ubiquity of outflows from accretion disks around black holes also remains a challenging problem and it is unclear from most theoretical models \\citep*[e.g.][]{blandpayne82,lovelace91,li92,wardkon93,ustyugova00} whether MHD disk turbulence plays a significant role \\citep[but see][]{heinzbeg00}. Nevertheless, numerical simulations of turbulent MHD accretion disks do in fact show that outflows become important in the innermost regions of turbulent accretion disks \\citep*[e.g.][]{stonepring01,hawley01,hawbalb02}. Unfortunately, the numerical models are restricted by their approach: non-Keplerian motions are defined {\\it a priori} as fluctuating quantities, so that vertical, mean-field transport is not self-consistently taken into account. Indeed, the outflows that emerge in some of the simulations \\citep[e.g.][]{hawbalb02} are defined as regions where the nett radial flow is outward, rather than inward. Furthermore, the  neglect of vertical angular momentum transport restricts the radial inflow to substantially subsonic speeds \\citep[see][]{balbhaw98}. Inevitably, the resulting accretion rates in the numerical models are typically very low \\citep[see e.g.][]{stonepring01,hawley01,hawbalb02}. Numerical models are also restricted by computational limitations: simulations that are global as well as vertically stratified are required to estimate the fraction of accretion power that can be vertically transported and this is not only computationally prohibitive, but also sensitive to numerical dissipation effects, which are difficult to quantify.  Thus, the present status of black hole accretion disk theory is that {\\em there is currently no formalism which self-consistently couples turbulent, MHD disk accretion with a magnetically-dominant corona in a framework that can accomodate a range of radial inflow and vertical outflow solutions.}  In this paper, we present the first fully analytic description of a turbulent MHD accretion disk coupled to a corona, self-consistently taking into account both vertical and radial mean-field fluxes of mass, momentum and energy. By deriving the relevant transport equations from first-principles, our formalism provides a non-phenomenological and non-empirical approach to the problem of energy transport to a corona and the associated outflow and the inter-relationship of this transport with accretion onto the central black hole. In this treatment, we focus on the effects of turbulent magnetic stresses and the mean magnetic flux density is assumed to be zero, for the sake of clarity. However, our formalism lends itself naturally to the inclusion of nonzero mean magnetic fields, which we intend to explore separately in a subsequent paper. In \\S~2, we present the relevant conservation equations for a resistive, viscous, and optically-thick MHD gas. In \\S~3, we statistically average these equations and derive the corresponding mean-field equations for a turbulent MHD gas. In \\S~4, we apply these mean-field equations to the dynamics of a geometrically-thin accretion disk that is stationary and axisymmetric in the mean, and we expressly examine the implications of vertical mean-field transport for the conservation of mass and momentum. In \\S~5, we utilize the results of \\S~4 to analyze the total disk energy budget. We conclude with a discussion of the main results in \\S~6. ",
        "conclusions": "In this paper, we have established a self-consistent  framework for the theory of magnetized, turbulent disk accretion around black holes. Our formalism is the first to consistently treat turbulent disks using a robust statistical averaging procedure that explicitly includes dynamical equations for the evolution of the magnetic field together with the conservation equations for mass, momentum and energy transport. We have paid special attention to  vertical transport of conserved quantities and consistently related such transport to the dynamical structure of the underlying disk. Although the nett magnetic flux is assumed to be zero, the formalism is nonetheless sufficiently general to allow the straightforward inclusion of a systematic nett mean-field component.  Our main results can be summarized as follows: \\begin{enumerate} \\item We have derived a comprehensive set of equations describing the transport of  mass, momentum, internal energy, turbulent kinetic and magnetic energies, and total energy. The statistically-averaged conservation equations are completely general and are applicable to other subject areas involving turbulent magnetic fields. \\item  We have applied the statistically averaged equations to a geometrically-thin, optically-thick accretion disk that is stationary and axisymmetric in the mean. We have demonstrated that when vertical transport of mass, radial, vertical and angular momentum, and energy is self-consistently treated, the general equations include additional terms related to a disk wind and turbulent azimuthal--vertical stresses on the disk surface. \\item We have shown that the total azimuthal-vertical stress can have a significant dynamical and energetic effect on the disk even though it may be numerically small compared to the radial--azimuthal stress that has dominated a large amount of accretion disk theory and simulation, to date. Note, however, that the importance of this stress has also been realized in accretion--outlfow models \\citep[see][ for a review]{konpud00}. \\item We have derived an expression for the radiative luminosity from the disk photosphere and shown clearly how this relates to the mechanical power in a wind and to the Poynting flux, thereby identifying the possible sources responsible for powering coronae and/or outflows from accreting black holes. This expression also entails a different distribution of radiative flux than a standard accretion disk. This in turn affects the integrated spectrum. Again, we defer the details to future work. \\item We have discussed the three main sources of Poynting flux into the corona -- a component associated with the product of the  azimuthal - vertical component of the turbulent magnetic stress, a component associated with wind advection of magnetic energy and a component associated with turbulent diffusion of magnetic field from the disk into the corona . The first component probably dominates in most cases even if the azimuthal - vertical stress is quite small in comparison to the radial - azimuthal stress. In the course of the analysis of the condition for a wind and the conditions for a significant Poynting flux into the corona, the ratio of the coronal Alfv\\'{e}n emerges as a critical parameter. When this ratio is of order unity, important magnetic effects are clearly present. However, there is also the prospect of significant effects when this parameter is less than unity. This region of parameter space is currently a relatively unexplored avenue of research in black hole accretion disks. \\item In the limiting case, when all of the angular momentum transport is through the vertical-azimuthal stress, we have shown that the wind power, at the base of the wind, is exactly equal to the  accretion power. Some of this power would be dissipated in the corona. This is the first time that a coupled  disk-corona model has demonstrated, in a physically consistent fashion, the possibility of significant power emanating from the corona. \\item This is also the first time that the power of a disk wind has been dynamically linked to the process of accretion. In models with a nett magnetic flux the wind power is related to the strength of the magnetic field. \\item The existence of a coronal wind and the production of intense coronal emission are inextricably linked. The major influence in heating the corona is the stress that is responsible for transporting angular momentum vertically. The wind is essential to transport this angular momentum away from the disk. \\end{enumerate}  In a realistic disk, we expect that both radial--azimuthal and azimuthal--vertical stresses would be involved in the transport of angular momentum, as well as a nett mass loss from the innermost regions. Nevertheless, our limiting solution provides a good physical basis for the commonly held notion that accretion power could be channelled into significant coronal emissivity and/or outflow comprised of electromagnetic and bulk kinetic components. Thus there is the prospect of explaining not only the coronal emission from radio-quiet AGN and Galactic Balck Hole Candidates (GBHCs), but also systems such as Ultra-Luminous X-ray (ULX) sources and analogous AGN sources such as Broad Absorption Line (BAL) quasars where large mass outflows (e.g. $\\Mdotw \\sim \\dot M_{\\rm Edd}$) are inferred \\citep[see for example][and references therein]{kingpounds03}.  With a self-consistent framework now established, we are in the position of being able to consider further, via specific models, the complex relationship between disk, corona and outflows in a variety of sources, and to examine more thoroughly the conditions for the initiation of a wind and the implications for the general  structure of the immediate environment of accreting black holes."
    },
    "0402/astro-ph0402617_arXiv.txt": {
        "abstract": "We summarize the multiwavelength properties of X-ray sources detected in the 80 ks XMM-Newton observation of the Groth-Westphal Strip, a contiguous strip of 28 HST Wide-Field Planetary Camera 2 (WFPC2) images. Among the $\\approx 150$ X-ray sources detected in the XMM-Newton field of view, 23 are within the WFPC2 fields. Ten spectroscopic redshifts are available from the Deep Extragalactic Evolutionary Probe (DEEP) and Canada-France Redshift Survey (CFRS) projects. Four of these show broad Mg II emission and can be classified as type 1 AGNs. Two of those without any broad lines, nevertheless, have [NeV] emission which is an unambiguous signature of AGN activity. One is a narrow-line Seyfert 1 and the other a type 2 AGN.  As a followup, we have made near-infrared (NIR) spectroscopic observations using the OHS/CISCO spectrometer for five of the X-ray sources for which we found no indication of an AGN activity in the optical spectrum.  We have detected H$\\alpha$+[NII] emission in four of them.  A broad  H$\\alpha$ component and/or a large [NII]/H$\\alpha$ ratio is seen, suggestive of AGN activity. Nineteen sources have been detected in the $K_{\\rm s}$ band and four of these are extremely red objects (EROs; $I_{814}-K_{\\rm s}>4$).  The optical counterparts for the majority of the X-ray sources are bulge-dominated. The $I_{814}-K_{\\rm s}$ color of these bulge-dominated hosts are indeed consistent with evolving elliptical galaxies, { while contaminations from star formation/AGN seems to be present in their $V_{\\rm 606}-I_{\\rm 814}$ color.} Assuming that the known local relations among the bulge luminosity, central velocity dispersion, and the mass of the central blackhole still hold at $z\\sim 1$, we compare the AGN luminosity with the Eddington luminosity of the central blackhole mass. The AGN bolometric luminosity to Eddington luminosity ratio ranges from 0.3 to 10\\%. ",
        "introduction": "Deep X-ray images of a patch of the sky show numerous X-ray sources, which mainly consist of a mixture of absorbed and unabsorbed active galactic nuclei (AGNs). It is now recognized that these AGNs make up the bulk of what has been called the ``X-ray Background''. While normal galaxies, whose X-ray emission is probably dominated by the integration of X-ray binaries, start to emerge \\citep{miy_fluct,horn03} at the faintest fluxes, the dominant X-ray source population in the deep Chandra and XMM-Newton Surveys comes from AGN activities.  Thus multiwavelength studies of these X-ray sources are key to understanding the detailed history and physical conditions of the formation and the growth of the supermassive blackholes (SMBHs), which are now known to reside in the centers of almost all galaxies with a bulge \\citep{magorrian,merritt}.   The region known as the ``Groth-Westphal Strip'' (GWS), consists of Hubble Space Telescope (HST) Wide-Field Planetary Camera 2 (WFPC2) medium-deep images of a strip of 28 contiguous fields\\citep{groth}. It is a particularly useful field for extensive multiwavelength studies. Numerous on-going and future followup projects have been/ are being conducted on and around this field. Existing morphological information from the original WFPC2 observations, combined with the Deep Extragalactic Evolutionary Probe (DEEP) \\footnote{http://deep.ucolick.org/}, Canada-France Redshift Survey (CFRS)\\footnote{http://www.astro.utoronto.ca/$\\sim$lilly/CFRS/} and the ongoing DEEP 2 \\footnote{http://deep.berkeley.edu/} redshift surveys, is providing us with the first clues to the nature of the X-ray source counterparts.  Because of the contamination from starlight in the host galaxy, optical searches for faint AGNs in deep survey fields such as GWS need elaborate efforts. Such attempts have been made by searching for an unresolved component at the centers of galaxies, searching for variable nuclei, and by selecting those with ultraviolet-excess cores \\citep{vicki99,vicki03,beck}.  These surveys reveal up to 10\\% of galaxies as AGN candidates with nuclei extending as faint as M$_B$$\\simeq$ --15. Optical searches, however, are less sensitive to AGNs obscured by dust around the active nucleus. On the other hand, X-ray surveys for AGNs are not hindered by the luminosity of the underlying host galaxy. In particular, X-ray surveys with hard band ($E>2$ keV) sensitivity are also sensitive to the obscured AGNs.  In view of this, we have obtained an 80 ks exposure the northeast part of the GWS with XMM-Newton. This field also has the advantage of a low column density of neutral gas in our galaxy, corresponding to $N_{\\rm H}=1.3\\;10^{20}$ cm$^{-2}$ \\citep{dicky}.  A quick look view and the preliminary Log N-Log S relations from these data have been presented in \\citet{miy_moriond} and \\citet{miy_sant}. In this paper, we present the nature of the 23 X-ray sources in the GWS, where morphological properties from the WFPC2 images and some redshifts are available from the DEEP/CFRS redshift surveys. We further make supplemental near infrared spectroscopic observation for some of these X-ray sources for which we did not observe AGN signatures in the optical spectra. $K_{\\rm s}$ band photometry of the X-ray sources are also presented.  The scope of this paper is as follows. In Section \\ref{sec:xobs}, we describe the X-ray data and analysis, including source detection and spectral analysis. In Section \\ref{sec:opt_ir}, we explain the source of the optical and near infrared (NIR) data. The optical and NIR nature of the XMM-Newton sources and related statistics are presented.  In \\ref{sec:disc}, we discuss the overall results on these X-ray sources and black hole mass and the relationship with the bulge luminosity. We also comment on selected individual sources.  A summary is given in \\ref{sec:sum}. Throughout this paper, we use $H_{\\rm 0}=70\\;{\\rm km\\,s^{-1}\\, Mpc^{-1}}$, $\\Omega_{\\rm m}=0.3$, and  $\\Omega_{\\rm \\Lambda}=0.7$. Unless otherwise noted, $L_{\\rm x}$ is the 2-10 keV rest-frame luminosity in units of ${\\rm erg\\;s^{-1}}$ calculated using the cosmological parameters shown above. ",
        "conclusions": "\\label{sec:disc}  \\subsection{X-ray Source Population}  The 23 sources detected in our 80 ks XMM-Newton observation of the GWS are representative of the X-ray sources that contribute most to the ``Cosmic X-ray Background'' and many of them represent the regime which marks the peak of accretion onto SMBHs in centers of galaxies.  A dominant population in this field consists of AGNs with ${\\rm Log}\\;L_{\\rm x}\\sim$ 44 at $z\\sim 1$. Only a few of them show signs of AGN activity in their optical spectra. Subaru OHS NIR spectroscopy of four of the X-ray sources with no previous optical signature of AGNs revealed  H$\\alpha$+[NII] emission lines showing hints of broad H$\\alpha$ and/or stronger narrow [NII] lines indicative of AGN activity.  The host galaxies of the X-ray sources tend to be bulge-dominated and four are extremely-red objects (EROs) (or Very Red Objects; VROs) ($I_{814}-K_{\\rm s}\\geq 4$). { Also one object (x130) has an upper limit $I_{814}-K_{\\rm s}< 4.8$, which is consistent with being an ERO.} Fig. \\ref{fig:col} shows $V_{\\rm 606}-I_{\\rm 814}$ and $I_{\\rm 814} - K_{\\rm s}$ colors of the X-ray sources in our sample as a function of redshift. We have excluded x33 (no $K_{\\rm s}$ and $V_{\\rm 606}$ photometry available),x46 (X-ray source is off-nucleus, see below), and x83 (Galactic star) from the plot.  Those without redshift information are plotted left of z=0.  For reference, we plot K- and evolution-corrected galaxy colors for elliptical (labeled as E1 in  Fig. \\ref{fig:col}), Sa and Sc galaxies from \\citet{poggianti97}. Model E1 corresponds to Poggianti's model ``E'', which has a star-formation rate with e-folding time of 1 Gyr. We have neglected the difference of cosmological parameters used by us and\\citet{poggianti97}, which corresponds to a $\\sim 10\\%$ difference in the age of the Universe. In plotting these color tracks, we have converted the $I$ and $V$ magnitudes to $I_{\\rm 814}$ and $V_{\\rm 606}$ using the formulae given by \\citet{cow99}. We also plot the colors of an elliptical galaxy model, which we have calculated for our filters and cosmology from the evolving SEDs by \\citet{ka97} (labeled as E2). The plotted model is for a passively evolving galaxy after a short burst at the formation epoch of $z_{\\rm f}=4.62$ and $L=0.1 L_*$ at $z=0$.  The difference between $K$ and $K_{\\rm s}$ magnitudes have been neglected, following \\citet{Cristobal03}.    \\begin{figure} \\epsscale{1.0} \\plotone{z_vs_color.eps} \\caption {The $V_{\\rm 606}-I_{\\rm 814}$ (upper panel) and $I_{\\rm 814}-K_{\\rm s}$ colors of the X-ray sources plotted as a function of redshift. Those without redshift information are plotted left of zero. The meanings of the symbols are the same as those in Fig. \\ref{fig:hr}. One sigma error bars are shown if $\\sigma \\geq 0.1$.  The solid, short-dashed, long-dashed curves are galaxy colors as functions of $z$, based on model spectra for Elliptical (E1), Sa, and Sc galaxies after K- and evolution corrections given by \\citet{poggianti97}. The thick dotted line show the colors of a passively evolved  elliptical model by \\citet{ka97}. The dot-dashed lines are the colors calculated from the mean radio-quiet QSO spectrum by  \\citet{elvis94}. } \\label{fig:col} \\end{figure}     Fig. \\ref{fig:col} shows that the { $I_{\\rm 814} - K_{\\rm s}$ color} of the X-ray sources with bulge-dominated hosts indeed traces that of elliptical galaxies and, in particular, those of EROs with redshifts are roughly consistent with the passively evolving elliptical model (E2). { The $V_{\\rm 606} - I_{\\rm 814}$ color, which is more sensitive to the contaminations from star formation and AGN activities, shows a scatter towards bluer colors from the elliptical regime.} The reddest one (x20), with $I_{\\rm 814}-K_{\\rm s}=5.0$ may be contributed by a dust enshrouded AGN or a starburst. Point-like sources tend to be distributed towards bluer (QSO) colors, although the scatter is large. The scatter may be contributed to by unresolved host galaxies, intrinsic scatter in QSO colors and/or reddening by intrinsic dust absorption.  \\subsection{Bulge Mass and X-ray Luminosity}  The HST imaging of this field has allowed us to find a significant population of X-ray sources at $z\\sim 1$ whose counterparts have a resolved bulge component, either as part of a disk+bulge structure or a pure bulge.  In view of the relationship found between bulge mass and the central blackhole mass in nearby galaxies \\citep{magorrian,merritt}, it is interesting to make a first-order estimation of blackhole mass ($M_{\\bullet}$) from the bulge-component of the host galaxy and compare it with the X-ray luminosity. In the rough estimation below, we take the approach of \\citet{aller} by first converting the bulge luminosity to the central velocity dispersion ($\\sigma$) of the bulge stellar component using an empirical relation. Then we use the tight $\\sigma - M_\\bullet$ \\citep{merritt,tremaine} correlation to obtain the estimated blackhole mass. We use the F814W K-correction for the E galaxy in Fig. 18(d) of \\citet{fukugita} to calculated the F814W absolute magnitude ($M_{\\rm F814W}$).  We also assume an early-type galaxy color of $$ M_{\\rm B_T}-M_{\\rm F814W} = 2.1 $$ \\citep{fukugita,gonzalez}. We then use the relations, \\begin{eqnarray} -M_{\\rm B_T}+5\\log h_{\\rm 70} = 20.5 + 7.7(\\log \\sigma - 2.3),\\nonumber\\\\ M_\\bullet/M_\\sun = 1.48\\; 10^8\\;(\\sigma/200)^{4.65},\\;\\;\\; \\label{eq:mass} \\end{eqnarray} \\citep{gonzalez,merritt} to obtain the estimated blackhole mass. Because the F814W band corresponds to the $B$ band at the source rest frame of $z\\sim 1$, the combination of the K-correction and the magnitude conversion  to $M_{\\rm B_T}$ is insensitive to the assumed galaxy spectral energy distribution.    \\begin{figure} \\plotone{mbhvslx.eps} \\caption {The X-ray luminosity (absorption-corrected) plotted as a function of the central blackhole mass ($M_\\bullet$) estimated from the bulge luminosity (see text) for the 10 X-ray sources in the sample with resolved bulge components. The data points are shown with X-name labels. Three lines correspond to $(L_{\\rm bol}/L_{\\rm Edd})(25/b)=10^{-1},10^{-2}$, and $10^{-3}$, where $L_{\\rm bol}=b\\,L_{\\rm x}$ is the bolometric luminosity of the AGN component and $L_{\\rm Edd}$ is the Eddington luminosity corresponding to $M_\\bullet$. } \\label{fig:mbhvslx} \\end{figure}    Fig. \\ref{fig:mbhvslx} shows the estimated $M_\\bullet$ versus $L_{\\rm x}$ (absorption corrected, see Sect. \\ref{sec:z_and_l}) for 10 X-ray sources in the sample which have a resolved bulge component in the HST WFPC2 F814W image (MDS).  It is interesting to estimate the Eddington ratio $L_{\\rm bol}/L_{\\rm Edd}$ for these X-ray sources, where $L_{\\rm bol}$ is the bolometric luminosity of the AGN and $L_{\\rm Edd}$ is the Eddington luminosity corresponding to the blackhole mass. Writing  $L_{\\rm bol}=b\\;L_{\\rm x}$, where $b\\approx 25$ \\citep{elvis94}, we overplot three lines showing $(L_{\\rm bol}/L_{\\rm Edd})(25/b)=10^{-1},10^{-2},$ and $10^{-3}$ in Fig. \\ref{fig:mbhvslx}. Fig. \\ref{fig:mbhvslx} shows that the estimated blackhole masses for the 10 AGNs range from  $10^7-10^{10}$ $M_\\sun$ and the Eddington ratios from 0.3\\%-10\\%.  These results have interesting implications on how AGN evolve and the current stage of evolution for the AGN represented here ($z\\sim 1$, ${\\rm Log}\\,L_{\\rm x}\\sim 44$).  This is indeed a characteristic redshift and luminosity marking the peak of the accretion history of the universe.  The result that these AGN are typically radiating at a few percent may have important implications on the accretion history and formation of SMBHs. If this is a typical Eddington ratio throughout the AGN phase of these objects, the growth of the blackhole occurs on a timescale of $\\sim t_{\\rm s}/(L_{\\rm bol}/L_{\\rm Edd}) \\sim$ a few $\\times \\;10^9$ yrs, where $t_{\\rm s}\\sim 5\\;10^7 (\\frac{\\epsilon}{1-\\epsilon})(\\frac{0.1}{1-0.1})^{-1}$ yrs is the Salpeter timescale, that is the timescale at which the mass of an object accreting at the Eddington Luminosity grows by a factor of $e$ (for a radiative efficiency $\\epsilon\\sim 0.1$ of a standard accretion disk model). This scenario has difficulty in that the timescale of a few $\\;10^9$ yrs may be too long, while the number density of AGN at ${\\rm Log} L_{\\rm x}\\lesssim 44$ decreases beyond $z\\sim 1$ \\citep{ueda1}. Alternatively, it is possible that these AGN have gone through a brief luminous phase in the past with near-Eddington accretion rates. These may have been observed as more luminous QSOs (${\\rm Log}\\;L_{\\rm x}\\gtrsim 45$) at $z>2$. Yet another more exotic possibility is that they are just accreting with a low radiative efficiency ($\\epsilon << 0.1$), allowing much less time for the SMBH to grow.  We note, however, that there are a number of caveats in interpreting these results and drawing conclusions relating the X-ray AGN evolution and growth of the SMBH. Firstly, the current MDS database shows the analysis for stellar (point-like) images or galaxies (pure bulge, pure disk or bulge+disk decomposition), but only limited analysis \\citep{vicki99} has been completed for decomposing point-like (stellar) nuclei from the host galaxy (stellar+bulge+disk, stellar+bulge or stellar+disk). Thus we select against those with strong AGN components (or with large  $(L_{\\rm bol}/L_{\\rm Edd}$), which are likely to be listed as ``stellar'' in the MDS database or the bulge luminosity in the database may be contaminated by the central AGN component. This situation should be improved in the future, where the HST images are analyzed with point-like nucleus+host galaxy decomposition. Extending this study to other deep fields with X-ray and HST (WFPC2 as well as ACS) data, including the Extended Chandra Deep Field-South (E-CDFS) and the COSMOS field, will be a next logical step. Secondly, a much more fundamental limitation is that we have assumed the local relations in Eq. \\ref{eq:mass} are still valid at $z\\sim 1$. This assumption is not guaranteed to be valid.  \\subsection{Comments on Selected Individual Objects} \\label{sec:indi}  \\begin{description} \\item[x8:] This is a typical type 1 QSO at z=1.22 with a broad Mg II line. \\item[x10:] The DEEP optical spectrum shows many narrow emission lines including high-excitation lines like [NeV]$\\lambda\\lambda 3346,3426$, which are unambiguous indicators of AGN activity. Permitted lines (Mg II, H$\\beta$) are also narrow. Since H$\\beta/$/[OIII]$\\lambda 5007 \\approx 0.7 > 1/3$, these lines are not dominated by a Seyfert 2. It is either a Seyfert 2 whose emission lines are heavily  contaminated by starburst activity or a narrow-line Seyfert 1 galaxy (NLS1) \\citep{oster_pogge95}, where H$\\beta$ is contributed to from the (narrow end of the) broad line region. Because our X-ray spectral analysis shows no X-ray absorption ($N_{\\rm H}<10^{21.5}{\\rm cm^{-2}}$; see Table \\ref{tab:spec}), the NLS1 interpretation is more plausible. \\item[x11:] The optical counterpart has a very bright stellar (point-like) morphology and there is no optical spectrum available for this source to discriminate between a galactic star and a QSO.  However, its X-ray spectrum is inconsistent with a thermal plasma, having a significant residual in the soft part. Also its X-ray to optical flux ratio ${\\rm Log}\\;(f_{\\rm x}/f_{\\rm R})\\approx -0.7$ ($f_{\\rm x}$ is measured in 2-10 keV) is well within the AGN regime (roughly between -1 and 1; see e.g. \\citealt{horn01}). Thus it is most likely to be a QSO. \\item[x20:] Thanks to the Chandra position, we can identify the X-ray source with the brightest bulge-dominated galaxy at $z=1.148$, among four candidates apparently interacting with one another indicated by tidal bridges (See Fig. \\ref{fig:poststamp}). It is an interesting case where galaxy interactions are possibly feeding the AGN activity at this early stage of the universe. It has a QSO luminosity (${\\rm Log} L_{\\rm x}=44.2$), but the optical image is dominated by bulge component of the host galaxy.  The DEEP spectrum shows a broad Mg II line.  The X-ray spectrum shows an absorption of ${\\rm Log}\\;N_{\\rm H}\\sim 22.3$. This is an example of optical type-1 X-ray type 2 AGN. This is also an ERO ($I_{\\rm 814}-K_{\\rm s}=5.0$) and a sub-mm source detected in a deep SCUBA survey \\citep{waskett}. \\item[x22:] No previous optical/IR spectroscopic observations existed for this source. Our Subaru OHS/CISCO observation detected H$\\alpha$ and $[NII]\\lambda\\lambda$ 6548,6583 emission lines, giving z=0.983. Based on the sign of broad H$\\alpha$ and strong $N{II}$, high luminosity ${\\rm Log}\\;L_{\\rm x}\\sim 44.2$. While we mark it as an AGN-IR, it may well be a type 1 AGN. \\item[x28:] This source is the most conspicuous hard X-ray source in the field with an intrinsic absorption of ${\\rm Log}\\;N_{\\rm H}\\sim 22.5$ [cm$^{-2}$].  It is a bulge-dominated galaxy with no optical/IR spectroscopy and a photometric redshift of $z=0.76$ \\citep{im} based on its V--I color alone. This is most likely a Seyfert 2 based on the X-ray properties. \\item[x46:] The position of the source has been determined with Chandra and the relative alignment of the WFPC2 image and X-ray sources has been achieved the three other CXO sources in the field. The X-ray source counterpart is identified with a hot spot just off the patchy irregular starforming galaxy DEEP gss 074\\_2638 (MDS u2ay1:0019) at z=0.432. If we assume that the X-ray source is at the same redshift as this irregular galaxy, the luminosity would be ${\\rm Log}\\;L_{\\rm x}\\approx 42.8$. This is too luminous for an ultraluminous X-ray source (ULX).  We also note that there is a nearby edge-on disk galaxy (DEEP gss 074\\_2237, z=0.156). The X-ray source is 4$\\arcsec$ away (projected distance of 10 kpc) from its nucleus and located towards the direction perpendicular to the disk.  Therefore the  X-ray source could be a ULX associated with the halo of the galaxy. However, even if the X-ray source was at the redshift of this disk galaxy, its luminosity would still be ${\\rm Log}\\;L_{\\rm x}= 41.8$, again well above the ULX regime. One possibility is that this irregular galaxy is undergoing a merging process and the X-ray source is at the nucleus of one of the merging galaxies, or it may simply be a background QSO. \\item[x52:] The DEEP spectrum show no broad lines and the location of MgII shows only absorption.  Other features include [OII]$\\lambda 3727$ and [NeIII]$\\lambda 3869$. The presence of [NeV]$\\lambda 3426$ is suggested but uncertain. We observed this object with Subaru OHS and found a moderately broad H$\\alpha$ (FWHM $\\sim 2000 {\\rm km\\,s^{-1}}$) and a strong [NII] doublet. (Fig. \\ref{fig:ohs}). { The fact that there is no broad MgII line but a moderately broad  H$\\alpha$ suggests that the AGN is obscured by a dust cloud. This is consistent with the hard color of this object, with the second largest HR(2-4.5 keV/0.5-2 keV) in the sample. See Fig. \\ref{fig:hr}.} This is an ERO ($I_{\\rm 814}-K_{\\rm s}=4.4$). \\item[x66:] We find no indication of an AGN in the DEEP spectrum. Like x52, our Subaru OHS observation revealed possible AGN activity through the detection of a moderately broad $H\\alpha$ (FWHM $\\sim 2000 {\\rm km\\,s^{-1}}$) line and a strong [NII] doublet. { This is an obscured AGN similar to x52 and has a hardest HR(2-4.5 keV/0.5-2 keV).} There is also a hint of [NeIII]$\\lambda 3869$, a line which is often stronger in AGN than starforming galaxies. \\item[x69:] A spectroscopic redshift of $z=0.995$ has been determined by our Subaru observation in close agreement with the photometric redshift determined by \\citet{brunner} of $z_{\\rm ph}=0.935$. The Subaru OHS spectrum seems to indicate either strong [NII]$\\lambda 6583$ or broad H$\\alpha$ suggestive of AGN activity. This is an ERO ($I_{\\rm 814}-K_{\\rm s}=4.1$). \\item[x83:] Because of the low X-ray to optical flux ratio (${\\rm Log}\\;(f_{\\rm x}/f_{\\rm R})\\approx -3$) for this bright optical object (F606W=14.8), this is certainly a Galactic star. \\item[x125:] Although the MDS database shows that it is a point source, its color is consistent with an elliptical galaxy at the photometric redshift of $z_{\\rm ph}=1.55$ \\citep{brunner}. The optical counterpart is probably dominated by the host galaxy. \\item[x146:] This object is detected only in the hard (2-8 keV) X-ray band. The DEEP optical spectrum clearly shows [NeV]$\\lambda 3426$ emission and strong [NeIII]$\\lambda 3869$. This is a typical Seyfert 2 galaxy with absorbed X-ray spectrum. \\end{description}"
    },
    "0402/astro-ph0402492_arXiv.txt": {
        "abstract": "{ We derive the amplification of the cosmological magnetic field associated with forming gravitational structure. The self-similar solutions of magnetohydrodynamic equations are computed both in linear and nonlinear regimes. We find that the relatively fast magnetic field enhancement becomes substantial in the nonlinear phase.} ",
        "introduction": "\\label{sec:wstep}   The hypothesis describing the dynamical role  of the primordial magnetic field in the formation and evolution of gravitational structure frequently occurs in the literature (e.g.~Wasserman 1978; Kim~et~al. 1996; Peebles 1995). On the contrary, the inverse trend i.e. the amplification of the magnetic field during density perturbation collapse is not often represented in the context of large-scale structure formation. Common practice is to refer to the constraints for magnetic field amplification set by the density of collapsed matter (e.g.~Zeldovich et al. 1980).  The hints of magnetic field existence on cosmological scales excites interest not only in the absolute value of a frozen field magnification but also in its growth rate. The early nonlinear and previrial phase of the gravitational formation is of particular importance, since it results in several megaparsec structures observed as superclusters or filaments. An understanding of the amplification rate of the structure is obviously related to the explanation of the appearance of sufficiently strong frozen-in magnetic fields expected at this stage of collapse.  We investigate the mildly nonlinear collapse of cylindrical gravitational structure and give the growth rate of the primordial magnetic field as a function of the accretion velocity field. The magnetic field growth proceeds intensively during the phase of fluid compression. The magnetic flux for collapsing plasma is conserved and the magnetic strength changes according to the induction equation. It clearly shows that substantial amplification occurs for strongly compressing flows i.e. for growing $\\divop \\textbf{v}$ --- analogously to the shock processes. The general expression for $\\divop \\textbf{v}$ is obtained thanks to the self-similar form of the hydrodynamic equation. The self-similar presentation of magnetohydrodynamics becomes possible in the case of rapid density contrast and velocity evolution, when $v$ and $\\delta \\propto a^n$ $(n > 1)$ i.e.~when the Lorentz force neglect in the Euler equation is justified.  The plan of the paper is as follows. In Section~\\ref{sec:podstawowe_rownania} we give the magnetohydrodynamic (MHD) equations for cylindrical structures in comoving coordinates of a flat universe. We also discuss the assumptions and symmetries allowing us to separate the induction equation. In Section~\\ref{sec:cylindryczne_perturbacje} we derive the self-similar set of hydrodynamic equations. Its linearization and the subsequent comparison with the known velocity solutions are also given. We obtain the general nonlinear relation $\\divop \\textbf{v}$ versus $\\delta$. The rate of magnetic field amplification is discussed in Section~\\ref{sec:nieliniowe_wzmocnienie} both in the linear and nonlinear regime. In the former we present the analytical expression for amplification and the results of numerical integration in the latter. ",
        "conclusions": "The goal of this paper is to present the amplification rate of the magnetic field associated with the forming gravitational structure of cylindrical symmetry. The widespread conviction that the large-scale structures are filled with microgauss cosmological magnetic fields motivates our interest in amplification processes during their evolution. On the other hand, we see a high degree of filamentarity in the galaxy redshift surveys (e.g. Sathyaprakash et al. 1998). This demonstrates that we deal with magnetized, elongated structures of axial symmetry.  To determine their magnetic structure growth, several simplifications are needed. We used here two categories of simplifying assumptions: physical (i.e. $p \\sim 0$ and $F_L \\sim 0 $), constraining the results to the early nonlinear phase and geometrical ones --- requiring the radial motions and thus the axial fields. On the basis of a formal solution of the induction equation we obtained the exact analytical expression for linear field amplification. The relationship between the density contrast and the magnetic field strength is established through the velocity field divergence. However the major conclusion concerns the nonlinear phase. The magnetic field may be effectively enhanced there. The density contrast growth is stronger than the velocity, achieving the nonlinearity regime earlier. The radial structure of the magnetic field and density contrast are identical, in general --- nonhomogeneous.  Contrary to this highly idealized model, in the realistic situation, the centrifugal forces will stop the collapse. This will however occur in the successive, virialization phase, when the matter will become collisional and then shocked. Therefore, a proper description requires more elaborate application of the fluid model. Within its current limitations the above applied symmetries seem to be less weighty than the physical assumptions. According to previous papers (e.g. Siemieniec \\& Woszczyna 2004, Bruni et al. 2003) the more degenerate, pancake geometry leads to comparable amplification results. Introducing cylindrical symmetry enables us instead to depict the magnetic field profile inside the structure. The substantial enhancement of the matter density accreting onto collapsing structure indicates that significant magnetic fields may be produced in its outer region --- the precursor of the future shock."
    },
    "0402/astro-ph0402171_arXiv.txt": {
        "abstract": "{We report new spectroscopic results, obtained with UKIRT/CGS4, of a sample of 14 candidate ultracool dwarfs selected from the DENIS (Deep Near-Infrared Survey of the Southern Sky) database. A further object, selected from the 2MASS Second Incremental Release, was observed at a later epoch with the same instrument. Six objects are already known in the literature; we re-derive their properties. A further four prove to be very nearby ($\\la$\\,10\\,pc) mid-to-late L-dwarfs, three unknown hitherto, two of which are almost certainly substellar. These findings increase the number of L-dwarfs known within $\\sim$\\,10\\,pc by $\\sim$\\,25\\%. The remainder of the objects discussed here are early L or very late M-type dwarfs lying between $\\sim$45 and 15\\,pc and are also new to the literature.  Spectral types have been derived by direct comparison with {\\it J-,H-} and {\\it K-} band spectra of known template ultracool dwarfs given by Leggett et al.\\thanks{\\tt ftp://ftp.jach.hawaii.edu/pub/ukirt/skl/dL.spectra/} For the known objects, we generally find agreement to within $\\sim$1 subclass with previously derived spectral types. Distances are determined from the most recent M$_{\\rm J}$ vs. spectral type calibrations, and together with our derived proper motions yield kinematics for most targets consistent with that expected for the disk population; for three probable late M-dwarfs, membership of a dynamically older population is postulated. The very nearby L-type objects discussed here are of great interest for future studies of binarity and parallaxes. ",
        "introduction": "The analysis of current near-infrared sky surveys such as 2MASS (Two Micron All Sky Survey; \\cite{skr97}), DENIS (\\cite{epc97}) and the Sloan Digital Sky Survey (SDSS; \\cite{yor00}) is rapidly revolutionising our knowledge of the very low-mass dwarf population in the solar neighbourhood. Observations have required the establishment of new spectral classes (L,T) to characterise the very coolest dwarfs (\\cite{kir99}; \\cite{mar99}) and very recently \\cite{cru03} have increased the number of known M7 -- L6 dwarfs by a further 127\\% (186 new objects), by exploitation of the 2MASS Second Incremental Release. It is the clear goal of current efforts, using existing data, to produce a complete, volume-limited sample of ultracool dwarfs, over the whole sky.  Such dwarfs (of spectral types M7 and later) are likely to have ages of a few Gyr and, with reference to theoretical models (e.g. \\cite{bar03} and references therein) are likely to be substellar. Indeed, as pointed out by \\cite{leg01} any object later than L5 has to be substellar; i.e. incapable of sustaining core hydrogen fusion at any point in its lifetime. The field population of L- and T-dwarfs thus represents an important link to even less massive, younger objects known in nearby star-forming regions.  The DENIS survey has demonstrated its ability to detect very cool stellar objects with the detection of the first L-dwarf populations (\\cite{del97}). To date, 5700 deg$^2$ of survey data have been explored, yielding a sample of 300 ultracool dwarf candidates selected to have $(I-J)$\\,$>$\\,3.0, complete to $I\\sim$\\,18 and reaching $I$\\,=\\,19.0. These objects are plotted in Fig.\\,1~(crosses). Optical spectroscopy of the complete sample is underway and will be published in a future paper. In this Letter, we discuss near-infrared spectroscopy obtained for selected relatively bright ($I\\sim$\\,15--17.5) objects with 3.0\\,$<(I-J)<$\\,4.0. The previously published DENIS L-dwarfs (\\cite{del99}; \\cite{mar99}) have $(I-J)>$\\,3.1.  \\begin{table*} \\begin{center} \\caption[] {Basic and derived data for observed targets. The abbreviated name Dnnnn will be used throughout this paper. Co-ordinates are Equinox J2000 and are given in unabbreviated form. Previously known objects are referenced in the final column. $IJK$ magnitudes are from the DENIS database and have typical errors 0.05--0.1\\,mag. The modified Julian date of the DENIS observation, galactic latitudes, derived spectral types and distances are given in cols. 6, 7, 8 and 9. Spectral types quoted are the mean of those derived from independent $K$- and $H$-band derivations, where both are available (see Col. 10). For 2MJ1112, the $I$-band magnitude is from UKST/SuperCosmos$^1$; $JK$ from 2MASS.} \\begin{tabular}{lllllllllll}\\hline\\hline Name & DENIS-P & $I$ & $J$ & $K$ & MJD & $b$ & Sp. & d/pc & Band & ref.\\\\ \\hline D1048 & J104842.81+011158.2 &  16.2 & 12.9 & 11.5 & 51828.0 & +50.79 & L4   & 9.1$^{-0.9}_{+1.0}$ & $K$  & H02 \\\\ D1411 & J141121.30--211950.6 &  15.5 & 12.5 & 11.3 & 51366.5 & +37.83 & M9   & 16.0$^{-0.9}_{+1.1}$ & $H,K$ & C03 \\\\ D1425 & J142527.97--365023.4 &  17.7 & 13.7 & 11.7 & 51828.0 & +22.32 & L5   & 10.6$^{-1.1}_{+1.2}$ & $K$ & -\\\\ D1456 & J145601.39--274736.4 &  16.4 & 13.2 & 12.2 & 51374.5 & +27.89 & M9   & 22.0$^{-1.3}_{+1.6}$ & $ H,K$ & C03 \\\\ D1510 & J151047.85--281817.4 &  16.0 & 12.8 & 11.4 & 51828.0 & +25.27 & M8   & 21.0$^{-1.5}_{+2.0}$ & $H,K$ & G02 \\\\ D1514 & J151450.16--225435.3 &  17.1 & 14.0 & 12.9 & 51828.0 & +29.14 & M7   & 44.1$^{-4.4}_{+6.0}$ & $K$ & - \\\\ D1539 & J153941.96--052042.4 &  17.5 & 13.8 & 12.4 & 51828.0 & +37.98 & L2   & 19.5$^{-1.5}_{+1.5}$ & $H$ & - \\\\ D1705 & J170548.38--051645.7 &  16.6 & 13.2 & 12.1 & 51698.6 & +20.62 & L4   & 10.7$^{-1.0}_{+1.1}$ & $H,K$ & - \\\\ D2036 & J203608.64--130638.3 &  18.2 & 14.7 & 13.5 & 51828.0 &  --29.07 & M9.5 & 41.7$^{-2.4}_{+2.7}$ & $H,K$ & - \\\\ D2057 & J205754.10--025229.9 &  16.6 & 13.2 & 11.6 & 51786.7 & --29.26 & L1.5 & 16.2$^{-1.1}_{+1.2}$ & $H,K$ & C03 \\\\ D2200 & J220002.05--303832.9 &  16.7 & 13.4 & 12.4 & 51776.7 & --52.54 & L0   & 21.7$^{-1.3}_{+1.4}$ & $H,K$ & - \\\\ D2229 & J222958.15--065043.2 &  18.0 & 14.5 & 13.2 & 51828.0 & --50.80 & M9.5 & 38.0$^{-2.2}_{+2.5}$ & $H,K$ & - \\\\ D2252 & J225210.73--173013.4 &  17.9 & 14.2 & 12.8 & 51435.6 & --60.88 & L7.5 & 8.3$^{-0.5}_{+0.7}$  & $H,K$ & - \\\\ D2254 & J225451.90--284025.4 &  17.4 & 14.1 & 12.8 & 51775.8 & --64.26 & L0.5 & 28.2$^{-1.7}_{+1.8}$ & $H$ & C03 \\\\ \\hline 2MJ1112 & 2MASS J11124910-2044315 & 18.3 & 14.9 & 13.5 & 50930.6 & +36.51  & L0.5: & 41.1$^{-10.2}_{+11.7}$ &  $J$  & -  \\\\ \\hline \\end{tabular} \\end{center} \\begin{flushleft} 1. {\\tt http://www-wfau.roe.ac.uk/sss}; C03: \\cite{cru03}; G02: \\cite{giz02}; H02: \\cite{haw02}\\\\ \\end{flushleft} \\end{table*}  \\begin{figure} \\centering \\includegraphics[angle=-90,width=9cm]{fl221_f1.ps} \\caption{DENIS colour-colour diagram. The objects discussed in this Letter are represented by large circles (L-dwarfs) and large squares (M-dwarfs). Small symbols are the previously published DENIS ultracool dwarfs (\\cite{del99}; \\cite{mar99}). The complete sample of DENIS candidate ultracool dwarfs are plotted as crosses. Note that the diagram excludes objects without a DENIS-$K$ magnitude: some such objects are retained in our overall sample on the basis of their $(I-J)$ colour only.} \\end{figure} ",
        "conclusions": "We present spectroscopic and kinematic data for 15 late M and  L-dwarfs, all but one taken from the DENIS catalogue. Spectral types have been determined by direct comparison to known L-type templates, in both $H$- and $K$-bands. Proper motions derived by comparison of 2MASS and SuperCosmos positions yield transverse velocities consistent with membership of the disk population, at least for all confirmed L-type objects. Three probable very late M-type objects have high transverse velocities ($\\sim$\\,100\\,km\\,s$^{-1}$) and are likely to belong to a dynamically older thick disk population. Nine objects in the sample are hitherto unpublished; three of these, and one further object, are shown to have have spectral types in the range L4--L7.5 and are relatively bright (17.9\\,$<$\\,$I$\\,$<$\\,16.2); hence, if single objects, they are extremely close, $\\la$\\,10\\,pc. This last finding represents an increase in the number of known L-dwarfs likely to be within $\\sim$\\,10\\,pc of the Sun from 12 to 16."
    },
    "0402/astro-ph0402347_arXiv.txt": {
        "abstract": "We have measured the transmission of the Ly$\\alpha$\\ forest produced by neutral hydrogen scattering in the intergalactic medium between redshifts 2 and 6.3 using high signal to noise, high resolution ($R \\ge 5000$) observations of 50 quasars spread over the redshift range. We use a uniform set of $15~{\\rm\\AA}$\\ intervals covering Ly$\\alpha$, Ly$\\beta$, and Ly$\\gamma$\\ absorption regions to tabulate the forest transmission as a function of redshift.  The transmitted fractions show a relatively smooth evolution over the entire range of redshifts, which can be modelled with a smoothly decreasing ionization rate. Previous claims of an abrupt change at $z \\sim 6$\\ appear in part to be a consequence of an incorrect conversion of Ly$\\beta$\\ to Ly$\\alpha$\\ optical depths.  The tabulated transmissions can be used to calculate the colors of objects with a specified input spectrum as a function of redshift.  We calculate the colors of a flat $f_{\\nu}$\\ galaxy with a large intrinsic continuum break, as an important example. ",
        "introduction": "\\label{intro}   The epoch of hydrogen reionization is one of the landmarks of the high redshift universe, providing constraints on the first light sources to have formed and a crucial input to models of structure formation. Recent dramatic progress has been made in determining limits on this redshift, leading to an interesting observational situation.  On the one hand, spectroscopy of the Lyman $\\alpha$\\ forest shows that the intergalactic medium (IGM) is highly ionized at $z < 6$, becoming progressively opaque to radiation blueward of the quasars' Ly$\\alpha$\\ emission, culminating in the spectrum of the $z = 6.28$\\ QSO SDSS~1030+0524 which exhibits a substantial dark region that is consistent with no transmitted flux, i.e. a Gunn-Peterson trough at $z = 6.05$\\ (Becker et al.\\ 2002, Songaila \\& Cowie 2002; Fan et al.\\ 2002).  This high opacity has been used to infer that this redshift is at the tail end of the hydrogen reionization epoch (Becker et al.\\ 2002; Djorgovski et al.\\ 2002) and comparison with simulations suggest the IGM could have been fully neutral at $z = 6.5$\\ (Cen \\& McDonald 2002; Gnedin 2002; Fan et al.\\ 2002) though, as we point out in this paper, some of the strongest claims in this regard are based on an incorrect conversion of Ly$\\beta$\\ opacities to equivalent Ly$\\alpha$\\ opacities (Becker et al.\\ 2002; White et al.\\ 2003).  On the other hand, the recent detection by WMAP of a large optical depth to electron scattering (Bennett et al. 2003) is consistent with the Universe having been reionized at $z \\sim 15$\\ and fully ionized thereafter.  The exact significance of the Gunn-Peterson trough in SDSS 1030+0524 is still not clear.  It could follow from the ongoing thickening of the Ly$\\alpha$\\ forest rather than representing the onset of reionization since it is extremely difficult to distinguish the two effects.  The discovery of three new $z > 6$\\ SDSS quasars (Fan et al.\\ 2003) has made the situation more complex since, as might have been expected, there is significant variance in opacity between different lines of sight at $z \\sim 6$ (Fan et al.\\ 2003; White et al.\\ 2003).  Although White et al.\\ (2003) have argued that apparent increased transmission in the spectrum of SDSS 1148+5251, currently the highest redshift quasar, could be a result of foreground emission, it is more likely to be real cosmic variance (Songaila 2004a).  In assessing the significance of the opacity at $z = 6$\\ we need to know how the ionization rate in the IGM is evolving at lower redshifts, since the strongest signature of reionization would be an abrupt change in ionization at some redshift.  In this paper, we use high signal-to-noise spectra of a large sample of quasars, including all the brightest $z > 4.5$\\ quasars observable from the northern hemisphere, to investigate the evolution of the transmission in the IGM from $z = 2$\\ to $z = 6.3$.  We find that the transmission is smoothly evolving with redshift throughout the range. ",
        "conclusions": "\\label{conc}  We have measured the transmission of the Ly$\\alpha$\\ forest produced by neutral hydrogen scattering in the intergalactic medium between redshifts 2 and 6.3 using high signal to noise, high resolution ($R\\ge 5000$) observations of 50 quasars spread over the redshift range.  The transmitted fractions show a relatively smooth evolution over the entire range of redshifts, which can be modelled with a smoothly decreasing ionization rate.  We have used the tabulated transmissions to calculate the colors of a flat-spectrum galaxy with a large intrinsic Lyman continuum break."
    },
    "0402/astro-ph0402037_arXiv.txt": {
        "abstract": "IRAS\\,16279$-$4757 belongs to a group of post-AGB stars showing both PAH bands and crystalline silicates.  We present mid-infrared images, that resolve the object for the first time.  The morphology is similar to that of the `Red Rectangle' (HD\\,44179), the prototype object with PAHs and crystalline silicates.  A two-component model and images suggest a dense oxygen-rich torus, an inner, low-density carbon-rich region and a carbon-rich bipolar outflow.  The PAH bands are enhanced at the outflow, while the continuum emission is concentrated towards the center.  Our findings support the suggestion that mixed chemistry and morphology are closely related.  We discuss the ISO/SWS spectra of IRAS\\,16279$-$4757.  Several bands in the ISO/SWS spectrum show a match with anorthite: this would be the first detection of this mineral outside the solar system.  Compared to HD\\,44179, the shapes of PAH bands are closer to those of planetary nebulae, possibly related to a population of small PAHs present HD\\,44179, but absent around IRAS\\,16279$-$4757.  Detailed examination of the spectra shows the individual character of these two objects.  The comparison suggests that the torus found in IRAS\\,16279$-$4757 may have formed more recently than that in HD\\,44179. ",
        "introduction": "Some post-AGB stars show both PAH and crystalline silicate bands in their infrared spectra.  The formation history of this mixed chemistry (oxygen-rich silicates versus carbon-rich PAHs) is not well understood.  A possibility, but implausible, is that these stars all evolved from oxygen-rich to carbon-rich within the last few hundred years \\citep{Zijlstra91}.  \\citet{Waters98} and \\citet{Molster99} propose that the silicate dust is stored in a long-lived circumbinary disk.  In this scenario, the PAHs form during a later mass-loss phase, after the star became carbon-rich, while the gas stored in the disk retains the chemistry of the earlier, oxygen-rich phase.  Part of the amorphous silicate dust crystallizes in the disk.  The scenario explains why there are relatively few post-AGB stars with this mixed chemistry, and why the silicates have a lower temperature than the carbon rich dust. It requires all such stars to be binaries, as is the case for the prototype of the class, the Red Rectangle (HD\\,44179; \\citet{Waelkens96}).   IRAS\\,16279$-$4757 (hereafter IRAS\\,16279) is classified as a post-AGB star, based on a double peak in its spectral energy distribution indicating a detached envelope \\citep{vanderVeen89}. Optical spectra suggest a spectral type of G5 \\citep{Hu93}.  The ISO/SWS spectra show crystalline silicates beyond 20\\,$\\mu$m \\citep{Molster99} and PAH bands are seen in the near-infrared \\citep{vanderVeen89} and ISO spectra.  IRAS\\,16279 is therefore a member of the group of mixed-chemistry post-AGB stars.   We present TIMMI-2 mid-infrared imaging and spectroscopic data, resolving IRAS\\,16279 for the first time.  In this paper we discuss the spatial distribution of the different dust components based on these images, and compare these with predictions from the circumbinary disk scenario.  We compare the spectra of this object with those of the prototype mixed-chemistry object, the Red Rectangle. ",
        "conclusions": "The mid-infrared images of IRAS\\,16279 show an elongation in the PAH band and a rectangular shape in the N11.9 and Q-band images.  At the center, the PAH-to-continuum ratio decreases.  There is some resemblance to the Red Rectangle (HD\\,44179) \\citep{Waters98}. In the Red Rectangle, the silicates are thought to be located in a disk and the PAHs are in the perpendicular polar flows. The images of IRAS\\,16279 strengthen this link between morphology and mixed chemistry as proposed by \\citet{Waters96} and \\citet{Molster99}.   The temperature of the crystalline silicate bands in HD\\,44179 is about 135\\,K (enstatite; \\citet{Molster02b}), while it is below 110\\,K in IRAS\\,16279.  Nevertheless, both objects show a lower temperature in the silicates than in the PAH region.  In fact, all crystalline silicates in post-AGB stars tend to show low temperatures (100--250\\,K) \\citep{Molster02b}.   In the case of IRAS\\,16279, the model suggests that the PAH emission in the SW and NE regions (Fig.\\ref{Fig-images}) cannot be shielded by the dusty oxygen-rich regions. The oxygen-rich material should be configured as an torus rather than a shell, allows radiation to leak towards the SW and NE directions.  The model explains the PAH excess at 2 arcsec: the decreasing temperature suppresses the continuum but not the PAH bands.   This structure, with a low-density carbon-rich region and an obscured, dense oxygen-rich torus, agrees with the (Red Rectangle) oxygen-rich-disk scenario of \\citeauthor{Waters96} and \\citeauthor{Molster99} But the oxygen-rich region in IRAS\\,16279 is several times larger than seen in the Red Rectangle.  A circumbinary disk is expected to be compact in order to store the oxygen-rich gas over a long time.  It is not clear whether this has happened in IRAS\\,16279.  The extended torus should disrupt faster than the disk of the Red Rectangle. Its formation may have happened relatively recent.       The crystallization of amorphous silicates occurs via heating and subsequent cooling of the grain.  This may occur slowly at low temperature in a long-term stable disk, under the influence of UV radiation (as in the Red Rectangle), or quickly in the AGB wind at very high mass-loss rates, through high temperature annealing \\citep{Waters96, Sylvester99}.  Assuming an expansion velocity of 20\\,km\\,s$^{-1}$, the mass-loss rate of IRAS\\,16279 was of order $5 \\times 10^{-4}\\,\\rm M_\\odot\\,yr^{-1}$, sufficient for high temperature annealing in the outflow to occur. Therefore, crystallization itself may occurred already in the AGB outflow.  Part of the crystallized silicate is stored in the disk later, and part of it may remain in the outflow.   It is not clear that the torus is stable long enough for long-term crystallization.  This may be solved by shocks addressed by \\citet*{Harker02} who found a 5\\,km\\,s$^{-1}$ shock is sufficient for annealing comets. This velocity range might be possible in AGB or post-AGB wind interaction with the torus.  In this case, the higher density of the torus is more likely to obtain a higher rate of crystallization, and this may compensate the short life of the torus than the disk.     The crystalline silicate bands in IRAS\\,16279 are significantly weaker than in the HD\\,44179 \\citep{Waters98}.  The weakness of the features indicate that either they are less abundant than in HD\\,44179, or the temperature difference with the hotter amorphous silicates (responsible for the continuum) is larger in IRAS\\,16279. Although a difference in abundance can clearly not be excluded, the low temperature of the crystalline silicates seems to be supported by the absence of detectable features below 30\\,$\\mu$m."
    },
    "0402/astro-ph0402201_arXiv.txt": {
        "abstract": "The aspects of the analysis of  photometric time--series obtained on double--mode or multiperiodic pulsating stars are briefly reviewed. In particular, the ratios between frequencies are used to pin cases revealing peculiarities. In addition to the Petersen diagrams, we also demonstrated that the period ratios can detect interesting objects. In particular, new results are obtained on High--Amplitude Delta Scuti contained in the OGLE-II database. ",
        "introduction": "In the recent years a huge collection of time--series has been obtained on variable stars, as a noticeable by--product of several microlensing projects. Therefore, the investigation of thousands of light curves is carried out by detecting the frequencies present in the time--series. It is not easy to give astrophysical depth to this kind of analysis, as we have monochromatic photometric data only at our disposal. Different classes of variable stars can show very different physical processes with very similar light curves. Moreover, instrumental terms are often superimposed to physical ones in the same frequency range. In this paper we will try to review some of these aspects. ",
        "conclusions": ""
    },
    "0402/astro-ph0402215_arXiv.txt": {
        "abstract": "{ We report on spectral and timing analysis of \\BSAX data of the 13.6\\,s period transient X-ray pulsar \\exo. Observations were carried out in March 1997 and October 1998, catching the source during a high and a low emission state, respectively. Correspondingly, the X-ray luminosity is found at a level of $ 4.2 \\times 10^{37} \\ergs $ and $1.5\\times 10^{36} \\ergs $ in the two states.  In the high state the X-ray emission in the energy range 1--100\\,keV is well fitted by an absorbed power--law with photon index $ \\Gamma\\sim 1.7$ plus a blackbody component with a characteristic temperature of $\\sim 3.5 $\\,keV. Moreover, we find an evidence of an iron emission at $\\sim$6.8\\,keV, typical feature in this class of sources but never revealed before in the \\exo\\, spectrum. In the low state an absorbed power--law with $\\Gamma\\sim 0.4$ is sufficient to fit the 1--10\\,keV data.  During \\BSAX observations \\exo\\, display variations of the pulse profile with the X-ray flux: it showed single peaked and double peaked profiles in the low and high state, respectively.  Based on these two observations we infer a spin--up period derivative of $- (1.14\\pm0.08)\\times 10^{-10}ss^{-1}$. By comparing these with other period measurements reported in literature we find an alternating spin-up and spin-down behaviour that correlates well with the X-ray luminosity. \\keywords { stars: individual: EXO 053109-6609.2 --- stars: neutron --- stars:binaries --- X-rays: stars } } ",
        "introduction": "\\exo\\, is a High Mass X-ray Binary (HMXB) hosting a neutron star and a B--emission spectral type (Be) star (Haberl, Dennerl, \\& Pietsch 1995; McGowan \\& Charles 2002).  The system is in the Large Magellanic Cloud (LMC), $\\sim 17^{\\prime}$ away from LMC X-4, another high mass X-ray binary pulsar (La Barbera et al. 2001).  Be stars are characterised by high rotational velocities (up to 70\\% of their break--up velocity), and by episodes of equatorial mass loss which produce a temporary ``decretion'' disk around the star (Slettebak 1987, Okazaki \\& Negueruela 2001). These X-ray binaries in which a neutron star is orbiting in a relatively wide orbit with moderate eccentricity are the most numerous among HMXBs. In the LMC, more than half of the confirmed HMXBs are variable or transient sources consisting of a neutron star with a Be companion (Haberl, Dennerl, and Pietsch 2003; Sasaki, Haberl \\& Pietsch 2000; Haberl \\& Sasaki 2000).  Variable X-ray emission is often observed which likely results from the large variations in the Be wind density and relative velocities along the neutron star orbit, which in most cases may results in regular X-ray outbursts near the periastron (Type I outbursts). Alternatively, aperiodic outbursts occur which often last longer than the neutron star orbit. These are probably caused by matter ejection outflowing from the equatorial plane of the Be star (Type II outbursts; Stella, White \\& Rosner 1986; Motch et al. 1991).  \\exo\\, was discovered in 1983 deep EXOSAT exposures of the LMC X-4 region (Pakull et al. 1985; Pietsch, Dennerl, and Rosso 1989). The luminosity was $\\sim 6 \\times 10^{36} \\ergs $ (0.15--4 keV; assuming a distance of 50 kpc). Subsequently, the coded mask X-ray telescope SL2 XRT flown on board of the shuttle Challenger between July and August 1985 detected a second outburst (Hanson et al. 1989), with a source luminosity of $\\sim 1 \\times 10^{37} \\ergs $ (2--10\\,keV).  The source was then monitored from June 1990 to July 1994 with the ROSAT PSPC (Haberl, Dennerl, \\& Pietsch 1995). Alternating high states ($L_{x}\\sim 10^{37}\\ergs $) and low states ($L_{x}\\sim 10^{36}\\ergs $) in the X-ray flux were discovered.  Haberl et al. 1996 detected another outburst from March to May 1993 with an average luminosity $\\sim 2.4 \\times 10^{36} \\ergs $ (0.1--2.4\\,keV).  Dennerl et al. (1996) reported on a ROSAT observation that took place in October 1991 and led to the detection of a spin period of 13.67133(5)\\,s with a period derivative of $ (1.5\\pm0.1)\\times10^{-8} ss^{-1}$ (calculated during their observation).  Haberl et al. 1995 proposed an orbital period of about 600--700 days and an orbital eccentricity of $e \\sim0.4-0.5$.  Under the assumption that period changes are caused by Doppler shifts, Dennerl et al. (1996) corrected the proposed orbital period, finding an orbital solution with $ P_{orb} = 25.4 $ days and an eccentricity of $ e \\sim 0.1 $.  Timing analysis of the \\BSAX\\ data obtained in March 1997 when the system was in high--state, revealed coherent pulsations at a period of 13.67590(8)s and two different period derivatives: a short-term period derivative, calculated during the 2 days of BeppoSAX observation, of $\\dot{P}_{loc}=(3.7\\pm0.5)\\times10^{-9} ss^{-1}$, and a secular period derivative, calculated from a comparison with a previous ROSAT measure of $\\dot{P}_{sec}=(3.67\\pm0.05)\\times10^{-11} ss^{-1}$ (Burderi et al. 1998).  Last observation of \\exo\\, was a deep XMM-Newton observation of LMC field of October 2000 (Haberl, Dennerl and Pietsch 2003) in which \\exo\\, was found to be the brightest source in the field ($7\\times10^{-12}\\ergscm2$ in 0.2-10\\,keV range, corresponding to a luminosity of about $2.1\\times10^{36}\\ergs$).  Across different observations this source varied its intensity up to a factor of about 10. Nevertheless, notwithstanding most of Be X--ray binary systems show a quiescent state, this source has not yet been detected in quiescence but only in a low--luminosity state with a minimum luminosity of $\\sim 10^{35}-10^{36}\\ergs$.  In this paper we report on spectral, spin period and pulse profile changes with the X-ray luminosity. We finally compare our results with previous findings.   \\begin{figure}[htb] \\centerline{\\psfig{figure=fig1.ps,width=8cm,height=4cm,angle=270} } \\caption{Epoch folding search of the data from the low state observation of 1998.} \\end{figure}  \\begin{figure}[htb] \\centerline{\\psfig{figure=fig2.ps,width=8cm,height=6cm,angle=270}} \\caption{Evidence of a quadratic component in the phase--fitting of the high--state observation (linear component, corrisponding to the local $\\dot{P}$ was removed in order to better show the quadratic residuals).} \\end{figure} ",
        "conclusions": "\\begin{figure*}[htb] \\centerline{\\psfig{figure=fig8.ps,width=10cm,height=10cm}} \\caption{Spin period and luminosity changes between ROSAT (Haberl et al. 1996), \\BSAX\\, (this paper) and XMM (Haberl et al. 2003) observations (all reported luminosities are extrapolated in 0.1--2.4\\,keV band).} \\end{figure*}  During the outburst phase of Be X--ray binaries, accretion disks are expected to be present, and indeed, evidence for an accretion disk, based on a correlation between the observed flux and spin--up or spin--down rate , has been found for several sources (Bildsten et al. 1997; Okazaki \\& Negueruela 2001).  The \\exo\\, spin period was detected four times: the first time by ROSAT (Dennerl et al. 1996), two times by \\BSAX\\, (Burderi et al. 1998 and this paper) and the fourth time by XMM-Newton satellite (Haberl et al. 2003).  In Fig.5 we plot the 0.1--2.4\\,keV band luminosity versus period for all these four observations; it is evident from the plot that periods and luminosities are directly correlated.  Comparing the spin periods of all the observations of \\exo\\, carried out from its discovery, we found that the source alternates spin--up and spin--down (Fig.5).  The secular period derivative between the ROSAT observation (Haberl, Dennerl, and Pietsch 1995) and the first \\BSAX\\, observation is $(2.9\\pm0.1)\\times 10^{-11}ss^{-1}$; comparing the two \\BSAX\\, observations we obtain $-(1.14\\pm0.08)\\times 10^{-10}ss^{-1}$, and the $\\dot{P}$ between the last \\BSAX\\, observation and the XMM observation (Haberl et al. 2003) is $-(3.7\\pm0.1)\\times 10^{-11}ss^{-1}$ (period derivative errors are at 90\\% confidence level).  The secular spin period derivative is changed from a spin--down to a spin--up trend, which has a clear correlation with the X-ray luminosity. We detected this change around 1997.  Double peaked pulse profile (Fig.3, first and second panels) was observed when the source was in high--state, while a single peaked profile was found in the low--state (Fig.3, third panel). This behaviour has been often seen in X--ray binary systems and has been ascribed to the transition of the source between a pencil beam and a pencil plus fan beam emission geometry, a behaviour which is correlated with the X--ray flux (Parmar et al. 1989a and 1989b). Timing behaviour of this source is similar to what seen in other Be X--ray binaries, making the period derivative changes a peculiar characteristic of these systems.  The spectrum in the high--state of emission is well fitted by an absorbed blackbody plus a power law and a Gaussian line at $\\sim$ 6.8\\,keV (see Tab.1 and Fig.4 first panel). The 6.8 keV emission line is probably due to K-shell emission from highly ionized iron (probably in the He-like and/or H-like ionization stages). This is the first evidence of an iron emission line in this source; iron emission lines are a common feature in HMXBs, although their centroid energies are usually detected in the range 6.4 - 6.7 keV (e.g. Parmar et al. 1989a).   From measured temperature of kT$\\sim3.5$ keV we infer an emission radius for the blackbody component of $\\sim 5.5$ km, about half neutron star radius.  The calculated blackbody radius is smaller than the neutron star radius, probably due to the fact that the thermal emission does not come from all neutron star surface but probably just from a small region around the polar caps. The energy flux of the two spectral components differs of about two orders of magnitude, $F_{bb}/F_{power} = 1.728\\times10^{-2}$.  The blackbody component probably disappear while the system is in the low--state and the power--law photon index becomes flatter (see Tab.1 and Fig.4 second panel).   \\vspace{1cm} We thank T.Mineo for useful helps with LECS response matrices. This work is born during an Astrophysical Laboratory of Prof. R.Buonanno in the University of Rome ``Tor Vergata''. This work is supported through ASI, CNR and Ministero dell'Universit\\`a e Ricerca Scientifica e Tecnologica (MURST-COFIN) grants."
    },
    "0402/astro-ph0402023_arXiv.txt": {
        "abstract": "We present a low energy expansion of the Kramers-Heisenberg formula for atomic hydrogen in terms of $(\\omega/\\omega_l)$, where $\\omega_l$ and $\\omega$ are the angular frequencies corresponding to the Lyman limit and the incident radiation, respectively. The leading term is proportional to $(\\omega/\\omega_l)^4$, which admits a well-known classical interpretation. With higher order terms we achieve accuracy with errors less than 4 \\% of the scattering cross sections in the region $\\omega/\\omega_l\\le 0.6$. In the neighboring region around Ly$\\alpha$ ($\\omega/\\omega_l >0.6$), we also present an explicit expansion of the Kramers-Heisenberg formula in terms of $\\Delta\\omega\\equiv (\\omega-\\omega_{Ly\\alpha})/\\omega_{Ly\\alpha}$. The accuracy with errors less than 4 \\% can be attained for $\\omega/\\omega_l \\ge 0.6$ with the expansion up to the fifth order of $\\Delta\\omega$. We expect that these formulae will be usefully applied to the radiative transfer in high neutral column density regions, including the Gunn-Peterson absorption troughs and Rayleigh scattering in the atmospheres of giants. ",
        "introduction": "Hydrogen is the most abundant element in the universe and therefore one may often encounter astronomical situations associated with the radiative transfer in a region with a very high neutral hydrogen column density $N_{HI}$. The extended atmosphere around a giant star is such an example, where near UV photons can be significantly scattered by atomic hydrogen (e.g. Isliker, Nussbaumer \\& Vogel 1989). Another example may be found in searches for the first objects that are responsible for the reionization of the universe. When the universe is still partially neutral before the completion of the reionization process, radiation should go through regions with a very high neutral hydrogen column density $N_{HI}$.  In particular, the interactions with low energy electromagnetic waves with angular frequency $\\omega$ are simply Rayleigh scattering, where the scattering cross section is known to be proportional to $\\omega^4$ in the limiting case where $\\omega$ is much smaller than the angular frequency $\\omega_{Ly\\alpha}$ corresponding to the Ly$\\alpha$ transitions.  In quantum mechanics, the Rayleigh scattering process is described by a second-order time dependent perturbation theory, where the scattering atom suffers level transition twice, one associated with the annihilation of the incident photon and the other associated with the creation of the scattered photon. The scattering cross section is known as the Kramers-Heisenberg formula, which is obtained by combining an infinite sum over all the bound $np$ states with the energy $E=-E_0/n^2$ and an integral over all continuum states $n'p$ with the energy $E=E_0/n'^2$, where $E_0$ is the Rydberg energy (e.g. Sakurai 1967). Since the wave functions are analytically known for a single electron atom, the Kramers-Heisenberg formula for hydrogen can be written explicitly in a closed form.  However, the Kramers-Heisenberg formula is unwieldy due to the presence of the infinitely many atomic levels contributing to the cross section.  In the case of a hydrogen atom in the ground state interacting with incident radiation with $\\omega$ much less than $\\omega_{Ly\\alpha}$, this inconvenience can be overcome by expanding the Kramers-Heisenberg formula in terms of $\\omega/\\omega_{Ly\\alpha}$. Devoid of any resonance in the red region of Ly$\\alpha$, the scattering cross section is a well-behaved monotonic function of $\\omega$, and the leading term is proportional to $\\omega^4$, which admits an immediate classical interpretation. Because $\\omega^4$ dependence is the limiting behaviour of $\\omega/\\omega_{Ly\\alpha} \\ll 1$, inclusion of higher order terms will be useful to obtain more accurate cross section values in the red vicinity of Ly$\\alpha$. However, very near the Ly$\\alpha$ resonance, the cross section is well approximated by a Lorentzian and hence a polynomical approximation becomes poor. Lee (2003) introduced an expansion of the Kramers-Heisenberg formula near Ly$\\alpha$, by computing the deviation from the Lorentzian.  In this paper, we will provide and compute the accuracy of the expansions that provide the Rayleigh scattering cross section redward of Ly$\\alpha$. ",
        "conclusions": "In this paper, we obtained an expansion in terms of $\\omega/\\omega_l$ of the Rayleigh scattering cross section by atomic hydrogen, which is applied in the low energy regime with $\\omega/\\omega_l<0.6$. By combining this with another expansion of the Kramers-Heisenberg formula around the Ly$\\alpha$ resonance in terms of $\\Delta\\omega=(\\omega-\\omega_{Ly\\alpha})/\\omega_{Ly\\alpha}$, we may have a wieldy and useful approximate formula for the Rayleigh scattering process redward of Ly$\\alpha$ by atomic hydrogen, which can be made arbitrarily accurate by inclusion of higher order terms directly calculated from the Kramers-Heisenberg formula.  Rayleigh scattering by atomic hydrogen is important only in the presence of a scattering region with a very high neutral hydrogen column density $N_{HI}$.  Such high column density media may be found in an extended atmosphere of a giant star where the mass loss process is already very important.  Isliker et al. (1989) considered the effect of Rayleigh scattering in binary systems containing a giant star. In these systems, for a given inclination and density distribution, the scattering optical depth is dependent on the wavelength, and therefore light curves differ according to the observed wavelength. This information may be quite important to investigate the mass loss process from a giant star.  In their analysis, Isliker et al. (1989) presented the Rayleigh scattering cross section given by \\begin{equation} \\sigma(\\omega) = \\sigma_T \\left[\\sum_{k=2}^\\infty {f_{1k}\\over \\left({\\omega_{1k}\\over \\omega}\\right)^2-1}\\right]^2, \\end{equation} where $f_{1k}$ is the oscillator strength between $1s$ and $kp$ states. In their work, they neglected the contribution from the continuum states. However, as is noted in the previous section, the contribution from the continuum states to the low energy regime is not negligible, and hence caution should be exercised.  In more than half of the symbiotic stars, Raman scattered O~VI 6827, 7088 features are seen, which are formed via Raman scattering of O~VI 1032, 1038 resonance doublet by atomic hydrogen. Being slightly less energetic than Ly$\\beta$, O~VI 1032, 1038 doublet may excite a hydrogen atom that can subsequently de-excite to excited $2s$ state re-emitting an optical photon redward of H$\\alpha$. This process was first identified by Schmid (1989), where the scattering cross section is of similar order to that for Rayleigh scattering (e.g. Lee \\& Lee 1997, Nussbaumer,  Schmid, \\& Vogel 1989).  A very high column density media may be found in an early universe when the reionization of intergalactic medium initiated by the first objects was not completed. In this partially ionized universe, a significant extinction around Ly$\\alpha$ is expected, which will result in a big absorption trough known as the Gunn-Peterson effect (Gunn \\& Peterson 1965, Scheuer 1965). Thus far several quasars with redshift $z>6.2$ have been idenfied with Gunn-Peterson troughs (Becker et al. 2001, Fan et al. 2003).  It appears that the H~I column density that is responsible for these Gunn-Peterson troughs are not sufficiently high for applications of our current work, but may be high enough to see the deviation from the Lorentzian approximation of the scattering cross section (see Lee 2003, Miralda-Escud\\'e 1998). Deeper IR search for higher redshifted objects may exhibit extremely high neutral hydrogen column density, where an accurate calculation of the cross section is required."
    },
    "0402/astro-ph0402509_arXiv.txt": {
        "abstract": "I present a new census of the members of the Chamaeleon~I star-forming region. Optical spectroscopy has been obtained for 179 objects that have been previously identified as possible members of the cluster, that lack either accurate spectral types or clear evidence of membership, and that are optically visible ($I\\lesssim18$). I have used these spectroscopic data and all other available constraints to evaluate the spectral classifications and membership status of a total sample of 288 candidate members of Chamaeleon~I that have appeared in published studies of the cluster. The latest census of Chamaeleon~I now contains 158 members, 8 of which are later than M6 and thus are likely to be brown dwarfs. I find that many of the objects identified as members of Chamaeleon~I in recent surveys are actually field stars. Meanwhile, 7 of 9 candidates discovered by \\citet{car02} are confirmed as members, one of which is the coolest known member of Chamaeleon~I at a spectral type of M8 ($\\sim0.03$~$M_{\\odot}$). I have estimated extinctions, luminosities, and effective temperatures for the members and used these data to construct an H-R diagram for the cluster. Chamaeleon~I has a median age of $\\sim2$~Myr according to evolutionary models, and hence is similar in age to IC~348 and is slightly older than Taurus ($\\sim1$~Myr). The measurement of an IMF for Chamaeleon~I from this census is not possible because of the disparate methods with which the known members were originally selected, and must await an unbiased, magnitude-limited survey of the cluster. ",
        "introduction": "The targets for observational studies of young stars and brown dwarfs are selected from surveys for members of nearby star-forming regions. The masses for these targets are typically estimated by measuring their spectral types and interpreting their positions on a Hertzsprung-Russell (H-R) diagram with theoretical evolutionary models. Thus, the success of work on young stars and brown dwarfs relies on clear evidence of membership and accurate spectral classifications for putative members of star-forming populations.  At a distance of 160-170~pc \\citep{whi97,wic98,ber99}, the Chamaeleon~I cloud complex is one of the nearest major sites of active star formation. A variety of methods have been used to identify members of Chamaeleon~I, including monitoring of photometric variability at optical \\citep{hof62} and infrared (IR) \\citep{car02} wavelengths, objective prism spectroscopy at H$\\alpha$ \\citep{hen63,men72,hm73,sch77,har93,com99,com00}, X-ray imaging with the {\\it Einstein Observatory} \\citep{fk89} and the {\\it R\\\"ontgen Satellite} \\citep{fei93,alc95,alc97,com00}, and near- to mid-IR photometry from ground-based telescopes \\citep{hjm82,jon85,cam98,ots99,gk01,kg01,per01,car02}, the {\\it Infrared Astronomical Satellite (IRAS)} \\citep{bau84,ass90,whi91,pru91,gs92}, and the {\\it Infrared Space Observatory (ISO)} \\citep{nor96,per99,per00,com00,leh01}. To measure spectral types and to check for signatures of youth and membership in the resulting candidate members, spectroscopy was employed by those authors and in subsequent followup work \\citep{app77,app79,ryd80,app83,wal92,hlf94,law96,cov97,nc99,gp02,gm03,saf03}.  Although Chamaeleon~I has been the target of a large number of studies, many of the objects that have been referred to as members of the cluster lack either conclusive evidence of membership or accurate spectral classifications. In addition, many of the candidates from recent surveys have not been observed with spectroscopy. To address these shortcomings in the current census of Chamaeleon~I, I present optical spectroscopy for most of the objects that have been previously identified as possible members of the cluster, that lack either accurate spectral classifications or evidence of membership, and that are sufficiently bright ($I\\lesssim18$) (\\S~\\ref{sec:obs}). I then measure spectral types from these data and use all available constraints to evaluate the membership status of most of the candidate members of Chamaeleon~I that have appeared in previous work (\\S~\\ref{sec:class}). The implications of this new census for recent searches for members of the cluster are described (\\S~\\ref{sec:imp}). Finally, I estimate extinctions, luminosities, and effective temperatures for the known members and use these data to construct an H-R diagram for Chamaeleon~I (\\S~\\ref{sec:prop}). ",
        "conclusions": "As a part of a new census of the Chamaeleon~I star-forming region, I have performed optical spectroscopy on most of the sources that have been previously identified as possible members of the cluster, that lack either accurate spectral types or clear evidence of membership, and that are optically visible ($I\\lesssim18$). After measuring spectral types for the 179 objects in this sample, I used the spectroscopic data and all other available constraints to evaluate the membership status of 288 potential members that have been presented in published surveys. This analysis has produced a list of 158 confirmed members, 8 of which are later than M6 and thus are likely to be brown dwarfs according to the evolutionary models of \\citet{bar98} and \\citet{cha00}. The membership of 41 of these sources is established for the first time by data in this paper. Many of the objects that have been referred to as members of Chamaeleon~I in previous work lack evidence of membership. For instance, I find that approximately half of the candidates that were classified as young stars through IR photometry and spectroscopy by \\citet{per00} and \\citet{gm03}, respectively, are field stars, predominantly background giants. Meanwhile, most of the candidates discovered by \\citet{car02} are spectroscopically confirmed as members, one of which is the coolest known member of Chamaeleon~I at a spectral type of M8 ($\\sim0.03$~$M_{\\odot}$).  For the known members of Chamaeleon~I that have accurate spectral types, I have estimated extinctions, luminosities, and effective temperatures and used these data to construct an H-R diagram for the cluster. Evolutionary models imply a median age of $\\sim2$~Myr for Chamaeleon~I, which is similar to that of IC~348 and slightly greater than the age of $\\sim1$~Myr for Taurus. In addition to ages, masses of the members of Chamaeleon~I can be estimated from the H-R diagram and evolutionary models. However, the current census of the cluster is not suitable for deriving an IMF because the known members were originally identified by signatures of youth (H$\\alpha$, IR excess, variability) that have different and ill-defined sensitivities as a function of mass. Indeed, I find that it is not possible to define an area of the cluster and an extinction limit within which the current census contains a significant number of members and approaches a high level of completeness for a useful range of masses, which is necessary for constructing a mass function that is unbiased in mass and thus a meaningful representation of the cluster. The measurement of an IMF for Chamaeleon~I from spectroscopic data will require a magnitude-limited survey for cluster members."
    },
    "0402/astro-ph0402353_arXiv.txt": {
        "abstract": "{ We identified new pre-main sequence stars in the region of high-latitude molecular clouds associated with the reflection nebula IC\\,2118, around $l \\sim 208\\degr$ and $b \\sim -27\\degr$. The stars were selected as T~Tauri candidates in objective prism plates obtained with the Schmidt telescope of Konkoly Observatory. Results of spectroscopic follow-up observations, carried out with the FLAIR spectrograph installed on the UK Schmidt and with ALFOSC on Nordic Optical Telescope, are presented in this paper. Based on spectral types, presence of emission lines and lithium absorption line, we identified  five classical T~Tauri stars and a candidate weak-line T~Tauri star projected on the molecular clouds, as well as two candidate pre-main sequence stars outside the nebulous region. Using the near infrared magnitudes obtained from the 2MASS All Sky Catalog (IPAC~2003) we determined  the masses and ages of these stars. We found that the five classical T~Tauri stars projected on the clouds are physically related to them, whereas the other stars are probably background objects. Adopting a distance of 210\\,pc  for IC\\,2118 (Kun et al.~2001) and using Palla \\& Stahler's~(1999) evolutionary tracks we derived an average age of $2.5\\times10^{6}$ yrs and a mass interval of 0.4--1.0\\,M$_{\\sun}$ for the members of the IC\\,2118 association. ",
        "introduction": "\\label{Sect_1}  Small molecular clouds at high galactic latitudes (Magnani, Blitz \\& Mundy~\\cite{MBM})  are usually devoid of star formation. A well-studied exception is MBM~12, containing the young association of low-mass pre-main sequence (PMS) stars MBM\\,12A (e.g. Luhman~\\cite{Luhman}). A less studied example of star forming high latitude molecular cloud is MBM~21 which harbours an infrared source, IRAS~04591$-$0856, associated with a faint nebulosity HHL\\,17 (Gyulbudaghian, Rodr\\'{\\i}guez, \\& Mendoza-Torres~\\cite{GRM}). Persi et al.~(\\cite{Persi}) have shown that HHL\\,17 is a low-mass YSO  between the protostellar and pre-main sequence evolutionary stage. MBM~21 and 22 are projected at an angular distance of some 10 degrees from the Orion~A molecular cloud. They lie at the southernmost part of an extended reflection nebula, IC\\,2118 (Witch Head Nebula), illuminated by $\\beta$~Orionis (Rigel). A $^{12}$CO survey performed with the 4-meter NANTEN radio telescope and covering the whole area of IC\\,2118 (Kun et al.~\\cite{KAY}, hereinafter Paper~I) resulted in the detection of six molecular clouds in the bright region, including MBM~21 (G\\,208.4$-$28.3) and MBM~22 (G\\,208.1$-$27.5). The most massive member of this small group of clouds, G\\,206.4$-$26.0, is not included in the MBM catalogue, but was studied by Bally et al.~(\\cite{Bally91}) and Yonekura et al.~(\\cite{Yonekura}). Low-mass star formation in G\\,206.4$-$26.0 is indicated by the small group of nebulous stars RNO\\,37 (Cohen~\\cite{RNO}), with H$\\alpha$ emission in the two brightest members (Nakano, Wiramihardja \\& Kogure~\\cite{NWK}). The northern of these two stars coincides with the IRAS source 05050$-$0614, having spectral energy distribution indicative of a PMS star (Paper~I). Yonekura et al.'s~(\\cite{Yonekura}) $^{12}$CO, $^{13}$CO and C$^{18}$O studies have shown that both clouds associated with IRAS point sources contain high-density cores, suitable for forming low-mass stars.  IC\\,2118 is a part of the Orion region surveyed for weak-line T~Tauri stars by Alcal\\`a et al.~(\\cite{A96}) on the basis of ROSAT all-sky survey. One wTTS of this sample, RXJ\\,0502.4$-$0744 is projected on the reflection nebula, and another one, RXJ\\,0507.8$-$0931 is located at about 1.5 degrees to the east of it.  The visual appearance of the clouds associated with IC\\,2118 suggests their interaction with Orion OB\\,1 (e.g. Ogura \\& Sugitani~\\cite{OS}), therefore they are usually thought to be as distant as the Orion~A and B molecular clouds, i.e. $\\sim$460\\,pc. The radial velocities of the clouds, however, are more negative ($-5.30  < v_\\mathrm{LSR} < +4.8$\\,km\\,s$^{-1}$) than both those of the main clouds Orion~A and Orion B, (3\\,km\\,s$^{-1} \\lse v_\\mathrm{LSR} \\lse 11$\\,km\\,s$^{-1}$, Bally~(\\cite{Bally89}); Aoyama et al.~(\\cite{AMY})) and the bright stars of the Orion~OB1 association ($v_\\mathrm{LSR} \\sim$ +5\\,km\\,s$^{-1}$, Brown, de Geus \\& de Zeeuw~(\\cite{BGZ})). This velocity pattern  suggests that whereas the Orion~A and B molecular clouds are situated in the receding hemisphere of the expanding interstellar structure around Orion~OB1 (Orion--Eridanus Bubble, Brown, Hartmann \\& Burton~(\\cite{BHB})), the IC\\,2118 clouds belong to its approaching side. These small clouds therefore are probably closer to us than the expansion centre of the Bubble, Ori~OB1a (336$\\pm$16\\,pc, de Zeeuw et al.~\\cite{ZHB}). Considering the cometary shapes of the clouds, Bally et al.~(\\cite{Bally91}) proposed that they are actually located inside the Bubble, whose radius is about 140\\,pc (Brown et al.~\\cite{BGZ}). Based on literature data, Kun et al.~(\\cite{KAY}) adopted 210\\,$\\pm20$\\,pc for the most probable distance of IC\\,2118. This result implies that the clouds are situated inside the Orion--Eridanus Bubble, and close to its surface nearest to us.  The age sequence of OB subgroups of Orion~OB\\,1 (e.g. Brown et al.~\\cite{BGZ}) as well as star formation observed in some cometary globules (e.g. Stanke et al.~\\cite{SSGS}) suggest that interactions of high-mass stars with the interstellar medium have played a significant role in forming the present appearance of the region. Several observed properties of the IC\\,2118 region suggest that low mass star formation has been triggered here by the Orion--Eridanus Bubble. We performed a search for additional PMS stars in order to explore the star forming history of the region. Objective prism Schmidt plates were used to search for H$\\alpha$ emission stars, and then  spectroscopic follow-up observations of the candidates were carried out in order to establish their nature.  In this paper we present the results of our spectroscopic survey. We describe our observations and data analysis in Section~\\ref{Sect_2}. Results on the new PMS stars are shown in Sect.~\\ref{Sect_3}. A brief summary of the paper is given in Sect.~\\ref{Sect_4}. As all of our target stars are included in the 2MASS All Sky Catalog~(IPAC~\\cite{2MASS}), we use the 2MASS source designation for identifying our objects. ",
        "conclusions": "\\label{Sect_3}  Figure~\\ref{Fig3} shows the surface distribution of the newly found PMS stars, together with other known YSOs of the region, overlaid on the IRAS 100$\\mu$m image. In addition to the objects listed in Tables~2 and 3 two wTTS, identified by Alcal\\`a et al.~(\\cite{A96}), and the embedded YSO IRAS~04591$-$0856 are plotted. The five cTTS listed  in Table~3, IRAS~04591$-$0856, as well as the candidate wTTS 2MASS J\\,05060574$-$0646151 are projected on the molecular clouds associated with IC\\,2118, while RXJ\\,0502.4$-$0744 is projected against a lower density part of the cloudy region.  \\subsection{HRD of the target stars} \\label{Sec_3.1}  We used $J$, $H$, and $K_s$ magnitudes obtained from the 2MASS All Sky Catalog~(IPAC~\\cite{2MASS}) to place our stars on the Hertzsprung--Russell diagram. For this purpose their effective temperatures and bolometric luminosities are to be determined. $T_{\\rm eff}$ comes from the spectral type (Kenyon \\& Hartmann~\\cite{KH}), whereas $L_{\\rm bol}$ can be determined from the near-infrared photometric data. Figure~\\ref{Fig4} displays their positions on the $H-K_s$ vs. $J-H$ colour-colour diagram together with the lines indicating the position of zero-age main-sequence, the giant branch, direction of the interstellar reddening and the locus of classical T~Tauri stars determined by Meyer, Calvet \\& Hillenbrand (\\cite{Meyer}). In addition to the stars found during the present survey, the ROSAT wTTS (Alcal\\`a et al.~\\cite{A96} are also plotted.  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{0510fig5.ps}} \\caption[]{Positions of the PMS stars in IC\\,2118 in the $J-H$ vs. $H-K_s$ diagram. Loci of zero-age main sequence, giant branch, and classical T\\,Tauri stars, as well as the slope of interstellar reddening are indicated. 2MASS\\,J05112460$-$0818320, displaying cTTS-like spectrum in the G\\,1200R image, is located on the giant branch.} \\label{Fig4} \\end{figure}  The four cTTS associated with the cloud G\\,206.4$-$26.0, and HHL\\,17 clearly display infrared excess, located to the right of the band of the reddened main sequence, whereas the positions of 2MASS J\\,05020630$-$0850467, associated with the cloud G\\,208.4$-$28.3 and 2MASS J\\,05060574$-$0646151, projected on G\\,206.8$-$26.5, are equally compatible with unreddened cTTS and reddened main sequence stars. 2MASS\\,J05112460$-$0818320, which displayed the outburst during the FLAIR observing run, is also marked as cTTS in Fig.~\\ref{Fig4}, though it lies on the giant sequence, rendering its nature somewhat uncertain.  We made the widely used assumption that the total emission of our target stars in the {\\sl J\\/} band originates from the photosphere ({\\sl e.g.}~Hartigan, Strom \\& Strom~\\cite{Hartigan}). Thus the colour index $J-H$ can be written as $$ J-H = (J-H)_{0} + E_{\\rm CS}(J-H) + E_{\\rm IS}(J-H), $$ where $(J-H)_{0}$ is the true photospheric colour of the star, $E_{\\rm CS}(J-H)$ is the colour excess due to the emission from the circumstellar disk in the {\\sl H\\/} band, and  $E_{\\rm IS}(J-H)$ is the colour excess originating from the difference of interstellar extinctions in the {\\sl J\\/} and {\\sl H\\/} bands.  We dereddened our cTTS onto the locus of unreddened T~Tauri stars in the $H-K_s$ vs. $J-H$ colour-colour diagram (Meyer et al.~\\cite{Meyer}) in order to determine $E_{\\rm IS}(J-H)$. Bolometric luminosities were derived from the {\\sl J} magnitudes and $E_{\\rm IS}(J-H)$ colour excesses by using the interstellar extinction law $A_{\\rm J}=2.65\\,\\times E_{\\rm IS}(J-H)$ (Rieke \\& Lebofsky~\\cite{RL}), and the bolometric corrections tabulated by Hartigan et al.~(\\cite{Hartigan}).  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{0510fig6.ps}} \\caption{Positions of the observed stars in the HRD, assuming a distance of 460\\,pc. Black dots indicate classical T~Tauri stars associated with IC\\,2118, crosses mark the  other target stars and  asterisks are for wTTS detected by ROSAT (Alcal\\`a et al.~\\cite{A96}). Dotted lines indicate the isochrones of 10$^6$, 3$\\times10^6$, 5$\\times10^6$ , 10$^7$, 5$\\times10^7$ and 10$^8$ years, and thin solid lines show the evolutionary tracks from Palla \\& Stahler's~(\\cite{PS}) model. The dashed line corresponds to the birthline and thick solid line indicates the zero age main sequence.} \\vskip -0.4cm \\label{Fig5} \\end{figure}    \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{0510fig7.ps}} \\caption{The T~Tauri stars of IC\\,2118 in the HRD, assuming a distance of 210\\,pc. Isochrones and evolutionary tracks, as well as the birthline and zero age main sequence are indicated as in Fig.~5.} \\vskip -0.4cm \\label{Fig6} \\end{figure}  The distribution  of our target stars in the HRD is displayed in Fig.~\\ref{Fig5}, with the assumption that all of them are located at 460\\,pc, at the distance of main Orion molecular clouds. Uncertainties of $\\log\\,T_\\mathrm{eff}$  are derived from those of the spectral classification. In assessing the uncertainty of $\\log\\,L$ the errors of photometric data, given in the 2MASS All Sky Catalog, and uncertainties of the bolometric corrections due to the error of spectral classification were taken into account. Further sources of uncertainty of $\\log\\,L$ are negligence of the excess luminosity arising from photospheric veiling and circumstellar dust emission. Both of these effects, however, have their minima around 1\\,$\\mu$m (Kenyon \\& Hartmann~\\cite{KH}). Evolutionary tracks and isochrones, as well as the position of the birthline and zero-age main sequence (Palla \\& Stahler~\\cite{PS}) are also shown. However, distances of the stars, in particular of those outside the IC\\,2118 molecular clouds, are actually  unknown. They may  either be low-mass members of Orion~OB1, closer to us than the A and B clouds, or may be situated at different distances as members of the Gould Belt system (Alcal\\`a et al.~\\cite{ACT}). Therefore this figure only indicates that if they are as distant as the giant clouds of most recent star formation, then most of them are PMS stars at different evolutionary stages. The only exception is 2MASS J\\,05094864$-$0906065, located far from the IC\\,2118 clouds on the sky, and below the ZAMS in Fig.~\\ref{Fig5}. This star is probably more distant ($d \\ga 800$pc) than the Orion star forming region, given that H$\\alpha$  and H$\\beta$ emission, seen in its spectrum, are indicative of PMS nature for this spectral type. In this figure the five cTTS found in the IC\\,2118 molecular clouds are located high above the 1\\,Myr isochrone, around the birthline. It was shown by Baraffe et al.~(\\cite{BCAH02}) that stellar ages and evolutionary tracks are very uncertain at this part of the HRD. The birthline shown here is considered as an upper limit for pre-main sequence luminosities, and even it is probable that the youngest accreting low-mass stars appear below this line (Hartmann et al.~\\cite{HCK}). Four of the five cTTS are located at the outer regions of the dense C$^{18}$O cores of their parent clouds (Yonekura et al.~\\cite{Yonekura}), suggesting that they  have already evolved off the birthline. Therefore their positions  in this diagram provide  further support for the result that IC\\,2118 is closer to us than Orion A and B.  Adopting that IC\\,2118 is located at a distance 210\\,pc from us, the HRD shown in Fig.~\\ref{Fig6} has been obtained. In this plot the cTTS projected on IC\\,2118 form a group in the mass interval 0.4--0.9\\,M$_{\\odot}$, and they  are scattered between the isochrones of 1$\\times10^{6}$  and 4$\\times10^{6}$ years. The candidate wTTS seen along the line of sight of a molecular cloud, 2MASS J\\,0506057$-$0646151, is probably a more distant object, lying  behind the clouds. Properties of these five T~Tauri stars, resulted from our study, are shown in Table~\\ref{Tab4}. Effective temperatures $T_\\mathrm{eff}$ corresponding to the spectral classes  are displayed in column 2, and visual extinctions $A_{\\rm V}=9.14\\,\\times E_{\\rm IS}(J-H)$ (Rieke \\& Lebofsky~\\cite{RL}) are shown in col. 3. Luminosities  derived from the 2MASS data are listed in col. 4, and masses and ages  resulting from the Palla \\& Stahler~(\\cite{PS}) model, are shown in columns 5 and 6, respectively. Minimum and maximum values of the derived quantities, resulting from the errors of spectral classification and photometry, are indicated in parentheses.  \\begin{table*} \\caption{Properties of the pre-main sequence stars associated with IC\\,2118, derived from spectroscopic and 2MASS data} \\label{Tab4} \\begin{flushleft} \\begin{tabular}{cccccc} \\hline \\hline \\noalign{\\smallskip} 2MASS J  & $T_\\mathrm{eff}$ & $A_V$ &  $L$ & $M$ & Age  \\\\ &  (K) & (mag) & ($L_{\\sun}$) &  ($M_{\\sun}$) & (10${^6}$ yr) \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 05020630$-$0850467 & 3580\\,($^{3720}_{3470}$) & 0.0\\,($^{+0.5}_{-0.5}$) &  0.33\\,($^{0.42}_{0.27}$)  & 0.38\\,($^{0.45}_{0.30}$) & 1.5\\,($^{2.0}_{1.0}$) \\\\ [.5ex] 05065349$-$0617123 & 4060\\,($^{4205}_{3955}$) & 4.0\\,($^{4.7}_{3.4}$) &  0.84\\,($^{1.04}_{0.66}$)  & 0.80\\,($^{0.90}_{0.66}$) & 2.5\\,($^{3.0}_{2.0}$) \\\\[.5ex] 05071157$-$0615098 & 3580\\,($^{3850}_{3370}$) & 6.8\\,($^{7.8}_{6.3}$) & 0.29\\,($^{0.36}_{0.23}$) & 0.38\\,($^{0.45}_{0.22}$)  & 2.0\\,($^{3.0}_{0.9}$) \\\\ [.5ex] 05073016$-$0610158 & 4205\\,($^{4350}_{4060}$) & 2.5\\,($^{3.1}_{2.1}$) & 0.77\\,($^{0.93}_{0.63}$)  & 0.90\\,($^{1.05}_{0.80}$) &  4.0\\,($^{5.5}_{2.5}$) \\\\ [.5ex] 05073060$-$0610597 & 4060\\,($^{4205}_{3955}$) & 1.7\\,($^{2.4}_{1.2}$) & 1.23\\,($^{1.56}_{1.00}$)  & 0.80\\,($^{0.90}_{0.66}$) & 1.0\\,($^{1.6}_{0.9}$)  \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{flushleft} \\end{table*}    \\subsection{The IC\\,2118 association} \\label{Sect_3.2}  Both the surface distribution of the newly identified cTTS and their position in the HRD suggest the presence of a young association of low-mass stars formed in the high latitude molecular clouds associated with IC\\,2118. The five cTTS identified in this work are projected on two different molecular clouds. The cloud G\\,206.4$-$26.0 hosts four of the stars. The mass of this cloud is 85\\,M$_{\\sun}$ (Paper~I), and its radial velocity of $v_\\mathrm{LSR} \\approx -2.2$\\,kms$^{-1}$, significantly more negative than those of Orion~OB1 and Orion~A and B, suggesting that it represents a distinct subsystem of the Orion star forming region. It contains an elongated dense core mapped in C$^{18}$O by Yonekura et al.~(\\cite{Yonekura}). The mass of the cores, traced by C$^{18}$O, is 25\\,M$_{\\sun}$ (scaled to 210\\,pc Yonekura et al.'s~(\\cite{Yonekura}) result). The stars associated with this cloud are aligned parallel to the long axis of the core, at a mean projected distance of $\\sim\\,0.3$\\,pc from each other, so that the two nebulous objects in RNO\\,37 as well as 05065349$-$0617123 are located at the outskirts, while the 05071157$-$0615098 is projected inside the core (see Fig.~\\ref{Fig3}: the IRAS 100$\\mu$m intensities show largely the same structure  as $^{13}$CO and C$^{18}$O maps). This surface distribution suggests an age sequence: the star closer to the centre of the core should be younger. This age sequence  is washed out by the uncertainties of $T_\\mathrm{eff}$ and $L_\\mathrm{bol}$, but the  signposts of strong accretion, observed in the spectrum of 05071157$-$0615098 may be indicative of its extreme youth indeed. This star may be significantly younger than the age derived from its position in the HRD. Both theoretical and observational studies suggest that strongly accreting PMS stars may be considerably less luminous than their coeval, non-accreting counterparts, mimicking an older age (Hartmann et al.~\\cite{HCK}; Comer\\'on et al. \\cite{CFBNK}).  Comparison of the $J$ magnitudes of the 2MASS and DENIS~(DENIS Consortium~\\cite{DENIS}) data bases, moreover, reveals the variability of this star: contrary to the 2MASS magnitude $J=13.017\\,\\pm0.026$, the same value for DENIS J\\,050711.5$-$061509 is $J=12.693\\,\\pm0.08$\\,mag. The variability in the $J$ band also contributes to the uncertainty of the derived luminosity.  The fifth member of the IC\\,2118 association, 2MASS J\\,05020630$-$0850467 is projected on the molecular cloud G\\,208.3$-$28.4 (MBM\\,21), whose mass was estimated to be 14\\,M$_{\\sun}$  (Paper~I). The radial velocity of this cloud is $v_\\mathrm{LSR} = +4.8$\\,kms$^{-1}$, close to the average value of Orion~OB1. The large velocity difference between the northern and southern clouds of the IC\\,2118 complex may suggest that they are unrelated objects at different distances. This is unlikely because both the illumination of the northern clouds by Rigel and the distance determination for the southern cloud by Penprase~(\\cite{Penprase}) converge to the same distance value adopted here. One may notice, however, that the velocity pattern of the IC\\,2118 complex is similar to that observed by Bally~(\\cite{Bally89}) in Orion~A: while the radial velocities of the southern parts of both Orion~A and IC\\,2118 are nearly the same as that of Orion OB\\,1, the northern portions  have  more positive velocities in Orion~A and  more negative velocities in IC\\,2118. Both regions are located to the south of Ori~OB1a, the centre of the expansion of the Orion--Eridanus Bubble; Orion~A resides in the receding hemisphere, and IC\\,2118 in the approaching one. Thus the observed velocity structures suggest that the northern parts of the clouds, closer to Ori~OB1, have experienced greater acceleration  than those farther from the origin of the shock wave, compressing and subsequently accelerating the clouds.  The cloud contains two dense C$^{18}$O cores having masses of 7.7 and 3.5\\,$M_{\\sun}$, respectively (scaled to 210\\,pc the values derived by Yonekura et al.~\\cite{Yonekura}). The star is located at the edge of the larger, eastern core, whereas the smaller, western core contains the embedded infrared source IRAS~04591$-$0856.  The large-scale geometry and kinematics of the Orion--Eridanus region suggests that star formation in the IC\\,2118 region propagates from the north-east toward the south-west, and also toward us. According to this picture, the two YSOs associated with  G\\,208.3$-$28.4 are probably somewhat younger than their counterparts in G\\,206.4$-$26.0. Both the derived age of  J\\,05020630$-$0850467  and the deeply embedded state of IRAS\\,04591$-$0856 support this hypothesis.  G\\,208.3$-$28.4 is one of the smallest known star forming molecular clouds in our galactic neighbourhood. Its star forming cores contain significantly less material than the average C$^{18}$O mass of $\\approx$12\\,M$_{\\sun}$, required to form a protostar in Taurus and Chamaeleon  (Onishi et al.~\\cite{Onishi}; Mizuno et al.~\\cite{Mizuno}). This may be the consequence of the high ambient pressure from the superbubble, leading to a smaller critical mass for gravitational collapse.  The velocity gradient along the cloud complex suggests that the compressed clouds are subsequently accelerated by the shock propagating from Ori~OB1, and that their acceleration continues after the onset of star formation. The probability of observing YSOs in very small, compressed clouds is low not only because small clouds disperse rapidly, but also because they are swept off the newly formed stars, which keep their velocities while their parental clouds are further accelerated.  Alcal\\`a et al.~(\\cite{ACT2000}) identified a subsample of  wTTS widely distributed over the Orion region with radial velocities $v_\\mathrm{LSR} < +6$\\,kms$^{-1}$, among them are the two stars, RXJ\\,0502.4$-$0744 and  RXJ\\,0507.8$-$0931, located within the field studied here. The presence of the young stars in G\\,208.3$-$28.4 gives some support to the speculation that these stars might have been born in small clouds compressed to form stars and then swept aside and dispersed by the approaching hemisphere of the superbubble. In this case the distance of these stars should be between 200--350\\,pc. The positions in the HRD of both RXJ\\,0502.4$-$0744 and  RXJ\\,0507.8$-$0931 favour the higher limit, because, according to the scenario of sequential star formation, they should be younger than $10^7$ years.  It was established in Paper~I that the clouds associated with IC\\,2118 lie on the surface of the Orion--Eridanus Bubble, being blown with variable powers by the stellar winds and supernova explosions of the massive stars of Orion~OB1 during the last ten million years. The ages of the PMS stars found in the clouds are compatible with the assumption that star formation has been triggered by the superbubble. The complicated geometry and wind history of the OB association (Brown et al.~\\cite{BHB}) hinders both any detailed speculation on the exact position and age of the sources of trigger and any accurate mapping of the shape of the bubble surface. Wherever the shock wave meets a dense medium, a new section of surface will arise.  The stars found during the present studies are probably the most massive members of the young stellar group born in the low-mass, high-latitude molecular clouds. Several faint and red 2MASS and DENIS sources are projected on the clouds, whose nature is uncertain due to the low S/N of the data. Further members of the IC\\,2118 association can be revealed by spectroscopic and deep near infrared observations of these sources."
    },
    "0402/astro-ph0402165_arXiv.txt": {
        "abstract": "{The gravitational lensing of gravitational waves should be treated in the wave optics instead of the geometrical optics when the wave length $\\lambda$ of the gravitational waves is larger than the Schwarzschild radius of the lens mass $M$. The wave optics is based on the diffraction integral which represents the amplification of the wave amplitude by lensing. We study the asymptotic expansion of the diffraction integral in the powers of the wave length $\\lambda$. The first term, arising from the short wavelength limit $\\lambda \\to 0$, corresponds to the geometrical optics limit. The second term, being of the order of $\\lambda/M$, is the leading correction term arising from the diffraction effect. By analysing this correction term, we find that (1) the lensing magnification $\\mu$ is modified to $\\mu ~(1+\\delta)$, where $\\delta$ is of the order of $(\\lambda/M)^2$, and (2) if the lens has cuspy (or singular) density profile at the center $\\rho(r) \\propto r^{-\\alpha}$ ($0 < \\alpha \\leq 2$), the diffracted image is formed at the lens center with the magnification $\\mu \\sim (\\lambda/M)^{3-\\alpha}$. ",
        "introduction": "The gravitational lensing of light is usually treated in the geometrical optics approximation, which is valid in almost all observational situations since the wave length of light is much smaller than typical scales of astrophysical lens objects. However for the gravitational lensing of gravitational waves, the wavelength is long so that the geometrical optics approximation is not valid in some cases. As shown by several authors (Ohanian 1974, Bliokh \\& Minakov 1975, Bontz \\& Haugan 1981, Thorne 1983, Deguchi \\& Watson 1986), if the wavelength $\\lambda$ is larger than the Schwarzschild radius of the lens mass $M$, the diffraction effect becomes important. Thus, the diffraction effect is important for the lens mass smaller than $10^8 M_{\\odot} ~( {\\lambda}/{1 {\\mbox{AU}}} )$, where $1$ AU is the wavelength for the planed laser interferometer space antenna (LISA: Bender \\etal \\cite{b00})   From the above discussion, for $\\lambda \\gsim M$ the diffraction effect is important and the magnification is small (the wavelength is so long that the wave does not feel the existence of the lens), and for $\\lambda \\ll M$ the geometrical optics approximation is valid. In this paper, we consider the case for $\\lambda \\lsim M$, i.e. the quasi-geometrical optics approximation which is the geometrical optics including corrections arising from the effects of the finite wavelength. We can obtain these correction terms by an asymptotic expansion of the diffraction integral in powers of the wavelength $\\lambda$.\\footnote{The asymptotic expansion of the diffraction integral has been studied in optics. See the following text books for detailed discussion: Kline \\& Kay (1965), Ch.XII; Mandel \\& Wolf (1995), Ch.3.3; Born \\& Wolf (1997), App.III, and references therein.}\\ The diffraction integral represents the amplification of the wave amplitude by lensing in the wave optics. It is important to derive the correction terms for the following two reasons: (1) calculations in the wave optics are based on the diffraction integral, but it is time consuming to numerically calculate this integral especially for high frequency (see e.g. Ulmer \\& Goodman \\cite{ug95}). Hence, it is a great saving of computing time to use the analytical expressions. (2) We can understand clearly the difference between the wave optics and the geometrical optics (i.e. the diffraction effect).   This paper is organized as follows: In \\S 2 we briefly discuss the wave optics in gravitational lensing of gravitational waves. In \\S 3 we show that in the short wavelength limit $\\lambda \\to 0$, the wave optics is reduced to the geometrical optics limit. In \\S 4 we expand the diffraction integral in powers of the wavelength $\\lambda$, and derive the leading correction terms arising from the effect of the finite wavelength. In \\S 5 we apply the quasi-geometrical optics approximation to the case of the simple lens models (the point mass lens, SIS lens, isothermal sphere with a finite core lens, and the NFW lens). Section 6 is devoted to summary and discussions. We use units of $c=G=1$. ",
        "conclusions": "We studied the gravitational lensing in the quasi-geometrical optics approximation which is the geometrical optics including the corrections arising from the effect of the finite wavelength. Theses correction terms can be obtained analytically by the asymptotic expansion of the diffraction integral in powers of the wavelength $\\lambda$. The first term, arising from the short wavelength limit $\\lambda \\to 0$, corresponds to the geometrical optics limit. The second term, being of the order of $\\lambda/M$ ($M$ is the Schwarzschild radius of the lens), is the first correction term arising from the diffraction effect. By analyzing this correction term, we obtain the following results: (1)The lensing magnification $\\mu$ is modified to $\\mu ~(1+\\delta)$, where $\\delta$ is of the order of $(\\lambda/M)^2$. (2)If the lens has cuspy (or singular) density profile at the center $\\rho(r) \\propto r^{-\\alpha}$ ($0 < \\alpha \\leq 2$) the diffracted image is formed at the lens center with the magnification $\\mu \\sim (\\lambda/M)^{3-\\alpha}$. Thus if we observe this diffracted image by the various wavelength (e.g. the chirp signal), the slope $\\alpha$ can be determined."
    },
    "0402/astro-ph0402486_arXiv.txt": {
        "abstract": "{ The expected metal enrichment of the intra--cluster medium (ICM) and the partition of metals between cluster galaxies and the hot ICM depends on the stellar Initial Mass Function (IMF). The choice of the IMF in simulations of clusters has also important consequences on the ``cold fraction'', which is a fundamental constraint on cluster physics.  We discuss the chemical enrichment and the cold fraction in clusters as predicted with different IMFs, by means of a straightforward approach that is largely independent of the details of chemical evolution models or simulations. We suggest this simple approach as a guideline to select the input parameters and interpret the results of more complex models and hydrodynamical simulations. ",
        "introduction": "The chemical enrichment of the intra--cluster medium (ICM) has been extensively discussed in literature by means of chemical evolution models of elliptical galaxies with galactic winds, and a variety of scenarios has been advanced (see the review by Matteucci, these proceedings). Very recently, cosmological hydro--dynamical simulations of cluster formation have been developed, that can follow self--consistently the star formation and chemical enrichment history of cluster galaxies and of the ICM (Valdarnini 2003; Tornatore et al.\\ 2004; Romeo et~al.\\ 2004). Star formation has important effects on the hydro--dynamical evolution of the cluster, via energy feedback from supernov\\ae\\ and metal enrichment of the ICM: the first effect contrasts, the second one boosts the cool-out of the hot gas. The chemical evolution of the ICM is not only an interesting issue {\\it per se}, but an important ingredient of the global physical evolution of clusters. As a consequence, the choice of input parameters and ``recipes'' related to star formation (notably, the stellar Initial Mass Function and the implementation of sub--grid feedback effects) is crucial for the results of the simulations.  In this paper we outline a simple procedure to estimate % the chemical enrichment of the ICM, the partition of the % metals between stars and ICM, and the cold fraction expected after an assumed Initial Mass Function (IMF). This provides a guideline to select the optimal input parameters of the simulation; and to distinguish, in the results of a complex and fully self--consistent simulation, what is merely a consequence of the adopted IMF, and what is an effect of the interplay between hydrodynamical evolution and star formation. % ",
        "conclusions": "We remark the following crucial points for modelling the cosmological evolution and chemical enrichment of clusters of galaxies self--consistently. \\begin{itemize} \\item The amount of mass {\\it and} metals locked in the stellar component is not necessarily negligible, depending on the assumed IMF and corresponding $M_*/L$. \\item The observed metal content in stars can be used as a constraint for the efficiency of feed--back and of metal dispersion in the simulations: the wind efficiency cannot be indefinitively enhanced, because the amount of metals in the stars should also be accounted for, and a sizable fraction of the metals produced may thus not be available to enrich the ICM. (Though the problem of present--day simulations is rather the opposite, namely to lock too much metals in the stars; Tornatore et~al.\\ 2004; Romeo et~al.\\ 2004). \\item Once an IMF is chosen for the simulations (preferably a bottom--light IMF with high-mass slope shallower than Scalo, see PMCS), we suggest to compute the corresponding partition of metals between stars and ICM with the simple procedure outlined in this paper. Such expected partition can then be used to test the numerical results. \\item The crucial constraint of the cold fraction should be consistent with the $M_*/L$ relevant to the assumed IMF. As the observational quantity is luminosity, we suggest to compute the luminosity of the simulated cluster consistently with the adopted IMF, and compare it directly to the {\\it observed} $M_{ICM}/L$. Red or IR luminosities are preferred, since they probe the stellar mass with a lower sensitivity to the age of the stellar populations and to recent star formation activity; $M_{ICM}/L_K$ estimates are becoming available (Lin et~al.\\ 2003). \\end{itemize}"
    },
    "0402/astro-ph0402329_arXiv.txt": {
        "abstract": "{We report on the June 2000 long (100 ks) BeppoSAX exposure that has unveiled above 10 keV a new very high energy component of the X--ray spectrum of $\\eta$ Car extending to at least 50 keV. We find that the 2--150 keV spectrum is best reproduced by a thermal $+$ non--thermal model. The thermal component dominates the 2--10 keV spectral range with kT$_h$=5.5$\\pm$0.3 keV and log\\,NH$_h$=22.68$\\pm$0.01. The spectrum displays a prominent iron emission line centred at 6.70 keV. Its equivalent width of 0.94 keV, if produced by the thermal source, gives a slightly sub--solar iron abundance ([Fe/H]=$-$0.15$\\pm$0.02). The high energy tail above 10 keV is best fitted by a power law with a photon index of 2.42$\\pm$0.04. The integrated 13--150 keV luminosity of $\\sim$12 \\LS\\ is comparable to that of the 2--10 keV thermal component (19 \\LS). The present result can be explained, in the $\\eta$ Car binary star scenario, by Comptonisation of low frequency radiation by high energy electrons, probably generated in the colliding wind shock front, or in instabilities in the wind of the S Dor primary star. It is possible that the high energy tail had largely weakened near the minimum of the 5.53 yr cycle. With respect to the thermal component, it probably has a longer recovering time like that of the highest excitation optical emission lines. Both features can be associated with the large absorption measured by BeppoSAX at phase 0.05. ",
        "introduction": "The peculiar southern object $\\eta$ Car is one of the most remarkable variables in our Galaxy due to dramatic changes in its brightness. In 1843 it was the second brightest star in the sky, then suffered a deep fading down to the eighth magnitude by the end of the 19th century (e.g. Viotti 1995). During the last century $\\eta$ Car was slowly and irregularly re--brightening up to the present V$\\simeq$5. Presently, according to the current distance estimates, $\\eta$ Car has a bolometric magnitude around 5$\\times$10$^6$ L$_{\\odot}$ (2$\\times$10$^{40}$ erg s$^{-1}$, e.g. Hillier et al. 2002). A mass loss rate of 10$^{-3/-4}$ M$_{\\odot}$ yr$^{-1}$ or larger has been estimated from observations (e.g. Hillier et al. 2002, van Boekel et al. 2003, Pittard \\& Corcoran 2002, Andriesse et al. 1978).  Optical spectroscopic observations unveiled a peculiar cyclic behaviour, showing regularly repeating excitation minima, with a period of 5.53 years (Damineli et al. 2000). A similar behaviour was also found at other wavelength bands, from radio to X--rays, which is commonly interpreted in terms of a highly eccentric binary model composed of an S Dor--type very luminous primary star, and an unseen early type secondary star. The binary system interacts through colliding winds producing the observed $\\eta$ Car's high temperature, luminous X--ray emission (e.g. Ishibashi et al. 1999, Corcoran et al. 2001).  Recently, thanks to the BeppoSAX unique broad--band X--ray coverage we were able to detect, for the first time, $\\eta$ Car above 10 keV (Viotti et al. 1998; Viotti et al. 2002a, Paper I). We reported the December 1996 observation showing a 13--20 keV flux in excess with respect to the extrapolated 5 keV thermal spectrum that dominates the 2--10 keV range. The presence of a high energy tail was confirmed by the following BeppoSAX observations of 31 December 1999--2 January 2000 (Rebecchi et al. 2001). In particular, in June 2000 a 100 ks exposure unveiled that the tail was probably non--thermal and extending to at least 50 keV (Viotti et al. 2002b). In this work we analyse in detail the latter observation in order to investigate the origin of these very high energy photons, and compare with the previous BeppoSAX observations of $\\eta$ Car. The results are summarised in Table 1.  \\begin{figure*} \\centering \\includegraphics[height=17cm,angle=270]{etasax2_3band.ps} \\caption{The BeppoSAX MECS images of the $\\eta$ Car region on 2000 June 21--23, in three energy bands: (from left to right) 2--5.5 keV (softer), 6--7.5 keV (iron line), and 7.5--10 keV (harder). North is on the top and east to the left; the size of the image field is 54 arcmin. The second brightest source to the west of $\\eta$ Car is the W--R star HD 93162 (WR25).} \\end{figure*} ",
        "conclusions": "We have presented the first in--depth analysis of the spectrum of the $\\eta$ Car system above 10 keV, based on a long exposure BeppoSAX observation. The power law best fit suggests a non--thermal origin of the hard X--ray tail. The integrated 13--150 keV luminosity ($\\sim$12 \\LS), is comparable to the luminosity of the 2--10 keV thermal component (19 \\LS), and suggests the presence of a very effective formation process.  Two models have been considered in the light of the proposed binary nature of $\\eta$ Car. In one model the high energy tail is produced by inverse Compton scattering of the UV stellar photons by relativistic electrons produced in the wind of the primary star, or in the shocked colliding wind region. Alternatively, the high energy photons are produced by self--Comptonisation of the thermal 5 keV emission from relativistic electrons with energies much lower than in the previous case. Though, as suggested by the referee, the weakness of the high energy tail in March 1998 could be better explained as high inverse Compton cooling of the reletivistc electrons during periastron passage, when the colliding winds shock is closer to the stars. Neither model has so far enough support from observations, like detection of non--thermal radio emission from the central source, a point which would deserve new very high resolution radio observations.  A crucial point for understanding the nature of the high energy tail of $\\eta$ Car, would also be to investigate the slope of the spectrum at higher energies and to measure the high energy cutoff of the spectrum, which is related to the energy of the scattering particles. Our PDS observations allow us to determine a lower limit to the cutoff energy of $\\sim$50 keV. INTEGRAL observations might allow measurement of the X--ray spectrum of $\\eta$ Car above 100 keV, and to determine up to what energy the power law spectrum is extending.  Finally, new high energy ($>$10 keV) X--ray observations, e.g. with the foreseen ASTRO--E satellite, of $\\eta$ Car near the periastron passage of the system will provide a clue of where the non--thermal source is located. Our March 1998 upper limit might suggest a recovering time slower than that of the thermal source, as observed in the optical spectra in the high energy emission lines. It would important to investigate the physical link beteen the two phenomena, and whether that behaviour is associated with the high NH$_h$ value found also at the eclipse egress."
    },
    "0402/astro-ph0402603_arXiv.txt": {
        "abstract": "Recent time-resolved X-ray spectra of a neutron star undergoing a superburst revealed an \\fe\\ line and edge consistent with reprocessing from the surrounding accretion disc. Here, we present models of X-ray reflection from a constant density slab illuminated by a blackbody, the spectrum emitted by a neutron star burst. The calculations predict a prominent \\fe\\ line and a rich soft X-ray line spectrum which is superimposed on a strong free-free continuum. The lines slowly vanish as the ionization parameter of the slab is increased, but the free-free continuum remains dominant at energies less than 1~\\kev. The reflection spectrum has a quasi-blackbody shape only at energies greater than 3~\\kev.  If the incident blackbody is added to the reflection spectrum, the \\fe\\ equivalent width varies between 100 and 300~eV depending on the ionization parameter and the temperature, $kT$, of the blackbody. The equivalent width is correlated with $kT$, and therefore we predict a strong \\fe\\ line when an X-ray burst is at its brightest (if iron is not too ionized or the reflection amplitude too small). Extending the study of reflection features in the spectra of superbursts to lower energies would provide further constraints on the accretion flow. If the \\fe\\ line or other features are relativistically broadened then they can determine the system inclination angle (which leads to the neutron star mass), and, if the mass is known, a lower-limit to the mass/radius ratio of the star. ",
        "introduction": "\\label{sect:intro} The reprocessing of external X-rays by accreting material has been observed from active galactic nuclei (AGN) and Galactic black hole candidates (GBHCs) for over a decade \\citep*[e.g.,][]{pou90,np94,gie99,bvf03,mil03}. In these objects, the origin and location of the illuminating X-ray power-law is largely unknown, although it is widely believed to be either within a magnetic accretion disc corona \\citep*[e.g.,][]{gal79,haa93,haa94,dm98} or a centrally located geometrically thick flow \\citep*[e.g.,][]{sle76,zdz99,zlgr03}. Nevertheless, many models of X-ray reflection from black hole accretion discs have been computed \\citep*{gf91,mpp91,ros93,zyc94,ros99,nkk00,nk01,brf01,roz02,btb04}, and, in some cases, successfully applied to data \\citep*{bif01,orr01,der02,lon03,bdn03,mil03}.  Recently, \\citet{sb02} discovered an \\fe\\ line and edge in the X-ray spectra of the low-mass X-ray binary (LMXB) \\fouru\\ as it was undergoing a superburst --- a powerful thermonuclear explosion on the surface of the neutron star that can last many hours \\citep*{corn00,sb03,kuul03}. Based on the strength of the \\fe\\ line, these authors suggested that the features were caused by reflection of the burst blackbody spectrum from the surrounding accretion disc \\citep[e.g.,][]{dd91}. This hypothesis was supported by spectral fitting with detailed reflection models \\citep{bs04}. The unique aspect of reflection during a burst from a LMXB is that there is no uncertainty in the location of the X-ray source, as it is located on the surface of the neutron star and outshines the persistent X-ray emission from the disc. This fact allows for far less ambiguity when analyzing changes to the reflection features, as they are more likely to trace the evolution in the accretion flow \\citep{bs04}.  X-ray emission lines have also been observed in the persistent emission from LMXBs \\citep[e.g.,][]{sma93,ang95,sch99,adnm00}, and these are also thought to be caused by reprocessing by matter in the accretion flow \\citep{ss73,mmh82}. The spectral shape of the persistent emission is consistent with bremsstrahlung with a temperature of 5--10~\\kev\\ \\citep*{lph01}, most likely arising from the boundary layer between the disc and the neutron star. Detailed models of the photoionized layer on the disc have been performed and predict a wealth of emission lines in the soft X-ray band \\citep*[e.g.,][]{kk94,jrl02} that can be compared against observations made with the grating spectrometers onboard \\textit{Chandra} and \\textit{XMM-Newton} \\citep[e.g.,][]{sch01,cott01,kabc03}. In contrast, a X-ray burst emits a blackbody spectrum \\citep*{swa77,hld77}, and can reach the Eddington luminosity ($L_{\\mathrm{Edd}} \\sim 10^{38}$~erg~s$^{-1}$), well over an order of magnitude greater than the persistent luminosity of many LMXBs \\citep{sb03}. Thus, disc reflection will be dominated by the blackbody component during the burst.  This paper presents reflection spectra from a uniform accretion disc illuminated by a blackbody, as in a LMXB during a Type~I X-ray (super)burst. These spectra were used to fit the superburst data of \\fouru\\ \\citep{bs04}, but, as shown below, they contain much more information than what was actually needed for those \\textit{Rossi X-ray Timing Explorer} (\\textit{RXTE}) data. Observations of reflection features from Type~I bursts or superbursts with sensitive, large bandwidth instruments such as the X-ray Telescope (XRT) on \\textit{Swift} could provide a wealth of new information on the structure and behaviour of accretion discs.  In the next section we outline the reflection calculations, and present the resulting spectra in Section~\\ref{sect:res}. We conclude by discussing the results in Section~\\ref{sect:discuss}. ",
        "conclusions": "\\label{sect:discuss} Searches for spectral features during X-ray bursts from neutron stars have been ongoing for many years. Since the explosion is occurring on the surface of the star, the spectral features would be expected to be redshifted and would therefore be a measure of the mass-to-radius ratio of the star (via $1 + z = (1-GM/c^2 R)^{-1/2}$), a number needed to constrain the many possible equations of state of nuclear matter. During the 1980s there were a few reports of X-ray absorption lines at 4.1~\\kev\\ during Type~I X-ray bursts \\citep*{wak84,nit88,mag89}, which were interpreted as redshifted Ly$\\alpha$ absorption from He-like iron. However, the equivalent width of these absorption lines were on the order of hundreds of eV, which theoretical models of spectral formation in burst atmospheres could not reproduce \\citep{frf87,dfr92}. Observations by more sensitive instruments on \\textit{RXTE} and \\textit{BeppoSAX} have been unable to discover any other examples of a 4.1~\\kev\\ absorption line. Recently, the Reflection Grating Spectrometer (RGS) on \\textit{XMM-Newton} has uncovered evidence for redshifted absorption lines during the X-ray bursts of EXO~0748-676 \\citep{cpm02}, although data from 28 bursts were needed in order to obtain the minimum signal-to-noise necessary to find the lines.  As first suggested by \\citet{dd91}, X-rays from an explosion on a neutron star can be reprocessed by the surrounding accretion disc, leading to features such as an \\fe\\ line and edge. The analysis of these features can lead to many new insights about the structure and evolution of the accretion disc \\citep{bs04}, as well as basic information about the neutron star itself. For example, if the emission lines are relativistically broadened then modeling can reveal the inclination angle of the system to the line of sight \\citep{fab89}, a parameter often needed to obtain masses of the neutron star and its binary companion. Modeling of relativistic lines also produce a radius where the reflection originates (assuming some emissivity profile). This radius is measured in gravitational units ($GM/c^2$), so if the mass of the neutron star is known, the radius can be converted to physical units, allowing a determination of a mass-to-radius ratio. Since the reflection may occur very close to the surface of the neutron star ($10$--$20$~$GM/c^2$ in 4U~1820-30; \\citealt{bs04}), the measurement may provide a useful lower-limit to the $M/R$ ratio for the neutron star.  \\textit{RXTE} has so far been the most successful telescope to obtain time-resolved spectroscopy of X-ray bursts, especially of the more energetic superbursts \\citep[e.g.,][]{sb02}. The telescope produces data for energies greater than 3~\\kev\\ which allows a good determination of the blackbody spectrum and, if detected, the \\fe\\ line. Yet, the reflection models predict a wealth of information at energies less than 3~\\kev\\ (Fig.~\\ref{fig:spectra}) that can provide further constraints on the abundances, ionization state and dynamics of the accretion disc. These features will also strongly depend on the density of the disc because of the importance of collisional effects and free-free emission/absorption. The spectra shown in Fig.~\\ref{fig:spectra} were calculated assuming a density of $n_{\\mathrm{H}}=10^{15}$~cm$^{-3}$, appropriate for the outer regions of the disc. The density at a distance of a few gravitational radii from the neutron star is expected to be $> 10^{21}$~cm$^{-3}$, even for a radiation-pressure dominated disc \\citep{ss73}, which is beyond the range of validity for the reflection code. To illustrate the effects of higher density on the soft X-ray spectrum, Figure~\\ref{fig:densitycompare} compares three reflection spectra computed with $n_{\\mathrm{H}}=10^{15}$~cm$^{-3}$ with ones computed at $10^{18}$~cm$^{-3}$, where the code remains valid. \\begin{figure} \\centerline{ \\includegraphics[width=0.5\\textwidth]{density_compare.eps} } \\caption{The solid lines show three reflection spectra computed with a density of $n_{\\mathrm{H}}=10^{18}$~cm$^{-3}$, while the dotted lines denote models with identical $\\xi$ but with $n_{\\mathrm{H}}=10^{15}$~cm$^{-3}$. A $kT=2.5$~\\kev\\ blackbody was the illuminating continuum for all cases. The increase in density results in only minor differences in the spectra at energies $\\ga 3$~\\kev, but there are significant changes at lower energies.} \\label{fig:densitycompare} \\end{figure} The figure shows very little difference in the spectra in the \\textit{RXTE} band including the \\fe\\ line. This implies that the \\fe\\ EWs computed from the $n_{\\mathrm{H}}=10^{15}$~cm$^{-3}$ models (Fig.~\\ref{fig:ews}) should still be useful guidelines for future observations. On the other hand, the three order of magnitude increase in density significantly alters the soft X-ray features predicted by the reflection spectra. The differences mainly lie at energies $< 0.5$~\\kev, where the continuum level is raised due to the increase in free-free emission and absorption. As a result, the EWs of many of the emission lines have been substantially decreased. The expectation is that increasing the density further will continue to raise the continuum level at soft energies, as more line emission is absorbed by the continuum, and decrease the EW of the remaining emission lines. Clearly, the soft X-ray emission features are a strong diagnostic for the accretion disc density.  It would be therefore be very interesting to obtain rapid, time-resolved broad-band spectra of X-ray bursts to exploit the information in those spectral features. In the near future only the XRT onboard the \\textit{Swift} satellite will have these features, plus the necessary rapid response capability. Unfortunately, soft X-ray spectroscopy of burst sources will be challenging. LMXBs are old systems and many of them lie toward to the Galactic center, thus suffering from significant absorption due to neutral hydrogen along the line of sight. However, there are a number of LMXBs which reside in globular clusters that may provide lower extinction columns and allow sensitive soft X-ray observations. Rapid follow-up of X-ray bursts from LMXBs globular clusters therefore provide the best chance of utilizing the full information obtained in the reflection spectra."
    },
    "0402/astro-ph0402435_arXiv.txt": {
        "abstract": "We have obtained spectra for 1273 stars using the 0.9m Coud\\'e Feed telescope at Kitt Peak National Observatory.  This telescope feeds the coud\\'e spectrograph of the 2.1m telescope.  The spectra have been obtained with the \\#5 camera of the coud\\'e spectrograph and a Loral 3K X 1K CCD. Two gratings have been used to provide spectral coverage from 3460~\\AA \\ to 9464~\\AA, at a resolution of $\\sim$1\\AA \\ FWHM and at an original dispersion of 0.44~\\AA/pixel.  For 885 stars we have complete spectra over the entire 3460~\\AA \\ to 9464~\\AA \\ wavelength region (neglecting small gaps of $<$ 50 \\AA), and partial spectral coverage for the remaining stars.  The 1273 stars have been selected to provide broad coverage of the atmospheric parameters T$_{eff}$, log g, and [Fe/H], as well as spectral type.  The goal of the project is to provide a comprehensive library of stellar spectra for use in the automated classification of stellar and galaxy spectra and in galaxy population synthesis.  In this paper we discuss the characteristics of the spectral library, viz., details of the observations, data reduction procedures, and selection of stars.  We also present a few illustrations of the quality and information available in the spectra.  The first version of the complete spectral library is now publicly available from the National Optical Astronomy Observatory (NOAO) via FTP and HTTP. ",
        "introduction": "The need for a comprehensive database of digital stellar spectra covering a large range in T$_{eff}$, log g, and [Fe/H] at moderate (1-2 \\AA \\ FWHM) spectral resolution has increased considerably in the past decade.  The uses of such a spectral library are wide ranging.  To begin with, synthetic stellar spectra generated from model atmospheres now incorporate such extensive line lists (e.g., Kurucz 1993, 1994; Bell \\& Gustaffson 1978, 1989; Tripicco \\& Bell 1992, 1995) that the synthetic spectra can be compared to empirical spectra at increasingly high spectral resolution. Once the basic accuracy of the synthetic spectra has been established, the synthetic spectrum technique can then be utilized to explore areas of atmospheric parameter space that are not adequately represented in Solar Neighborhood stars.  For example, recent results from galaxy population synthesis studies (e.g., Kuntschner \\etal 2002; Terlevich \\& Forbes 2002; Trager \\etal 2000 and references therein) indicate that the integrated spectra of most early-type galaxies are dominated by metal-rich stars with non-solar abundance ratios that are not found in the Solar Neighborhood. Second, with the advent of large databases of both stellar and galaxy spectra, generated by multi-fiber and multi-aperture spectrographs, methods for the automated parametrization of spectra in a fast and reliable manner are becoming essential.  Indeed, various researchers have been exploring the application of Artificial Neural Networks to the automated parametrization of stellar spectra (e.g., Singh, Bailer-Jones, \\& Gupta 2002).  In addition, Katz \\etal (1998) have developed the TGMET software, which establishes the best match between a target spectrum and a library of reference spectra.  For these techniques to be applied, a comprehensive set of reference spectra with known atmospheric parameters is essential.  Finally, in the field of spectral synthesis of the integrated light of galaxies, the need to resolve individual features in galaxy spectra, to effectively resolve the problems of non-uniqueness in extracting mean age and metallicity information from integrated spectra, has become increasingly evident (Worthey 1994; Vazdekis \\& Arimoto 1999).  To carry out the modeling of the composite spectra of galaxies of different ages and metallicities requires a spectral database that covers all areas in the HR diagram sampled by theoretical isochrones.  Despite the clear necessity for a comprehensive spectral library, there is surprisingly little existing material to select from.  In fact, until recently the existing spectral libraries have primarily consisted of low resolution (5 - 20 \\AA \\ FWHM) spectrophotometry of typically $\\sim$100 stars, covering a range in spectral type, but largely restricted to solar chemical composition (Pickles 1998 and references therein).  This traditional emphasis on accurate spectrophotometry with broad wavelength coverage at low resolution has likely reflected both the low number of pixels in spectroscopic detectors until recently, coupled with the low resolution approach to spectral synthesis studies of stars and galaxies that supplied the chief driver for the libraries.  However, the advent of large format CCD detectors in spectroscopy, along with the above-mentioned trends toward higher resolution modeling of the spectra of stars and of galaxies in integrated light, is now leading to the production of several extensive libraries of stellar spectra at higher spectral resolution. In particular, Jones (1999, see also Leitherer \\etal 1996) completed a spectral database of 684 stars at 2 \\AA \\ resolution (FWHM) using the Coud\\'e Feed telescope at the Kitt Peak National Observatory (KPNO).  Due to the small format of the existing CCD detector at that time, coverage is restricted to two $\\sim$700 \\AA \\ spectral regions.  More recently, Prugniel \\& Soubiran (2001) published a spectral library, based on the ELODIE echelle spectrograph at the Observatoire de Haute-Provence, that covers the wavelength interval $\\lambda$$\\lambda$4100-6800 \\AA \\ at a resolution of R=42000 for 708 stars that cover a large range in atmospheric parameters.  Finally, Cenarro \\etal (2001) have observed 706 stars in the wavelength region $\\lambda$$\\lambda$8350-9020~\\AA \\ at 1.5 \\AA \\ resolution, to characterize the behavior of the near-IR Ca II triplet for galaxy population synthesis.  In this paper we describe a new database of stellar spectra, encompassing more than 1000 stars, and covering the spectral region $\\lambda$$\\lambda$3400--9500 \\AA \\ at a resolution of $\\sim$1 \\AA \\ FWHM.  The primary emphasis of this new spectral library is on broad wavelength coverage, particularly well into the blue, at a resolution sufficient to resolve numerous diagnostic spectral features that can be used in the automated parametrization of spectra and in the population synthesis of galaxies.  We have also emphasized coverage in atmospheric parameter space, particularly at lower metallicity. The purpose of this paper is to provide a guide to the new spectral library, which is now publicly available from the National Optical Astronomy Observatory (NOAO) via ftp and http.  The paper presents technical information about the observations and the sample of stars observed to make the database readily useful to the astronomical community.  In \\S2 we describe the observational procedures used in creating the library, and in \\S3 we describe how the sample of stars was selected.  In \\S4 the data reduction procedures are described.  In \\S5 we describe the actual format of the archived spectral database as well as the tables which summarize the key information about the stars.  Finally, in \\S6 we show examples of the stellar spectra, some of which are meant to illustrate the challenges facing the automated parametrization of stellar spectra. ",
        "conclusions": ""
    },
    "0402/astro-ph0402059_arXiv.txt": {
        "abstract": "{The accretion disc eclipse mapping method is an astrotomographic inversion technique that makes use of the information contained in eclipse light curves to probe the structure, the spectrum and the time evolution of accretion discs in cataclysmic variables. This paper presents examples of eclipse mapping results that have been key to improve our understanding of accretion physics. ",
        "introduction": "Cataclysmic Variables (CVs) are close interacting binaries in which mass is fed to a white dwarf (the primary) by a Roche lobe filling companion star (the secondary) via an accretion disc, which usually dominates the ultraviolet and optical light of the system (Warner 1995). Accretion discs in CVs cover a range of accretion rates, \\.{M}, from the low-mass transfer dwarf novae (the discs of which show recurrent outbursts on timescales of weeks to months) to the high-mass transfer nova-like variables (the discs of which seems to be stuck more or less in a steady state).  One of the difficulties in studying accretion disc physics comes from the fact that the physical conditions in an accretion disc are expected to vary by large amounts with disc radius (the temperature distribution in a steady-state disc decreases as $T\\propto R^{-3/4}$), making the integrated disc spectrum a complex combination of light emitted from regions of distinct physical conditions (e.g., Frank, King \\& Raine 1992).  In this regard, CVs are excellent sites for studying accretion physics because their binary nature and relatively short orbital periods enable the application of powerful indirect imaging techniques. These techniques overcome the intrinsic ambiguities associated with the composite disc spectra by providing spatially resolved information about accretion discs on angular scales of micro arcseconds -- well beyond the current direct imaging capabilities.  The Eclipse Mapping Method (Horne 1985) is one of these techniques. It assembles the information contained in the shape of the eclipse into a map of the accretion disc surface brightness distribution. The {\\em eclipse map} is defined as a grid of intensities centred on the white dwarf and usually contained in the orbital plane. The eclipse geometry is specified by the inclination $i$, the binary mass ratio $q$ (=$M_2/M_1$, where $M_2$ and $M_1$ are the masses of, respectively, the secondary star and the white dwarf) and the phase of inferior conjunction. Given the geometry, a model eclipse light curve can be calculated for any assumed brightness distribution in the eclipse map. A computer code then iteratively adjusts the intensities in the map (treated as independent parameters) to find the brightness distribution the model light curve of which fits the data eclipse light curve within the uncertainties. The quality of the fit is checked with a consistency statistics, usually $\\chi^2$. Because the one-dimensional data light curve cannot fully constrain a two-dimensional map, additional freedom remains to optimize some map property. A maximum entropy procedure (e.g., Skilling \\& Bryan 1984) is used to select, among all possible solutions, the one that maximizes the entropy of the eclipse map with respect to a smooth default map. Details of the mathematical formulation of the problem can be found in Horne (1985) and Baptista (2001). A movie ilustrating the iterations of an eclipse mapping experiment from start to convergence is available at \\underbar{\\bf www.astro.ufsc.br/$\\sim$bap/slide1.gif}. ",
        "conclusions": "Eclipse mapping is a powerful tool to probe the radial and vertical disc structures, as well as to derive the physical conditions in accretion discs. Partly thanks to the many experiments performed over the last two decades, our picture of accretion discs was enriched with an impressive set of new details such as gas outflow in disc winds, gas stream overflow, flared discs with azimuthal structure at their edge, chromospheric disc line emission, ellipsoidal precessing discs, sub-Keplerian spiral shocks, and moving heating/cooling waves during disc outbursts.  Additional interesting eclipse mapping results are expected in the near future. For example, fitting state-of-the-art disc atmosphere models to the spatially-resolved spectra is an obvious next step to the spectral mapping experiments and will allow estimates of fundamental physical parameters of accretion discs, such as gas temperature, surface density, vertical temperature gradient, Mach number and viscosity, setting important additional constrains on current disc models."
    },
    "0402/astro-ph0402573_arXiv.txt": {
        "abstract": "s{We review cm and mm observations of thermal molecular line emission from high redshift QSOs. These observations reveal the massive gas reservoirs (10$^{10}$ to 10$^{11}$ M$_\\odot$) required to fuel star formation at high rates. We discuss evidence for active star formation in QSO host galaxies, and we show that these high redshift, FIR-luminous QSOs follow the non-linear trend of increasing L$_{FIR}$/L$'$(CO) with increasing L$_{FIR}$. We conclude with a brief discussion of the recent CO detection of the most distant QSO at $z=6.42$, and its implications for cosmic reionization.} ",
        "introduction": "Over the last few years, the study of high redshift QSOs has been revolutionized in three ways. First, wide field surveys have revealed 100's of high z QSOs, right back to the epoch of cosmic reionization ($z > 6$; e.g., Fan et al. 2003). Second, it has been shown that most (all?) low redshift spheroidal galaxies have central super-massive black holes (SMBH), and that the black hole mass correlates with bulge velocity dispersion. This M$_{BH}$-$\\sigma_{\\rm v}$ correlation suggests coeval formation of galaxies and SMBH, thereby making SMBHs a fundamental aspect of the galaxy formation process (Gebhardt et al. 2000).  And third, mm surveys of high redshift QSOs find that 30$\\%$ of the sources are `hyper-luminous infrared galaxies' ($L_{FIR} = 10^{13}$ L$_\\odot$), corresponding to thermal emission from warm dust, and that this fraction is {\\sl independent of redshift out to $z = 6.4$} (Omont et al.  this vol.). If the dust is heated by star formation, the implied star formation rates are extreme ($> 10^3$ M$_\\odot$ year$^{-1}$), consistent with the formation of a large elliptical galaxy on a dynamical timescale of 10$^8$ years.  On the other hand, the FIR luminosity constitutes typically only 10$\\%$ of the bolometric luminosity of the sources, such that dust heating by the AGN remains an alternative. Demographic studies show that SMBHs acquire most of their mass during major accretion events marked by the QSO phenomenon (Yu \\& Tremain 2002).  Molecular line observations (typically CO) of FIR-luminous high $z$ QSOs have revealed large gas masses in most cases observed to date (see Table 1).  Such gas reservoirs are a prerequisite for star formation models for dust heating in FIR-luminous high $z$ QSOs. The typical gas depletion timescales are of order 10$^7$ to 10$^8$ years, if the dust is heated by star formation.  In this review we consider this question in more detail.  We restrict ourselves to $z > 2$ QSOs (see also Barvainis 1999; and see Scoville et al. 2004 and Sanders \\& Mirabel 1996 for observations of lower redshift QSOs). We assume a standard concordance cosmology. ",
        "conclusions": ""
    },
    "0402/hep-ph0402126.txt": {
        "abstract": "We present a general slow-roll formalism within braneworld-motivated cosmologies with non-standard effective Friedmann equations. Full towers of parameters involving either the inflaton potential or the Hubble parameter are constructed and the dynamics of non-tachyonic and tachyonic fields are considered in detail; exact cosmological solutions and the inflationary attractor condition are provided. We compare scalar-driven and tachyon-driven accelerating eras through slow-roll correspondence and the observational imprint on early-universe structures. ",
        "introduction": "Motivated by recent developments of string, superstring and M theories, several models for a multidimensional target spacetime have been proposed \\cite{HW1,HW2,ADD98,AADD,LOSW,RSa,RSb}. Soon they found important applications in cosmology by revitalizing the idea that the visible universe is a (3+1)-dimensional variety (a 3-brane) embedded in a bulk with some extra dimensions, either non-compact or compactified \\cite{aka82,RuS,vis85}. Typically, the background metric on the brane is the Friedmann-Robertson-Walker (FRW) metric and the Einstein equations are modified in accordance with the gravity model permeating the whole spacetime. This in turn produces the basic FRW equations for the cosmological evolution; for a comprehensive review, see \\cite{rub01,mar03}. In particular, the Randall-Sundrum scenarios \\cite{RSa,RSb} and their Gauss-Bonnet generalization \\cite{KKL1,KKL2} (see also \\cite{LN,GrP} and references therein) seem promising because of their high-energy features, while preserving the standard cosmology in the low-energy limit.  In order to reconcile cosmological and astrophysical observations with the standard big bang scenario, it is sufficient (but not necessary) to invoke the inflationary paradigm; according to it, the early Universe experienced a phase of accelerated expansion driven by an effective cosmological constant. This mechanism is triggered by the dynamics of a scalar field rolling down its potential and may also provide an explanation for the present phase of acceleration; in the most famous version of inflation, the rolling is slow enough to justify the adoption of the slow-roll (SR) formalism. For a review of primordial 4D inflation and the SR approach, see \\cite{lid97}. Recently, due to many progresses made in understanding the vacuum structure of string theory (in particular, see \\cite{sen1,sen2,sen3,sen4,HKM,GS,KMM1,KMM2,sen5,sen6,sen7,sen8}), the eventuality that the scalar field is tachyonic has been explored.  In this paper we develop a general SR formalism starting from a FRW equation which is polynomial in the brane energy density, $H^2 \\propto \\rho^q$; by this way we obtain a Hamilton-Jacobi and SR formulation of the cosmological evolution which is valid for many known braneworld theories either in a particular energy limit or time interval. We will restrict ourselves to a brane filled (1) with a scalar or (2) with a tachyon field with an effective Dirac-Born-Infeld action provided by the low coupling limit of non-perturbative string theory. The advantages of this approach are several. It provides a concise and versatile formalism to explore different cosmological models and determine their main features, such as exact classes of solutions, the inflationary attractor and the inflationary imprint on the structure formation of the early Universe.  A general example of such a braneworld-generated cosmological equation is given in \\cite{ChF}. This model has been used to describe the post-inflationary evolution and in this case has been dubbed ``Cardassian cosmology'' \\cite{car1,car2,car3,car4,car5,car6,car7,car8,car9,car10,car11,car12}. Here, we will take a rather different perspective and ask how a period of non-standard expansion can modify the usual early-Universe picture.  The plan of this work is the following: in Sec. \\ref{background} we motivate the ``patch cosmology'' approach and set the background equations, both for a scalar and a tachyon field. Section \\ref{slowroll} is devoted to the slow-roll formalism and its consequences in the scalar and tachyon scenarios. The inflationary attractor is then considered in Sec. \\ref{attractor}, while in Sec. \\ref{exact} we construct some exact cosmological solutions; these sections are presented in a self-contained manner. In Sec. \\ref{pert} we calculate next-to-lowest order scalar spectral amplitudes and indices and compare 4D and Randall-Sundrum models through consistency equations. The intertwine between scalar-driven and tachyon-driven cosmologies is stressed in Sec. \\ref{vs} and in the Appendix we use the slow-roll correspondence between tachyonic and non-tachyonic dynamics to estimate the non-Gaussianity produced during a four-dimensional tachyon-driven inflationary period. Conclusions are in Sec. \\ref{concl}.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\label{concl}  In this paper we have explored the behavior of scalar field and tachyon field dynamics by reformulating the Hamilton-Jacobi equations in terms of towers of slow-roll parameters; exact cosmological solutions and the behavior of the inflationary attractor have been provided, as well as the spectra of cosmological perturbations. Scalar-driven and tachyon-driven scenarios are compared and a first-SR-order correspondence is found. In the Appendix some comments on the inflationary bispectrum are made.  It has been assumed that a non-standard Friedmann equation is originated in a non-trivial gravitational context, in which the presence of extra dimensions and the confinement of the visible Universe on a 3-brane determines the effective geometry a brane observer experiences. Indeed this is what explicit gravity models may produce, for example the Randall-Sundrum and the Gauss-Bonnet scenarios. In particular energy regimes, the effective Hubble parameter scales as a power of the energy density on the brane, $\\widetilde{H}^2 = \\rho^q$; we have restricted our analysis to such a ``patch'' cosmology with constant $q$. This approach does not account for the full structure of the effective Friedmann equation, that must recover the four-dimensional behavior at a sufficiently low density-to-brane tension ratio; however, this is sufficient to formulate most of the non-standard predictions concerning the brane content that arise when the 4D Friedmann equation is violated, $q \\neq 1$. On the contrary, quantities such as the gravitational spectrum must be computed with a complete gravity theory at hand.  With no reference to the gravitational sector, two important assumptions, intimately connected with the evolution of the matter content, emerge; namely, to consider an empty bulk and neglect the Weyl tensor contribution. In particular, there is no source term in the continuity equation (\\ref{conti}). In the Randall-Sundrum model, several works have shown that bulk physics mainly affects the small-scale or late-time cosmological structures, i.e., that part of the spectrum which is dominated by post-inflationary physics \\cite{GRS,GoM,LCML,IYKOM,koy03}. However, it is possible that a non-zero brane-bulk flux would modify the inflationary spectra. For instance, production of particles when the inflaton does not lie in its vacuum state may generate a non-Gaussianity signature during the accelerated expansion \\cite{MRS,GMS}. Cosmic microwave background (CMB) observations strongly constrain the maximum number density of these particles and the $n$-point correlation functions of the resulting perturbations; with a brane-bulk exchange mechanism and interactions at the KK energy scale, this number density, as well as the predicted non-Gaussianity, may vary non-trivially. Thus, the adoption of a modified continuity equation may lead to a richer scenario. See, e.g., \\cite{VDMP,LSR,KKTTZ,LMS} for Randall-Sundrum cosmologies with non-diagonal bulk stress-energy tensor and \\cite{LT} for a six-dimensional example.  We have focused our attention on patch cosmologies with positive $q$ index, since braneworld models, including the above-mentioned scenarios, generally lay in this range of values. This choice may be justified by other considerations which in turn lead to interesting possibilities.  First, according to Eq. (\\ref{dotH0}), it is not possible to pass from an energy regime A to another B with $\\text{sgn}(q_B) =-\\text{sgn}(q_A)$ without spoiling the monotonicity of the Hubble parameter. If the dominant energy condition holds, this eventuality might be discarder if one believes in dS-CFT correspondence \\cite{str1,SIT,str2,NO,BBM,SSV,kle02,haly1,van04,LVL,haly2,DLS,LuP}, relating time evolution in a de Sitter cosmology to the renormalization group flow of a dual boundary field theory. From a cosmological point of view, the infrared fixed point corresponds to an early inflationary period driven by an effective large cosmological constant $\\Lambda_{IR}$ and a small number of effective degrees of freedom, $n_\\text{\\tiny DOF} \\sim 1/\\Lambda_{IR}$, while the ultraviolet fixed point is a late-time de Sitter universe with a small cosmological constant $\\Lambda_{UV}$ and high number of degrees of freedom. In $dS_4$, the flow is governed by a  monotonically varying $c$-function or central charge $c \\propto H^{-2}$. Then, since standard cosmology has $q=1$ and $H>0$, the irreversibility of the flow imposes that $q$ should be always positive. However, apart from intrinsic theoretical problems for this conjectured equivalence, transitions between two such (inflationary) regimes are beyond the scope of the simplified energy-patch cosmology we adopted and, indeed, of the RG approach explored so far (see references for details).\\footnote{More generally \\cite{bou02}, the holographic principle experiences several difficulties in a cosmological context, in particular with inflation \\cite{KaL}.} The author does not know how to reasonably apply such a duality to the present case of an effective 4D quasi--de Sitter brane universe embedded in a higher-dimensional bulk;  therefore, a comparison with a holographic picture appears still unclear and will not be done here.  Secondly, consider the case of a scalar field, Eq. (\\ref{rhop}); as discussed in the previous section, an ESR expansion of the energy density yields $\\rho^q \\propto q \\dot{\\phi}^2_{eff}/2+V_{eff}$, where the effective theory includes the dimensional contribution of $\\beta_q$. If $q<0$, the kinetic term has the wrong sign and this may lead to unitarity problems when quantizing the brane field (also: particles with negative energy propagate forward in time); for a scalar field with positive energy density and equation of state $p=w\\rho$, this corresponds to a violation of the null energy condition, $\\rho+p \\geq 0$, a reasonable assumption according to which light rays are focused by matter.  However, negative kinetic energies arise in supersymmetric models and higher-derivative-gravity theories \\cite{nil84,pol88}, while string models can describe brane physics in which the effective 4D null energy condition is not preserved \\cite{CM}; last but not least, anti--de Sitter configurations do violate the dominant energy condition. Until now, the standard lore of well-established energy conditions has been adopted. What about abandoning the old path in favour of more speculative scenarios? In particular, can some of the most popular objections against embarrassing forms of matter be circumvented \\cite{pha6}? Recently, many people have been considering scenarios in which the dark energy content of the observable Universe is of a non-conventional nature, namely, violating the null energy condition (see \\cite{pha1,pha2,pha3,pha4,pha7,pha8,pha11,pha12,pha13,pha14,pha15,pha16,pha17,pha19,pha24,pha25,pha26,pha27} for the case of a non-canonical scalar field and \\cite{pha18,pha20,pha23} for the tachyon case). This ``phantom cosmology''\\footnote{The zoology of ``phantom physics'' would be quite rich and, indeed, drawn by modifications of just two kinds of fields. Let $[\\text{sgn}(\\rho),\\,\\text{sgn}(w+1),\\,\\text{sgn}(B)]$ describe a field with equation of state $\\rho+p=\\rho(1+w)$ and kinetic energy proportional to the parameter $B$.  Given a scalar field with energy density $\\rho = B\\dot{\\phi}^2/2+V$, violations of the null energy condition can be achieved with either $[+,-,-]$ (type-I phantoms, commonly known as dark energy ``phantoms'' or ``ghosts'') or $[-,+,-]$ (type-II phantoms). Given a tachyon field with $\\rho+p=\\rho B\\dot{T}^2$, violations corresponds to either $[+,-,-]$ or $[-,+,+]$ (semiphantoms). Conversely, if $\\rho+p>0$, then one has either standard matter $[+,+,+]$, type-III phantoms $[-,-,+]$ (scalar case), or superphantoms $[-,-,-]$ (tachyon case).} \\cite{pha1}, while successfully exploring previously forbidden regions in the space of parameters, does not avoid criticisms \\cite{pha9,pha28}. Anyhow, it would be interesting to ask what is the phenomenology of an early-universe short phantom era with ultra-negative equation of state and generalized Friedmann equation.  From a mathematical point of view, the phantom universe displays interesting properties such as the presence of a finite-time singularity when $w$ in constant [big smash or big rip \\cite{pha1,pha5,pha10}, see Eq. (\\ref{conw}) with $w<-1$] and a correspondence \\cite{chi02,agu03,pha21,pha22} resembling (but not equivalent to) the scale-factor duality of pre-big-bang cosmology which is a symmetry of the low-energy string effective action and is achieved with the mapping $a(t)\\rightarrow 1/a(-t)$ \\cite{ven91,GV1} (for some reviews on string and pre-big-bang cosmology, see \\cite{gas99,LWC,GV2}). Also patch cosmologies with negative $q$ have a finite-time singularity with divergent scale factor, even if the density evolution shows the opposite trend. This fact, together with the non-canonical effective theory which seems to characterize such models, invites to investigate if there is some relation between $q<0$ patch cosmologies and models with phantom fluids or, more generally, if suitable low-energy theories with dilatonic coupling may reproduce this class of effective Friedmann equations.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "0402/astro-ph0402655_arXiv.txt": {
        "abstract": "{ Photoevaporation may provide an explanation for the short lifetimes of disks around young stars. With the exception of neutral oxygen lines, the observed low-velocity forbidden line emission from T Tauri stars can be reproduced by photoevaporating models. The natural formation of a gap in the disk at several AU due to photoevaporation and viscous spreading provides a possible halting mechanism for migrating planets and an explanation for the abundance of observed planets at these radii. }    \\addkeyword{accretion} \\addkeyword{accretion disks} \\addkeyword{stars: formation}  \\begin{document} ",
        "introduction": "\\label{sec:intro}  Most low-mass stars form surrounded by circumstellar disks. Observations of infrared excess (e.g., Haisch, Lada, \\& Lada 2001) suggest disk lifetimes $\\tau_{\\rm disk} \\approx 6 \\times 10^6\\,$yr, requiring an efficient mechanism for disk dispersal.  Hollenbach, Yorke, \\& Johnstone (2000) considered a variety of disk dispersal mechanisms and concluded that viscous accretion of the disk onto the central star (e.g. Hartmann et al.\\ 1998) together with photoevaporation of the disk at moderate radii (Shu, Johnstone, \\& Hollenbach 1993; Hollenbach et al.\\ 1994; Johnstone, Hollenbach, \\& Bally 1998) must act together efficiently to remove the entire disk. Numerical calculations by Clarke, Gendrin, \\& Sotomayor (2001) and Matsuyama, Johnstone, \\& Hartmann (2003a) have shown that these processes may remove the disk in $10^{5-7}\\,$yr.  The photoevaporation model fits observational data well in the case of external heating via nearby massive stars (Bally et al.\\ 1998; Johnstone, Hollenbach, \\& Bally 1998; St\\\"orzer \\& Hollenbach 1998).  With the exception of a few cases (e.g., MWC349A), evidence for disk photoevaporation due to the central star (Shu et al.\\ 1993; Hollenbach et al.\\ 1994) is largely circumstantial. Observations of blue-shifted, low-velocity emission from forbidden lines of oxygen, nitrogen, and sulfur in the spectra of many T Tauri stars (Hartigan, Edwards, \\& Ghandour 1995) provide useful diagnostics. Font et al.\\ (2004) calculated the flow properties of the photoevaporative disk wind and found them to compare reasonably with the strengths and profiles of the nitrogen and sulfur lines. The oxygen lines, however, are underabundant in the model.  Along with disk dispersal, photoevaporation and viscous accretion produce structure in the disk such as gaps and rings. Matsuyama, Johnstone, \\& Murray (2003b) showed that the formation of gaps within the gaseous disk during the dispersal era places constraints on the migration of planetary orbits.  The following sections review the key results of Matsuyama et al.\\ (2003a,b) and Font et al.\\ (2004). ",
        "conclusions": "\\label{sec:conc} Photoevaporation of disks around young stars may be responsible for short observed disk lifetimes, especially when powered by FUV  or ionizing radiation from nearby massive stars. Together, photoevaporation and viscous accretion naturally lead to the formation of gaps and ring structures within disks, without the need to invoke unseen planets.  Hydrodynamic models of the photoevaporative disk wind reveal that the launching point for the flow, $r_g$, has been overestimated in the analytic models by a factor of $\\sim\\,3$ and that the launch velocity has been also somewhat overestimated.  While most observed low-velocity forbidden line profiles are reasonably matched by the model, the predicted neutral oxygen forbidden line flux is too low since almost all oxygen in the model is ionized.  Interestingly, a hot $10^4\\,$K wind which is not photoionized would contain mostly neutral oxygen and the line strengths would produce an excellent fit to the observations.  The existence of a gap in the disk at $r_g$ provides a halting mechanism for migrating planets. Given the predicted location of $r_g$ at several AU from the central star, such gaps might explain the recent abundance of observed planets with these radii."
    },
    "0402/astro-ph0402149_arXiv.txt": {
        "abstract": "We review the evidence for substructures from the anomalous flux ratios in gravitational lenses. Using high-resolution numerical simulations, we show that at typical image positions, the fraction of surface mass density in substructures is $\\la 0.5\\%$ with mass above $10^{-4}$ virial masses in the ``concordance'' $\\LCDM$ cosmology. Substructures outside the virial radius (but projected at typical lens image positions) only increase the fraction moderately. Several effects, in particular baryonic settling and the requirement of compactness, may further decrease the predictions by a factor of few. The predicted fraction with appropriate properties thus appears to be lower than that required by lensing, although both are still uncertain. More speculative substructures such as massive black holes ($M \\sim 10^5-10^6 M_\\odot$) in the halo may offer viable alternatives. ",
        "introduction": "Gravitational lenses on arcsecond scales provide a unique sample to probe the mass distribution in the lensing galaxies at intermediate redshift ($z\\sim 0.5-1$). Image positions in most lenses can be fitted adequately using simple smooth galaxy mass models. But observed flux ratios are  more difficult to match (e.g., Kochanek 1991). The discrepancy between the predicted and observed flux ratios is commonly referred to as the ``anomalous flux ratio problem.'' The most apparent cases are found in quadrupole lenses where we observe a close pair or a close triple of images. Here we know that the lensed source is close to either a fold or a cusp caustic. The {\\it asymptotic} magnification behavior in such cases is well understood -- a close pair must have equal brightness, while for a close triple, the flux of the middle image should be equal to the total fluxes of the two outer images. Virtually all the observed pairs and triples disagree with these relations (\\S2). This has been argued as evidence for substructures on the scale of the separations of the images (a few tenths of arcseconds, e.g., Mao \\& Schneider 1998; Metcalf \\& Zhao 2002). Another piece of evidence for substructures is the fact that saddle images are preferentially dimmed compared to model predictions (Kochanek \\& Dalal 2003). This is expected from milli-lensing by substructures (Keeton 2003) or microlensing by stars (Schechter \\& Wambsganss 2002)\\footnote{We refer to lensing by substructures and lensing by stars as milli-lensing and microlensing respectively because the angular scales involved are $\\sim$ mas and $\\mu$as in these two cases.}; such a preferential de-magnification of saddle images is not expected from other propagational effects.  The Cold Dark Matter (CDM) structure formation model predicts the existence of just such substructures from both semi-analytical studies and numerical simulations (e.g., Kauffmann et al. 1993; Klypin et al. 1999; Moore et al. 1999; Ghigna et al. 2000). About 5-10\\% of the mass is predicted to be in substructures with a mass spectrum of $n(M)dM \\sim M^{-1.8}dM$. Intriguingly, the predicted  number of subhaloes in CDM exceeds the observed number of {\\it luminous} satellite galaxies in a Milky-Way type galaxy (e.g., Klypin et al. 1999; Moore et al. 1999; see also Stoehr et al. 2002). One solution for this disagreement may be that some substructures, especially those of lowest mass, are dark. If this is true, then gravitational lensing may be the best way to detect them. In this paper, we examine the required amount of substructures in gravitational lenses (\\S2) and compare it with the predictions from CDM (\\S3). Finally, in \\S4, we discuss several effects that affect the predictions of numerical simulations. Throughout this paper, we adopt the ``concordance'' $\\LCDM$ cosmology (e.g., Ostriker \\& Steinhardt 1995; Spergel et al. 2003 and references therein), with a density parameter $\\omega0=0.3$, a cosmologically constant $\\lambda0=0.7$, a baryon density parameter $\\Omega_{\\rm b}=0.024h^{-2}$, and we take the power-spectrum normalization $\\sigma_8=0.9$. We write the Hubble constant as $h=H_0/(100\\kms\\mpc^{-1})$ with $h=0.7$. ",
        "conclusions": "We have reviewed the evidence for substructures from close pairs and triples in quadruple lenses. As emphasized by Kochanek \\& Dalal (2003), the fact that saddle images are frequently dimmer than expected is difficult to accommodate by other means. Quantitatively, the anomalous flux ratios in lenses appear to require a few percent of the surface mass density in substructures at typical image positions (Dalal \\& Kochanek 2002; Metcalf et al. 2003). The required fraction is higher than that provided by globular clusters and luminous satellite galaxies (Mao \\& Schneider 1998; Chiba 2002) and it also appears to be higher than the predicted values ($\\fsub \\la 0.5\\%$) from the $\\LCDM$ cosmology at typical image positions. However, at present it is unclear how serious the discrepancy is because of uncertainties in both observations and theoretical predictions.  There are a number of issues that need to be understood better in current numerical simulations. Even the basic question of the identification of substructures  needs to be explored further. The {\\tt SUBFIND} algorithm we adopted only identifies bound subhaloes, however, tidal streams from disrupted systems (for examples in the Milky Way, see, e.g., Ibata et al. 2001; Yanny et al. 2003) may also contribute to the budget of substructures. Another important issue is whether our results have achieved  convergence  as a function of the spatial and mass resolutions. New simulations are underway to address this issue. Presumably when the numerical resolution becomes higher, the inner parts of subhaloes are resolved better into higher-density regions that can survive tidal disruptions longer. However, the survival of substructures may be linked to another small-scale problem of CDM: the central density profiles of low-surface brightness galaxies (usually with circular velocities of $\\la 100\\kms$) seem to be too concentrated compared with observed galaxies (e.g., Bolatto et al. 2002; Weldrake et al. 2003; see, however, Swaters et al. 2003). Therefore, if we put in observed mass profiles, substructures may actually be more easily destroyed by tidal forces due to their lower central concentrations. There is another effect that makes the survival of substructures in the central part more difficult. In collisionless numerical simulations, the density profile can be approximated by an NFW profile; the density scales as $\\propto r^{-1}$ out to $\\sim 0.1\\rv$ ($\\sim 25$ kpc) for typical galactic-sized haloes. However, the observed velocity dispersion is nearly constant in the inner part implying $\\rho \\propto r^{-2}$, i.e., the density in real galaxies rises more quickly as the radius decreases. Hence substructures will be disrupted more easily if they come close to the center and dynamical frictions dragging them into the center would also be larger in realistically simulated galaxies.  Our simulations resolve substructure masses from $10^8 M_\\odot$ to $10^{11} M_\\odot$ for a $10^{12}M_\\odot$ parent halo. However, according to Metcalf et al. (2003), the required substructure mass is in the range of $10^4M_\\odot-10^8 M_\\odot$ in the case of 2237+0305. If the mass spectrum of substructures $n(M)dM \\propto M^{-1.8} dM$ extends all the way down to $10^4 M_\\odot$, one can estimate that the fraction of mass in substructures with $10^4M_\\odot<M <10^8 M_\\odot$ is about a factor of 3 smaller than that in substructures with $10^8M_\\odot < M <M^{11}M_\\odot$, i.e., the mass fraction in substructures in the range required by Metcalf et al. (2003) will be even lower than our predicted value. This conclusion depends on the mass spectrum. However, for any power-law spectrum with a slope shallower than $-2$, the mass fraction in substructures with $10^4 M_\\odot<M <10^8 M_\\odot$ is likely smaller than that in higher-mass substructures identified in numerical simulations.  There is another effect that may reduce the utilizable mass in substructures even further: the need for compactness, an issue we already touched upon in \\S2 (see Fig. 1). In order to affect the flux ratios significantly, the physical size of substructures must be sufficiently compact (cf. Metcalf et al. 2003). The most efficient substructures should be of the same order of the image separation of the pairs or triples. For the five CLASS quadrupole lenses that have known lens and source redshifts (see \\S4), the closest pairs have projected separations from $0.6h^{-1}$kpc (for B1422+231) to $2.3h^{-1}$kpc (for MG0414+0534). At redshift of 0.5, dark matter haloes with $M \\la 4\\times 10^6M_\\odot$ will have $\\rv \\sim 1h^{-1}$kpc, and so they will be efficient in causing flux anomalies if they are located between close pairs or triples. For larger masses, the effect is reduced. We estimate the reduction by assuming that only the mass enclosed within $1h^{-1}$kpc will affect the flux ratios. We adopt a mass spectrum for dark matter haloes of $n(M) dM \\propto M^{-1.8}$ for $10^4M_\\odot < M < 10^{11} M_\\odot$. Each halo is described by an NFW profile with a concentration parameter given by $c\\approx 11(M/10^{12}h^{-1})^{0.15}$ (e.g., Zhao et al. 2003). Note that the substructures in the central parts are tidally truncated and so cannot be fitted well by NFW profiles, but as most substructures have $r>0.1\\rv$ in spherical radius, the effect of tidal truncation may be modest. We find that the mass in substructures that can efficiently cause flux anomalies is reduced by an additional factor of five compared with the total mass in all substructures -- most substructures which are in the outer parts are too extended to cause flux anomalies efficiently. Our rough estimate shows that the compactness requirement is an issue that needs to be addressed more carefully. The two additional effects noted can reduce the likely mass fraction in substructures having the required masses and sizes to as low as $\\sim 0.03\\%$, uncomfortably low compared with the observational requirement.  Progress can be made from both the observational and theoretical fronts to reduce the uncertainties. Observationally, in the radio, the effect of scattering can be studied with observations at high frequency where it is expected to be unimportant. In the infrared, it would be interesting to have more observations with integral field spectroscopy similar to that reported by Metcalf et al. (2003). This method offers a way to separate stellar microlensing from substructure milli-lensing. More astrometric signatures of substructures will be important as well (see \\S2). Theoretically, higher-resolution simulations are needed and are already under-way. If future observations and numerical simulations still indicate a discrepancy between lensing requirements and CDM predictions, then alternatives must be sought. One possibility is that these substructures are massive black holes with $M \\sim 10^5-10^6 M_\\odot$ (Lacey \\& Ostriker 1985; Xu \\& Ostriker 1994), which satisfy the mass and compactness requirements. We require only a few percent of the surface density in the substructures, so the density parameter in these black holes is only $\\fsub\\Omega_0 \\sim 0.012 (\\fsub/0.04)$, which is about $30\\%(\\fsub/0.04)$ of the baryon density in the universe. These massive black holes will have other observable signatures (e.g., Wambsganss \\& Paczy\\'nski 1992; Tremaine \\& Ostriker 1999) and can be further tested."
    },
    "0402/astro-ph0402463_arXiv.txt": {
        "abstract": "{ Relevant studies of the non-thermal components of the intracluster medium are  performed at radio wavelengths. A number of clusters, indeed, exhibits cluster-wide diffuse radio emission, which is indication of the existence of large scale magnetic fields and of relativistic electrons in the cluster volume. There is strong evidence that the presence of diffuse radio emission is related to cluster merger processes. The details of the halo-merger connection are discussed and a brief outline of current models of halo formation is presented. ",
        "introduction": "\\begin{figure*} \\centering \\includegraphics[width=16cm]{fig1.ps} \\caption{Diffuse radio emission in the Coma cluster, obtained at 90 cm with the Westerbork Synthesis Radio Telescope. The discrete sources have been subtracted. The cluster center is approximately located at the position RA$_{1950}$ = 12$^h$ 57$^m$ 24$^s$, DEC$_{1950}$ = 28\\degrees 15\\arcmin 00\\arcsec. The radio halo Coma C is at the cluster center, the radio relic 1253+275 is at the cluster periphery. An angular size of 10\\arcmin~ corresponds to a linear size of $\\sim$ 400 kpc. Contour levels are at 2.5, 4, 8, 16 mJy/beam (FWHM = 55\\arcsec $\\times$ 125\\arcsec; RA $\\times$ DEC). } \\label{Fig1}% \\end{figure*}   The main components of clusters of galaxies are the galaxies (2-3\\%), the hot gas (13-15\\%) and the dark matter (82-85\\%).  In addition, a relativistic component may be present which plays an important role in the cluster formation and evolution.  The most detailed studies of this component come from the radio observations.  A number of clusters of galaxies is known to contain large-scale diffuse radio sources which have no obvious connection with the cluster galaxies, but are rather associated with the intracluster medium (ICM).  These sources are classified in two groups, radio halos and relics, according to their location at the cluster center or cluster periphery, respectively. The synchrotron origin of the emission from these sources requires the presence of cluster-wide magnetic fields of the order of $\\sim$ 0.1-1 $\\mu$G, and of a population of relativistic electrons with Lorentz factor $\\gamma >>$ 1000 and energy density of $\\sim$ 10$^{-14}$-10$^{-13}$ erg cm$^{-3}$.   The importance of halos and relics is that they are large scale features,  related to other cluster properties in the optical and X-ray domains, and thus  connected to the cluster history and evolution.    These sources are found in clusters which have recently undergone a merger event, thus leading to the idea that they originate from particle acceleration in cluster merger turbulence and shocks.  The formation and evolution of these sources is however still under debate: the radio emitting electrons could be reaccelerated cosmic rays, or accelerated from the thermal population, or could be produced as a result of the interaction between cosmic-ray protons and the ICM. We summarize the current knowledge on these sources from an observational point of view.  The instrinsic parameters quoted in this paper are computed with a Hubble constant H$_0$ = 50 km s$^{-1}$ Mpc$^{-1}$ and a deceleration parameter q$_0$ = 0.5.  \\begin{figure} \\centering \\includegraphics[width=7cm]{fig2.eps} \\caption{Total radio spectrum of the radio halo Coma C (from Thierbach et al. 2003). } \\label{Fig2} \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=7cm]{fig3.eps} \\caption{Total radio spectum of the relic 1253+275 (from Thierbach et al. 2003). } \\label{Fig3} \\end{figure}  \\section {Radio halos and relics: the Coma cluster}  The Coma cluster is the first cluster where a radio halo and a relic have been detected (Willson 1970, Ballarati et al. 1981). The halo in this cluster, Coma C (see Fig. 1), is the prototypical example of halo sources: it is located at the cluster center, it is characterized by a regular shape with a total extent of $\\sim$ 1 Mpc, and by a low radio surface brightness (\\ltsim~ $\\mu$Jy arcsec$^{-2}$ at 1.4 GHz).  It is unpolarized down to a limit of a few percent, and shows a steep radio spectrum, typical of aged radio sources ($\\alpha$ \\gtsim~ 1), with a steepening at higher frequencies (Fig. 2).  The spectral index distribution of Coma C shows a radial decrease (Giovannini et al. 1993) from $\\alpha \\sim$ 0.8 at the cluster center, to $\\alpha \\sim$ 1.8 beyond a distance of about 10\\arcmin.  By assuming that there is energy equipartition between relativistic particles and magnetic field, a minimum energy density of 1.62 10$^{-14}$ erg cm$^{-3}$ is derived from the radio data.  The corresponding equipartition magnetic field is 0.4 $\\mu$G.  The radio source  1253+275 in the Coma cluster (Fig. 1) is the prototype of the class of radio relics, which are extended diffuse radio sources associated with the ICM, located in the cluster peripheral regions. This source is similar to the halo Coma C in its low surface brightness, large size and steep spectrum (Fig. 3). Unlike halos, it shows an elongated structure and it is highly polarized ($\\sim$ 25\\%).   \\begin{figure*} \\centering \\includegraphics[bb=105 185 550 630,width=9cm,clip]{fig4.ps} \\caption{The cluster A2163 in radio and X-ray.  The grey scale image represents the radio emission in A2163 at 20 cm, showing an extended radio halo.  The contours represent the ROSAT X-ray emission. The extended irregular X-ray structure indicates the presence of a recent cluster merger. } \\label{Fig4} \\end{figure*}  The radiative lifetime of the relativistic electrons in Coma C, considering synchrotron and inverse Compton energy losses, is of the order of 10$^8$ yr. This is too short to allow the particle diffusion throughout the cluster volume. This implies that the radiating electrons cannot have been injected at some particular point of the cluster, but they must undergo {\\it in situ} energization. This is a general problem for all the halo and relic sources. Feretti (2002) argued that halos and relics are not the same objects seen in projection, i.e. halos are really at the cluster center and not simply projected onto it.  Halos and relics may indeed have different physical origins.  Coma is one of the few clusters where hard X-ray emission has been detected with the BeppoSAX and  Rossi X-ray Timing Explorer (RXTE) satellites (Fusco-Femiano et al. 2004 and references therein). This emission is expected in clusters with diffuse radio sources, as the high energy relativistic electrons responsible for the radio emission ($\\gamma$ $\\sim$ 10$^4$) scatter off the cosmic microwave background, boosting photons from this radiation field to the hard X-ray domain by inverse Compton (IC) process. Measurements of this radiation provide additional information that, combined with results of radio measurements (i.e. the ratio of hard X-ray IC emission to radio synchrotron emission), enables the determination of the electron density and mean magnetic field directly, without invoking equipartition.  The 20-80 keV flux in Coma is $\\sim$ 1.5 10$^{-11}$ erg cm$^{-2}$ s$^{-1}$, which leads to a volume averaged intracluster magnetic field of $\\sim$ 0.2 $\\mu$G (Fusco-Femiano et al. 2004). This value is consistent with that obtained by the radio emission assuming equipartition (see above). It is inconsistent, however, with the value of $\\sim$ 6 $\\mu$G deduced by Feretti et al. (1995) from Faraday Rotation Measure (RM) data (see Sect. 4).  It is worth mentionining here that alternative models have been suggested to explain  the hard X-ray tails (e.g. non-thermal bremsstrahlung).  These models were motivated by the discrepancy between the value of the ICM magnetic field derived by the IC model and the value derived from RM.  However, these models may have serious difficulties as they would require an unrealistic high energy input (Petrosian 2001). ",
        "conclusions": "The existence of cluster-wide diffuse radio emission indicates that there are important non-thermal components in the ICM: magnetic fields and  relativistic particles. The presence of magnetic fields in galaxy clusters is additionally demonstrated by RM studies, which indicate that magnetic fields are rather common in all clusters, not only those with radio halos. There is convincing evidence that radio halos and relics are linked to cluster merger processes. Violent mergers provide the energy necessary to reaccelerate the radio emitting electrons."
    },
    "0402/astro-ph0402239_arXiv.txt": {
        "abstract": "We discuss free-free radio emission from ionized gas in the intergalactic medium.  Because the emissivity is proportional to the square of the electron density, the mean background is strongly sensitive to the spatial clumping of free electrons.  Using several existing models for the clumping of ionized gas, we find that the expected free-free distortion to the cosmic microwave background (CMB) blackbody spectrum is at a level detectable with upcoming experiments such as the Absolute Radiometer for Cosmology, Astrophysics, and Diffuse Emission (ARCADE).  However, the dominant contribution to the distortion comes from clumpy gas at $z \\la 3$, and the integrated signal does not strongly constrain the epoch of reionization.  In addition to the mean emission, we consider spatial fluctuations in the free-free background and the extent to which these anisotropies confuse the search for fluctuations in 21 cm line emission from neutral hydrogen during and prior to reionization.  This background is smooth in frequency space and hence can be removed through frequency differencing, but only so long as the 21 cm signal and the free-free emission are uncorrelated.  We show that, because the free-free background is generated primarily at low redshifts, the cross-correlation between the two fields is smaller than a few percent.  Thus, multifrequency cleaning should be an effective way to eliminate the free-free confusion. ",
        "introduction": "Plans for upcoming low-frequency radio experiments able to measure the 21 cm background associated with neutral hydrogen during, and prior to, reionization era (e.g., \\cite{Scott90} 1990; \\cite{Madetal97} 1997; \\cite{Zaletal03} 2003; \\cite{MorHew03} 2003) have motivated study of the foregrounds that may contaminate such measurements. Synchrotron emission from the Milky Way, low frequency radio sources (\\cite{DiMetal02} 2002), and free-free emission from free electrons in the intergalactic medium (\\cite{OhMac03} 2003) are now thought to be the chief sources of confusion.  While the free-free background contaminates 21 cm studies, the emission itself captures important physics of the ionized component of the intergalactic medium (IGM). In particular, because the emissivity is proportional to the square of the electron number density, free-free emission is strongly sensitive to whether electrons are spatially clumped or distributed smoothly in the IGM. Initial estimates of the free-free background suggest that gas clumping boosts the specific intensity by a factor from order unity at $z\\sim20$ to over 100 at $z \\la 3$ (\\cite{Oh99} 1999).  Here, we consider the direct detection of free-free emission through the distortion it creates in the cosmic microwave background (CMB) blackbody spectrum at low radio frequencies (\\cite{Loeb96} 1996; \\cite{Oh99} 1999).  We calculate the mean distortion temperature using a variety of both analytic and numeric clumping models already existing in the literature.  While the expected brightness varies by factors of a few between the different models, we suggest that the distortion to the CMB is at a level detectable with the planned Absolute Radiometer for Cosmology, Astrophysics, and Diffuse Emission (ARCADE; \\cite{Kog03} 20003) experiment. A detection of the distortion could help to discriminate between existing clumping models and to constrain the integrated ionization history of the universe. We also show that ionized halos at $z \\la 3$ dominate the background, with only a relatively insignificant contribution from the reionization era.  In addition to the mean background, we estimate the magnitude of spatial fluctuations in the free-free intensity that appear because the ionized clumps are a biased tracer of the dark matter density field. We find that the fluctuations are comparable to or greater than those from the 21 cm signal on the relevant scales.  Fortunately, as shown by \\cite{Zaletal03} (2003), the smoothness of the free-free background in frequency space should allow one to clean these sources in 21 cm maps (see also \\cite{MorHew03} 2003); the cleaning is quite efficient but relies on the foreground and 21 cm signals being uncorrelated.  In fact the two should be anticorrelated, because the free-free emission comes from ionized halos while the 21 cm signal comes from neutral regions.  Here we show that, because nearly all of the free-free background arises at $z \\la 3$, the anti-correlation is at the level of few percent, suggesting that adequate cleaning is possible.  The discussion is organized as follows: in \\S 2, we briefly discuss free-free emission from ionized halos and its potential detectability as a distortion to the CMB.  In \\S 3, we discuss the detectability of spatial fluctuations in the free-free background and the extent to which free-free emission contaminates 21 cm measurements. Throughout the paper, we make use of the WMAP-favored LCDM cosmological model (\\cite{Speetal03} 2003).  \\begin{figure}[t] \\psfig{file=Fig1v3.eps,width=3.5in,angle=0} \\caption{(a) The clumping factor of electrons in the IGM as a function of redshift.  These clumping factors come from \\cite{Haietal01} (2001) (long-dashed line: a reionization model that includes minihalos as well as star-forming halos with $T_v > 10^4$~K), \\cite{Benetal01} (2001) (dashed line: semi-analytic model including cool neutral gas, or galaxies, in ionized halos as well as the clumping of gas outside halos), and \\cite{GneOst97} (1997) (dot-dashed line: a direct measurement in numerical simulations).  The dotted line is the clumping factor presented in \\cite{Oh99} (1999), while the solid line is a similar calculation described in the text based on the star formation history of the universe.  (b) The mean brightness temperature of free-free emission at an observed frequency of 2 GHz. Here, we show both the differential (thin lines) and cumulative brightness temperatures (thick lines) as a function of the redshift. The curves refer to the clumping models in the top panel. For reference, we show the case where $C(z)=1$ and the ionized gas distribution is taken to be spatially smooth.  We also show the current limit on the free-free distortion of the CMB (top line with $Y_{\\rm ff} \\equiv T_\\nu/T_{\\rm CMB} (h\\nu/kT_{\\rm CMB})^2 < 1.9 \\times 10^{-5}$~K; \\cite{Beretal94} 1996) and the expected constraint from the ARCADE experiment (\\cite{Kog03} 2003).} \\label{fig:clump} \\end{figure} ",
        "conclusions": "To summarize, we find that the cumulative specific intensity of the free-free background at low radio frequencies is strongly sensitive to the spatial clumping of free electrons in the IGM. Using a variety of existing models for the clumping of ionized gas and its redshift evolution, we find that the expected free-free distortion to the CMB blackbody spectrum, at frequencies of order a few GHz and below, is within the reach of upcoming experiments such as ARCADE.  The dominant contribution to the background is from ionized gas at $z \\sim 3$, because the clumping factor is generally large below this redshift. The free-free background varies across the sky at the level of a few to at most ten percent.  These fluctuations can be detected with experiments such as the SKA, though a careful separation of CMB anisotropies and other low-frequency radio point sources will be required.  At frequencies of order 200 MHz, where observations of 21 cm radiation from neutral hydrogen during and before the epoch of reionization are planned, angular fluctuations from free-free emission could become a significant source of confusion.  However, we have shown that the cross-correlation between the 21 cm and free-free signals is small, suggesting that the free-free foreground can be removed to good accuracy."
    },
    "0402/astro-ph0402525_arXiv.txt": {
        "abstract": "{ The Rees-Sciama effect produced in mergers of galaxy clusters is discussed, and an analytical approximation to compute this effect from numerical simulations is given. Using this approximation and a novel toy model describing the physics of the merger, we characterize the spatial properties and symmetries of the Rees-Sciama signal. Based on these properties, we propose a method to extract the physical parameters of the merger, which relies on the computation of the quadrupole moment of the observed brightness distribution on the sky. The relationships between the quadrupole coefficients and the physical parameters of the merger (physical separation, projection angle on the sky and angular momentum) are discussed. Finally, we propose a method to co-add coherently the RS signals from a sample of cluster mergers, in order to achieve an statistical detection of the effect for those cases where individual signals are masked by the kinetic SZ effect, the primordial CMB components, and by observational noise. ",
        "introduction": "}  The integrated Sachs-Wolfe \\citep[ISW;][]{1967ApJ...147...73S, 1994PhRvD..50..627H} effect produces CMB temperature fluctuations due to the accumulation of red- or blue-shift of photons travelling in time dependent gravitational potentials. Several recent studies \\citep{2003astro.ph..5001B,2003astro.ph..5097N,2003astro.ph..5468F} have claimed evidence for its detection in the WMAP data \\citep{bennett03}. The non-linear contribution to the ISW effect, in which the density contrast producing the gravitational potential is in its non-linear regime, is usually called the Rees-Sciama \\citep[RS;][]{rees_sciama68, seljak96} effect.  Here we discuss the RS effect in the extremely non-linear regime of present day galaxy cluster mergers. The regime of moving single galaxy clusters was already examined by \\citet{birkinshaw83}, \\citet{1998A&A...334..409A}, \\citet{2000ApJ...537..542M, 2003ApJ...586..731M}, and \\citet{2002PhRvD..65h3518C}.  However, galaxy clusters reach the largest velocities, and therefore the strongest RS signal, during mergers in which two or more clusters are invoked.  Due to the slow centre-of-mass velocity of the merging system, the RS signals of the merging subclusters can partly cancel or increase each other, depending on the merger geometry and the location of the line-of-sight (LOS).  Furthermore, the spectral signature of the RS effect is a temperature change of the measured CMB photons, and therefore indistinguishable form the kinematic Sunyaev-Zeldovich \\citep[kSZ;][]{1972ComAp...4..173S,1980ARA&A..18..537S} effect spectrally. However, both effects have distinct morphologies, which in principle allow to discriminate them.  The focus of this work is to examine methods to extract the RS signature from merging pairs of galaxy clusters. The layout of the paper is as follows. In Sec. 2 we provide an approximation to compute the RS effect either from theoretical cluster models or numerical simulations. In addition, a brief description of the kSZ and the thermal Sunyaev-Zeldovich (tSZ) effects is also given, emphasizing their phenomenological differences with the RS effect. In Sec. 3 we present and use an analytic toy model of a merger of two galaxy clusters to characterize the typical morphologies of these effects (RS, kSZ and tSZ). Based on its spatial properties, we present in Sec. 4 a method to extract both the RS signal and the physical parameters describing the system from a given merger. A method to stack the signal from a sample of mergers is also proposed. Conclusions are presented in Sec. 5. ",
        "conclusions": "In this paper we have presented a formalism that can be easily incorporated to N-body simulation codes in order to predict the Rees-Sciama effect in a merging systems of clusters of galaxies. For the typical range of velocities in those systems ($v/c \\ll 1$), the obtained expression (Eq. \\ref{eq:dTRS2}) can also be seen as a gravitational lensing effect produced by a moving lens.  Using simple modelling for the cluster merger, we have illustrated the morphology and symmetries of this effect, and we have developed a method to extract the signal, which should be applicable to realistic maps. This method is based on the computation of the (weighted) dipole and quadrupole moments of the brightness distribution. In particular, it has been shown how the quadrupole moment is related to the kinematic properties of the merger, so we can extract information about the dynamical state of the system: pre-merger or post-merger, and the magnitude of the angular momentum.  Since we expect that in the near future the observation of a single cluster merger will be extremely difficult, given the weak signal strength, we have proposed a simple method of stacking the signal from a large number of clusters in order to extract their RS signature statistically. The procedure is straight-forward: for a sample of clusters (e.g. a complete, or a merging cluster sample) a coordinate system is attached to the center of light (X-ray or tSZ map) so that the X-axis is aligned with the major elongation of the gas. Then the quadrupole moments of the CMB temperature fluctuations (outside the gas region, defined by the tSZ effect) are calculated for each cluster. Finally the individual quadrupole moment components are co-added for the sample. Intrinsic CMB fluctuations and kSZ effect residuals should cancel out statistically, leaving only a signature of the RS effect.  A detailed investigation of this stacking method will be carried out in a follow-up paper."
    },
    "0402/astro-ph0402341.txt": {
        "abstract": "Stellar bars drive gas  into the circumnuclear (CN) region of galaxies. To investigate the fate of the CN gas and star formation (SF), we study a sample of barred non-starbursts  and  starbursts with high-resolution CO, optical, H$\\alpha$, RC,  Br$\\gamma$, and  $HST$ data, and find the following. (1) The inner kpc of bars differs markedly from the outer disk. It hosts molecular gas  surface densities $\\Sigma_{\\rm gas-m}$ of 500-3500  M$_{\\tiny \\sun}$ pc$^{-2}$,  gas mass fractions of  10 to 30\\%, and epicyclic frequencies of several 100--1000 km~s$^{-1}$~kpc$^{-1}$. Consequently, in the CN region, gravitational instabilities can only grow at high gas densities and on short timescales, explaining in part why powerful starbursts reside there. (2) Across the sample,  we find bar pattern speeds with upper limits of  43 to 115 km~s$^{-1}$~pc$^{-1}$  and  outer inner Lindblad resonance (OILR) radii of $>$ 500 pc. (3) Barred starbursts and  non-starbursts have CN SFRs of  3--11 and 0.1--2 M$_{\\tiny \\sun}$ yr$^{-1}$, despite similar CN gas mass. $\\Sigma_{\\rm gas-m}$ in the starbursts is larger  (1000--3500 $M_{\\tiny \\sun}$ pc$^{-2}$) and close to the  Toomre critical density over  a large region. (4) Molecular gas makes up 10\\%--30\\% of the  CN dynamical mass, and fuels large CN SFRs in the starbursts, building young, massive, high $V/\\sigma$ components. Implications for secular  evolution along the Hubble sequence are discussed. \\end {abstract} ",
        "introduction": "It is widely recognized that  non-axisymmetries,  such as large-scale stellar bars, %such as  those of the galactic halos and bulges) facilitate  the radial transfer of angular momentum in disk galaxies, thus  driving their dynamical and secular evolution  (e.g., Pfenniger \\& Norman 1990; Friedli \\& Benz 1995; Athanassoula 1992; Kormendy 1993;  Kormendy \\&  Kennicutt 2004; review by Jogee 2004 and references therein). Numerous studies (e.g., Knapen et al. 2000; Eskridge et al. 2000) have  shown  that the  majority ($>$ 70 \\%) of nearby disk galaxies  host large-scale stellar bars. Using $HST$ Advanced Camera for Surveys (ACS) data, Jogee \\etal (2004a,b)  find that strong bars remain frequent from the present day out to lookback times of 8 Gyr ($z \\sim$~1), and infer that bars are long-lived features with a lifetime well above 2 Gyr.  A non-declining optical bar fraction out to $z \\sim$~1 is also confirmed by Elmegreen \\etal  (2004). These findings suggest that bars can influence a galaxy over a significant part of its lifetime. Compelling evidence that   bars efficiently drive  gas from the outer disk into the inner kpc comes from the larger central molecular  gas concentrations observed in barred galaxies compared to unbarred galaxies (Sakamoto \\etal 1999).     However, few  high resolution studies based on  large samples exist of the fate of gas once it reaches the inner kpc of a bar. In this study, we  use  high resolution  ($2 \\arcsec$; $\\sim$~200 pc) CO  ($J$=1--0) observations, optical, NIR, H$\\alpha$, radio continuum (RC), Br$\\gamma$,  and archival $HST$ data in order to characterize  the molecular environment, the onset of starbursts, and the   dynamical evolution  in the  circumnuclear  region of barred galaxies. Our sample consists of ten  carefully selected  nearby barred non-starbursts  and starbursts, including some of the  most luminous  starbursts within 40 Mpc. Given the time-consuming nature of   interferometric $2 \\arcsec$ CO observations, this is one of the largest  sample studied at this resolution. CO observations of  nearby galaxies have in the past often been limited to a a few individual systems,  and it is only in recent years  that we have seen a systematic mapping of sizable  samples  (10--12) of   nearby galaxies  (e.g.,  Baker 2000;  Jogee 1999). Complementary surveys such as the $7''$ resolution  BIMA CO(1--0) Survey of  Nearby Galaxies  (SONG; Regan  et al. 2001; Helfer et al. 2001) or the NMA--OVRO  $4''$ resolution CO(1--0)  survey of galaxies  (Sakamoto et al.  1999) are better suited  for studying extended structures or the average properties within the inner kpc.     The main sections of the paper are as follows. $\\S$ 2 outlines  the sample selection and $\\S$ 3  the properties of stellar bars in the sample galaxies. The imaging  and  interferometric CO  observations are covered in $\\S$ 4. %In $\\S$ 5, we compute the molecular gas content in the % circumnuclear  region of the barred starburst and non-starbursts. $\\S$ 6 highlights the extreme  molecular environment which has built up in the inner kpc of barred galaxies and discusses its implication. In $\\S$ 7, we estimate the SFR in the circumnuclear region using different  tracers  such as  Br$\\gamma$, RC, and FIR luminosities In $\\S$ 8, we compare the  circumnuclear molecular gas in the barred starbursts and non-starbursts in order to investigate  why they have such different SFR/$M_{\\tiny \\rm H2}$ in the inner few kpc. In $\\S$ 9, we push these investigations further by comparing the observed circumnuclear gas surface density  to the critical density for the  onset of gravitational  instabilities. In $\\S$ 10, we investigate where the molecular gas has piled up with respect to  the  dynamical  resonances of the bar. In $\\S$ 11, we  discuss the results  within the context of bar-driven  secular evolutionary  scenarios. For readers interested in specific galaxies, $\\S$ 12 describes the molecular gas distribution and kinematics of individual galaxies. $\\S$ 13 summarizes our main results. ",
        "conclusions": "There is compelling  observational and theoretical evidence that  bars efficiently redistribute  angular momentum in galaxies and drive gas inflows into the circumnuclear (inner 1--2 kpc) region. However, only few  high resolution studies, based on a large sample of galaxies, have been carried out on the fate of gas in this region. In this study, we  characterize  the molecular environment, the onset of starbursts, and the secular evolution  in the  circumnuclear  region of barred galaxies. We use a sample of   local ($D<$ 40 Mpc) barred  non-starbursts  and starbursts having   high resolution ($\\sim$~200 pc) CO  ($J$=1$\\rightarrow$0), optical, NIR, H$\\alpha$, RC,  Br$\\gamma$, and archival $HST$ observations. Our results are summarized below.     \\indent \\bf (1) \\rm %%% xxx galaxies. The circumnuclear  regions  of barred galaxies  host $3 \\times  10^{8}$ to $ 2 \\times 10^{9}$ M$_{\\tiny \\sun}$ of molecular gas and have developed  a molecular environment that \\it differs markedly \\rm from that  in the outer disk of galaxies. It includes molecular gas  surface densities of 500-3500  M$_{\\tiny \\sun}$ pc$^{-2}$, gas mass fractions of  10 to 30 \\%,  epicyclic frequencies of several 100 to several 1000  km s$^{-1}$  kpc$^{-1}$, and velocity dispersions of 10 to 40  km s$^{-1}$. In this environment, gravitational instabilities set in only at very high gas densities (few 100-1000 M$_{\\tiny \\sun}$ pc$^{-2}$), but  once triggered, they grow rapidly on a timescale of a few Myrs. This    high density,  short timescale, `burst ' mode may explain why the most intense starbursts tend to be in the central parts of galaxies. Furthermore, the  high pressure, high turbulence ISM can lead to the formation of  clouds with high internal dispersion and mass, and hence may favor the formation of massive clusters as suggested by Elmegreen et al (1993). The  molecular environment  in the  inner kpc of  the ULIRG galaxy  Arp 220  is a scaled-up version of the one in these barred galaxies, suggesting  that interactions build up even more extreme conditions.  \\\\ \\indent \\bf (2) \\rm We suggest that the wide variety in CO morphologies is  due to different stages  of bar-driven inflow. A non-starburst like NGC 4569   which is  in the \\it early \\rm stages of bar-driven inflow has a highly extended molecular gas distribution where a large fraction of the circumnuclear gas is still along the large-scale stellar bar, outside the outer inner Lindblad resonance (OILR). This gas   shows large non-circular kinematics and  is not forming stars efficiently. Several other galaxies  studied by others such as NGC 7479 (Laine et al. 1999), NGC 7723 (Chevalier \\& Furenlid 1978), NGC 1300,  and NGC  5383 (Tubbs 1982) may be in a similar phase. In contrast, we present   dynamical and morphological that the other non-starbursts and starbursts are in \\it later \\rm stages of  bar-driven inflow.  Most of their circumnuclear  gas is  inside  or near the OILR  of the bar and has predominantly circular motions.  Across the sample, we estimate upper limits in the range 43 to 115  km s$^{-1}$  kpc$^{-1}$ for the bar pattern speed and an OILR radius of  $>$ 500 pc.\\\\ \\indent \\bf (3) \\rm The barred  starbursts and  non-starbursts have   circumnuclear SFRs of  3 to 11   and  0.1-2 M$_{\\tiny \\sun}$ yr$^{-1}$, respectively. For a given amount of molecular hydrogen ($M_{\\tiny \\rm H2}$) in the inner 1--2 kpc (assuming a standard  CO-to-H$_{\\rm 2}$ conversion factor), barred galaxies can show an order of magnitude variation in the SFR/$M_{\\tiny \\rm H2}$ over this region. This range  seems related to the fact that the  gas surface densities in the starbursts are larger (1000--3500 $M_{\\tiny \\sun}$ pc$^{-2}$)  and close to the  Toomre  critical density over  a large region. The Toomre $Q$ parameter  reaches its minimum value of $\\sim$  1--2 in the region of star formation, despite an order of magnitude  variation in the gas surface density  and epicyclic frequency. This   suggests that the onset of gravitational instabilities, as characterized by $Q$,  plays an important role even in the inner kpc region.\\\\ \\indent \\bf (4) \\rm The  dynamical mass  enclosed  within  the inner kpc radius of the barred galaxies in our sample  is 6--30  $ \\times 10^{9}$ $M_{\\tiny \\sun}$. Molecular gas makes up 10\\%--30% of the dynamical mass  and in the circumuclear starbursts, it is fueling a SFR of 3--11  M$_{\\tiny \\sun}$ yr$^{-1}$ in the inner kpc. As these starbursts use up their gas and evolve into the post-starburst phase, they  will  build  young, massive, high $V/\\sigma$  stellar components within the inner kpc, inside the OILR of the large-scale bar. \\it Such compact stellar components will likely belong to the  class of pseudo-bulges \\rm (Kormendy 1993) whose light distribution and kinematics are more consistent with a disk than with a spheroidal bulge component. We present evidence of such a component in  NGC 3351, which seems to be in a post-burst phase. The observations are consistent with the idea that over a galaxy's lifetime, it can experience numerous episodes of bar/tidally driven gas inflows, which lead to  a  gradual buildup of its  central mass concentration, the formation of pseudo-bulges, and perhaps even secular  evolution along the Hubble sequence."
    },
    "0402/hep-th0402218_arXiv.txt": {
        "abstract": "We analyze the cosmological consequences of {\\it locked inflation\\/}, a model recently proposed by Dvali and Kachru that can produce significant amounts of inflation without requiring slow-roll. We pay particular attention to the end of inflation in this model, showing that a secondary phase of {\\it saddle inflation\\/} can follow the locked inflationary era.  However, this subsequent period of inflation results in a strongly scale dependent spectrum that can lead to massive black hole formation in the primordial universe. Avoiding this disastrous outcome puts strong constraints on the parameter space open to models of locked inflation. ",
        "introduction": "The inflationary paradigm~\\cite{inf} provides a compelling account of early universe cosmology. The universe emerges from the inflationary phase with large-scale homogeneity and endowed with a nearly scale invariant spectrum of density fluctuations, consistent with current observations. Despite these impressive phenomenological achievements, designing successful models of inflation within supergravity and string theory has proven to be a frustratingly difficult task~\\cite{lyth}. In spontaneously broken supergravities lifted flat directions are natural inflaton candidates. However, the various moduli have (stable, protected) masses $m$ of order $H$, the Hubble constant during inflation~\\cite{note,nima}. This perversely spoils slow-roll inflation since the slow-roll condition, $\\eta \\sim m^2/H^2\\ll 1$, is then violated.  In other words, the generic outcome in supergravity theories is $\\eta\\;\\gsim\\; {\\cal O}(1)$. This so-called $\\eta$-problem is encountered, for instance, in attempts to embed inflation in the stringy landscape~\\cite{mald}. Recent developments, however, indicate that a simple change to the K\\\"ahler potential might alleviate this problem~\\cite{newstuff}.  In a recent paper, Dvali and Kachru \\cite{Dvali:2003vv} (henceforth, DK) introduced {\\em locked inflation\\/} as a possible way out of this dilemma. Its distinguishing feature is that it does away with the slow-roll constraints. Instead locked inflation relies on the rapid oscillations of one scalar field which lock a second field at the top of a saddle point. The potential energy at the saddle then drives inflation. At the very least, this model is an intriguing alternative to slow-roll inflation and, if consistent, overcomes the hurdles faced by slow-roll inflation in supergravity and string theories. Locked inflation needs no intrinsically small parameters, although it does exploit -- like most two-field models -- the ratio of the two widely separated scales. The existence of widely separated scales is not a new tuning, however, as this hierarchy must be explained even in the absence of inflation.  In this paper, we examine the termination of locked inflation and the subsequent evolution of the universe. It is perhaps natural to assume that inflation ends as soon as the field-point is no longer trapped at the saddle point.  However, for the parameter values natural in broken supergravities, we find that this is not necessarily the case. Instead, the universe can undergo a second period of inflation as the field point moves orthogonally to the direction about which it was previously oscillating. We dub this phase {\\it saddle inflation\\/}.  Saddle inflation has potentially disastrous observational consequences for the DK scenario.  As we will show, modes that leave the horizon at the onset of saddle inflation typically have an amplitude of order unity and thus give rise to a phenomenologically dangerous number of massive black holes when they re-enter the horizon during the subsequent radiation or matter-dominated eras. The formation of these black holes must thus be avoided at all costs. There are two ways to do so. One is to demand that there be no saddle inflation at all. This requires that the (tachyonic) mass of the saddle field be much larger than $H$, or $\\eta\\gg 1$. The other way out of the black hole problem is to make the secondary phase of saddle inflation last long enough to move the dangerous range of scales outside the present cosmological horizon. This renders the prior period of locked inflation unobservable, although the latter remains useful as a mechanism for resolving the initial conditions problem related to the onset of inflation. Long saddle inflation naturally occurs for $\\eta \\ll 1$.  The generation of large perturbations at a saddle point of the potential has been discussed for general two-field models by Garcia-Bellido, Linde and Wands \\cite{Wands:1996}, and in the specific context of supernatural inflation by Randall, Soljacic and Guth \\cite{Randall:1995}. Adopting the analyses of these earlier models to the specific case of locked inflation, we can readily deduce the phenomenological consequences associated with the end of locked inflation.  We conclude that a viable model of locked inflation requires either $\\eta\\;\\gsim\\; 30-1500$ (no saddle inflation), depending on the reheating temperature, or $\\eta\\; \\lsim \\; 0.01-2.5$ (long saddle inflation), similarly depending on the reheating temperature. Thus, at a naive level, it appears that both models necessitate a similar degree of tuning.  It is noteworthy that long saddle inflation with $\\eta\\sim {\\cal O}(1)$ is allowed for reheating temperature in the range $1-10^9$~GeV. Locked inflation only sets up the desired initial condition for saddle inflation in this case. Nevertheless, this is a successful inflationary model with supergravity-inspired potential. As such, it is a clear candidate for embedding inflation in string theory and supergravity. To reproduce the observed scale invariant spectrum of density perturbations, however, it does require that the fluctuations arise from an alternative mechanism. ",
        "conclusions": "We have shown that, in order for locked inflation to be phenomenologically viable, it must either: i) end without any subsequent saddle inflation; or ii) be followed by a long phase of saddle inflation. This constrains the $(\\eta,M)$ parameter space as illustrated in Fig.~\\ref{cons2}. In the former case, no black holes are formed. In the latter case,  saddle inflation lasts sufficiently long to push the dangerous modes outside our observable universe. However, locked inflation would only serve the purpose of setting up the initial conditions for saddle inflation in this case, but would have no directly testable observational consequences in itself.  In case i), shown as ``No Saddle Inflation'' in the Figure, $\\eta$ and $M$ are required to satisfy (see Eq.~(\\ref{Nsaddle2})) \\begin{eqnarray} \\frac{f^2(\\eta)(f(\\eta)+3)^2}{6(2f(\\eta)+3)^2} \\exp[2f(\\eta)] & \\gsim & \\frac{M_{Pl}^4}{M^4} \\nonumber \\\\ \\stackrel{\\eta \\gg 1}{\\Longrightarrow}~~~~~~~~~ \\frac{\\eta}{8} \\exp[2\\sqrt{3\\eta}] & \\gsim &\\frac{M_{Pl}^4}{M^4}\\,, \\label{eq:1} \\end{eqnarray} which implies $\\eta \\; \\gsim \\; 30-1500$ for $M \\approx 10^{16}\\;{\\rm GeV}-1$~TeV.  In case ii), shown as ``Long Saddle Inflation'' in the Figure, the bound on $\\eta$ and $M$ is (see Eq.~(\\ref{conddgz})) \\begin{eqnarray} \\label{eq:2} \\nonumber \\frac{\\sqrt{\\eta}}{f(\\eta)}\\left(\\frac{M}{M_{Pl}}\\right)^2\\left(\\frac{M}{T_0}\\right)^{f(\\eta)} &<& 10^{-5} \\\\ \\stackrel{\\eta \\ll 1}{\\Longrightarrow}~~~~~~~~~ \\frac{1}{\\sqrt{\\eta}}\\left(\\frac{M}{T_0}\\right)^{\\eta} &<& 10^{-5}\\left(\\frac{M_{Pl}}{M}\\right)^2\\,, \\end{eqnarray} or $\\eta \\; \\lsim\\; 0.01-2.5$ for $M \\approx 10^{15}\\;{\\rm GeV}-1$ TeV.  The general expectation in spontaneously broken supergravity is that $\\eta$ should be of order unity. This is indeed what has been found so far in attempts to embed inflation within string theory~\\cite{mald}. Thus, the above conditions on $\\eta$ amount to a non-trivial tuning that is required for phenomenological viable scenarios of locked inflation with or without a subsequent phase of saddle inflation.  Case ii) does allow for $\\eta\\sim {\\cal O}(1)$ if $M$ is of order TeV scale. This is encouraging for attempts to embed inflation in string theory. It is important to keep in mind, however, that this assumes that density perturbations with the correct amplitude can be generated via the DGZ mechanism, even for such a low reheating temperature. For larger values of $M$, the model is forced towards $\\eta\\ll 1$, corresponding to the slow-roll inflationary regime.  \\bigskip  We conclude with a comment on observational consequences of locked inflation. Here i) is the interesting case, since for case ii) the secondary ``Long Saddle Inflation'' phase erases all observational characteristics of the locked period. Locked inflation without subsequent saddle inflation either solves the standard cosmological problems in a single step with a low reheating scale or in a series of steps with a higher inflation scale, when $N_{locked}$ (Eq.~(\\ref{Nlocked})) for each phase of inflation is too small on its own. In either case one still has to satisfy the constraint on $\\eta$ in Eq.~(\\ref{eq:1}) {\\em at each step.} In multistage locked inflation this constraint does become weaker as $M$ increases, as can be seen from Fig.~\\ref{cons2}. But the cost of this is that the constraint must be satisfied at the end of each of the multiple phases of locked inflation. Instead of needing to solve a stringent constraint once as is the case with locked inflation at low mass scale, one must solve a weaker constraint several times in order to construct a viable model.  If there are several phases of locked inflation, this could have immediate observational consequences. In DK's proposal, perturbations produced during locked inflation arise via the previously mentioned DGZ mechanism, whereby the amplitude of the perturbations depends on the coupling of the inflaton field to another field.  In general, this coupling and the resulting spectrum of perturbations will have a different amplitude during each phase of locked inflation. Consequently, the resulting perturbation spectrum could contain discontinuities corresponding to the different phases of locked inflation.  While there is no guarantee that one of these discontinuities will appear in the portion of the spectrum probed by observations, if $M$ is large enough the amount of inflation produced during each phase is small enough to make this unavoidable.  This argument is similar to that considered in~\\cite{Adams:2001vc} for a single field model where the potential contains a number of steps -- putting one  ``feature'' in the perturbation spectrum requires tuning, whereas adding many features makes it likely that at least one will fall in the range of $k$ accessible to cosmological measurements. This topic deserves further study, particularly if the evidence for a running spectral index seen in the first year data of WMAP survives."
    },
    "0402/astro-ph0402243_arXiv.txt": {
        "abstract": "``Diffuse'' gamma rays consist of several components:  truly diffuse emission from the interstellar medium, the extragalactic background, whose origin is not firmly established yet, and the contribution from unresolved and faint Galactic point sources.  One approach to unravel these components is to study the diffuse emission from the interstellar medium, which traces the interactions of high energy particles with interstellar gas and radiation fields. Because of its origin such emission is potentially able to reveal much about the sources and propagation of cosmic rays. The extragalactic background, if reliably determined, can be used in cosmological and blazar studies.  Studying the derived ``average'' spectrum of faint Galactic sources may be able to give a clue to the nature of the emitting objects. ",
        "introduction": "\\label{connection}  The Galactic diffuse \\gray\\ continuum emission, which arises from cosmic-ray proton and electron interactions with gas and interstellar radiation fields, is the dominant feature of the \\gray\\ sky.  This emission in the range 50 keV -- 50 GeV has been systematically studied in the experiments OSSE, COMPTEL, EGRET on the CGRO as well as in earlier experiments, such as SAS 2 and COS B.  A review of CGRO observations was presented by Hunter et al.\\ (1997).  \\begin{figure}[t]% \\vskip 2.65in \\special{psfile=diffuse_f1_small.ps voffset=0 hoffset=-15 vscale=79 hscale=79} \\caption{EGRET all-sky map in continuum \\gray\\ emission for energies $>$100 MeV (A.~W.~Strong, unpublished). \\label{fig:skymap}} \\end{figure}%  The great sensivity and spatial and energy resolution of the EGRET instrument allowed for detailed spatial and spectral analysis  of the diffuse emission (Fig.\\ \\ref{fig:skymap}). Because the  Galaxy is transparent to high energy \\grays, the diffuse \\gray\\ emission is the line-of-sight integral over the emissivity of the interstellar medium. The latter is  essentially the product of the cosmic ray density and the density of the gas or radiation field. The hydrogen distribution (H$_2$, H {\\sc i}, H {\\sc ii}) is derived from radio surveys and an assumed Galactic rotation curve, where the distribution of molecular hydrogen is derived indirectly from CO radio-emission and the assumption that the conversion factor H$_2$/CO is the same for the whole Galaxy.  The Galactic radiation field consists of contributions of stars, dust, and cosmic microwave background (CMB). Its spectrum varies  over the Galaxy and (apart from the CMB) cannot be measured directly.  The first detailed analysis of the diffuse emission from the plane $|b|\\leq10^\\circ$ was made by Hunter et al.\\ (1997).  The basic assumptions of this calculation were that  (i) the cosmic rays are Galactic in origin, (ii) a correlation exists between the cosmic ray density and interstellar matter in the Galaxy, and (iii) that the spectra of nucleons and electrons in the Galaxy are the same as observed in the solar vicinity.  This analysis confirmed results of earlier experiments \\citep{knifen73,fichtel75,mayer82} that the great majority of the emission is clearly correlated with  the \\emph{expected} Galactic diffuse emission.  It was also shown \\citep{strong88} that, on average, there is a generally decreasing \\gray\\ emissivity per H atom, and hence a decreasing cosmic ray density, with Galactic radius.  The observations have confirmed main features of the Galactic model derived from cosmic rays, however, they brought also new puzzles.  The \\grays\\ revealed that the cosmic ray source distribution required to match the \\gray\\ data apparently should be distinctly flatter \\citep{strong96} than the (poorly) known distribution of supernova remnants (SNRs), the conventional sources of cosmic rays. The spectrum of \\grays\\ calculated under the assumption that the proton and electron spectra in the Galaxy resemble those measured locally reveals an excess at $>$1 GeV in the EGRET spectrum (Fig.\\ \\ref{fig:hunter97}).  \\begin{figure}[t]% \\vskip 2.05in \\special{psfile=diffuse_f2_small.ps voffset=-200 hoffset=-10 vscale=70 hscale=90} \\narrowcaption{Average diffuse gamma-ray spectrum of the inner Galaxy region, $300^\\circ<l<60^\\circ, |b|\\le10^\\circ$. The contribution from point sources detected with more than $5\\sigma$ significance have been removed. The individual components of this calculation are nucleon-nucleon (NN), electron bremsstrahlung (EB), inverse Compton (IC), and isotropic diffuse emission (ID). Adapted from Hunter et al.\\ (1997). \\label{fig:hunter97}} \\end{figure}%  The puzzle of the ``GeV excess'' has lead to an attempt to re-evaluate the reaction of $\\pi^0$-production in $pp$-interactions. However, a calculation \\citep{mori97} made using modern Monte Carlo event generators to simulate high-energy $pp$-collisions has shown that the \\gray\\ flux agrees rather well with previous calculations.  Leaving the possibility of a instrumental artefact aside, another leading reason for the discrepancy discussed is that the local cosmic ray particle spectra (nucleons and/or electrons) may be not representative of the Galactic average.  The local source(s) and propagation effects (e.g., electron energy losses) can change the spectrum of accelerated particles.  A flatter Galactic nucleon spectrum has been suggested as a possible solution to the ``GeV excess'' problem \\citep{mori97,gralewicz97}. Explaining the excess requires the power-law index of proton spectrum of about --2.4--2.5.  A flatter electron spectrum has been proposed by Porter and Protheroe (1997) and Pohl and Esposito (1998).  The average interstellar electron spectrum can be harder than that locally observed due to the spatially inhomogeneous source distribution and energy losses.  The \\gray\\ excess in this case may be explained in terms of inverse Compton emission.  However, the average energy spectrum of the diffuse \\gray\\ emission alone does not tell much about the underlying processes.  Instead, the spectrum of diffuse continuum \\gray\\ emission  from different directions and its distribution on the sky carry unique information about the particle fluxes, mostly protons and electrons, in different locations.   The $\\pi^0$-decay and bremsstrahlung photons are gas related and thus should mimic the distribution of interstellar matter. In contrast, inverse Compton emission is broad since the density of background photons is high at even large distances from the Galactic plane.  In practice, however, this simple picture is complicated because the H\\,{\\sc ii} gas distribution is broad with typical scale height $\\sim$1 kpc, while the distribution of high energy electrons is narrow and concentrated near the Galactic plane due to the large energy losses.  To decode the wealth of information provided to us by diffuse \\grays\\ from different directions one needs a proper propagation model to calculate the particle spectra and the corresponding \\gray\\ flux on a large Galactic scale.  A self-consistent model of particle propagation and generation of diffuse \\grays\\ should include cosmic-ray transport as the first step. Knowing the number density of primary nuclei from satellite and balloon observations, the production cross sections from laboratory experiments, and the gas distribution from astronomical observations, one can calculate the production rate of secondary nuclei. The observed abundance of radioactive isotopes determines then the value of the diffusion coefficient, halo size and other global parameters.  The detailed procedure was described, e.g., by Ptuskin and Soutoul (1998) and Strong and  Moskalenko (1998). Having fixed the propagation model  and assuming some particle spectra in the cosmic-ray sources, this allows one to calculate the spectrum of the diffuse  \\gray\\ emission \\citep{strong00}. ",
        "conclusions": ""
    },
    "0402/astro-ph0402596_arXiv.txt": {
        "abstract": "The recent claim of a high degree of linear polarization in the prompt gamma-ray burst (GRB) emission of GRB 021206 might have important implications for the underlying mechanism ultimately responsible for the GRB radiation. While the claim itself remains controversial, a full characterization of the GRB polarization has become a scientific imperative. A review of past and present polarimetry missions motivates a set of guidelines for future dedicated GRB polarization experiments. It is also argued that polarization in GRBs could be measured by a relatively simple instrument using readily available technology. ",
        "introduction": "The nature of the prompt GRB emission remains one of the outstanding astrophysical mysteries of the past 35 years \\cite{kle}. The central problem concerns the lack of a proven mechanism that can both extract energy from a collapsing GRB progenitor, and generate prime conditions for the production of relativistic outflows \\cite{meszaros}. Three models have emerged as the leading candidates for explanation of the prompt GRB emission. In conventional hydrodynamic models, internal or external shock fronts accelerate particles that radiate high-energy photons \\cite{rees,meszaros-93,dermer,piran}. In the case of electromagnetic models, the GRB progenitor loses much of its spin in the form of an  electromagnetically dominated outflow that can extend all the way out to the \\g-ray emission region \\cite{lyutikov}. The cannonball model relies on inverse Compton scattering of photons to GRB energies \\cite{dar}. Although various numerical simulations have been  successful in reproducing the complex structure observed in GRB light curves \\cite{koba1997,liang2003}, the observations needed to validate any of these models are still sparse.  Perhaps the greatest observational challenge for prompt emission models has been put forward by the recent claim of polarization at the $\\Pi$ = 80 $\\pm$ 20\\% level in observations of GRB 021206 by $RHESSI$ (Ref.~\\cite{coburn2003}, CB03 hereafter). While the interpretation of the data remains controversial (Ref.~\\cite{rut2003}, RF03 hereafter, and Ref.~\\cite{boggs03}, BC03), it is clear that a significant measurement of polarization in GRBs might provide fundamental insight into the GRB mechanism. For example, electromagnetic models could be favored if $\\Pi \\approx$ 30--40\\% is established \\cite{lyutikov,granot}, since hydrodynamical models have difficulties generating more than 10\\% fractional polarization (but see Ref.~\\cite{nakar}). Only the cannonball model with some fine tuning may accommodate fractional polarization as large as $\\Pi \\approx$ 80\\%. Any measurement of $\\Pi <$ 20\\% would not exclude any of the three models since for each of them depolarization effects can enter into play. Nonetheless, an upper limit in polarization would help placing stringent boundaries for the ongoing theoretical effort.  The exciting prospects of polarization measurements in the prompt GRB emission is moving research into new directions \\cite{mcconnell2}. In particular, consensus is growing that a dedicated  polarization experiment is an inevitable step in constraining the GRB mechanism. In this work, after a general introduction (Sec.~2), we discuss the general guidelines for any dedicated GRB polarization experiments in the future (Sec.~3); we also argue that the technology is readily available to develop a new generation of instruments that can achieve improved sensitivity and sky coverage (Sec.~4); lastly, we outline the basic design for a relatively simple experiment that meets the conditions for a significant polarization measurement (Sec.~5). ",
        "conclusions": "The next generation of proposed instruments with the capability of measuring polarization in soft \\g-rays, including ACT \\cite{kurfess,ACT:proposal:2003}, GRAPE \\cite{mcconnell2}, and the design discussed here, promise important constraints of the prompt GRB emission mechanism. While the ACT and GRAPE proposals have a much broader scope than just measuring GRB polarization, the concept put forward in this work is the one of an instrument which will just measure GRB polarization. Even an upper limit on polarization will impose stringent limits for different GRB models. Polarization measurements might also provide a novel way to explore the mechanism responsible for the hard, short class of GRB events since the instrumental resolution is not limited to long events. In particular, it is important to explore if there is more than one class of GRBs in terms of polarization. To assess the potential of a future experiment requires a full-fledged Monte Carlo simulations including a detector mass model, GRB flux distribution and a proper background model. This is of vital importance when evaluating the need for rotation of the detector and modeling the sensitivity of large area experiments. The estimates given in this work are based on simple experimental arguments, scaling the performance of proven instruments like $RHESSI$ and BATSE, and should be useful in setting the initial guidelines for new experimental approaches to GRB polarization measurements. \\\\ Summarizing, the minimal requirements in our analysis are: \\begin{itemize} \\item[1. ] an effective area of 2000 cm$^2$ for ``triggering'' GRBs; \\item[2. ] enough sensitivity to a large fractional polarization for GRBs with peak flux of $\\sim$10 ph cm$^{-2}$ s$^{-1}$. \\end{itemize}  There is at least one last point that has been overlooked in most of the current polarization literature including this work, which is the possibility of energy measurements. Given a fully contained double event, Compton kinematics must be obeyed and this constraint should improve background rejection. In principle, it might even be possible to apply Compton imaging techniques \\cite{scho} to image the GRB, as would be the case (at a very refined level) for the proposed ACT. Lastly, we have pointed out that the technology to take the next step in polarimetry experiments is readily available and a working instrument may be built in relatively short time. The simplicity of the proposed instrument derives from the fact that its goal is narrowly focused on a  single, well-defined measurement without stringent requirements on energy resolution, source imaging capability and low background level. Alas, a simple instrument does not imply a simple experiment."
    },
    "0402/astro-ph0402075_arXiv.txt": {
        "abstract": "We examine the periodic nature of detailed structure (particularly dips) in the RXTE/ASM lightcurve of Circinus X-1. The significant phase wandering of the X-ray maxima suggests their identification with the response on a viscous timescale of the accretion disk to perturbation. We find that the X-ray dips provide a more accurate system clock than the maxima, and thus use these as indicators of the times of periastron passage. We fit a quadratic ephemeris to these dips, and find its predictive power for the X-ray lightcurve to be superior to ephemerides based on the radio flares and the full archival X-ray lightcurve. Under the hypothesis that the dips are tracers of the mass transfer rate from the donor, we use their occurrence rate as a function of orbital phase to explore the (as yet unconstrained) nature of the donor. The high $\\dot{P}$ term in the ephemeris provides another piece of evidence that Cir X-1 is in a state of dynamical evolution, and thus is a very young post-supernova system. We further suggest that the radio ``synchrotron nebula'' immediately surrounding Cir X-1 is in fact the remnant of the event that created the compact object, and discuss briefly the evidence for and against such an interpretation. ",
        "introduction": "\\label{}  Cir X-1 is a galactic X-ray binary located in the galactic plane at a distance of 6-8 kpc (Goss \\& Mebold 1977). Its most well-known feature is strong X-ray variability on a (presumed orbital) cycle of 16.6 days (Kaluzienski et al 1976). The source distance and X-ray lightcurve imply variation of X ray luminosity $L_x$ of a factor $\\sim$100 throughout this cycle (e.g. Kaluzienski et al 1976). Radio flares were found at peak levels of 1Jy, on a similar period to the X-ray variations (Haynes et al, 1978), and a similar cycle is evident in the IR activity (Glass 1994). An arcsec-scale asymmetric jet has been imaged (Fender et al, 1998), raising the possibility that outflow from the system is relativistic. During an interval of low X-ray activity, EXOSAT detected eleven X-ray bursts, and their identification with Cir X-1 strongly suggests that the compact object is a neutron star (Tennant et al 1986a,b). Quasi-periodic oscillations (QPO's) have been detected and characterised on two occasions (Tennant 1987, Shirey et al 1998). The earlier EXOSAT observations suggested Atoll-like behaviour on the X-ray colour-colour diagram, whereas RXTE/PCA observations in the late 1990's show Z-source behaviour (Shirey, Bradt \\& Levine 1999), implying an external magnetic field strength $\\ga 10^9G$ (van der Klis 1995). X-ray maxima are often preceded by dips in the X-ray lightcurve. Spectral fits with ASCA, RXTE and BeppoSAX during these dips are consistent with a partial covering model, in which the X-ray emisison consists of a bright component undergoing varying absorption, plus a fainter component not attenuated by absorbing matter (Brandt et al 1996, Shirey, Levine \\& Bradt 1999, Iaria et al 2001).  Very little is known about the mass-losing star in this system, due to a high but uncertain amount of extinction to the source. Early classification of this system as an HMXB was based on an optical identification (Whelan et al. 1977) which was later disproven (Moneti 1992) as the single object originally identified was subsequently resolved into three distinct stars. The best optical spectra to date are likely dominated by the accretion disk, with the donor an insignificant component (Johnston, Fender \\& Wu 1999). The highly uncertain extinction (5 $<$ $A_v$ $<$ 11 mag) prevents the IR colours from constraining the nature of the donor (Glass 1994). Recent high-resolution IR spectra with IRIS2 on the AAT (Clark et al 2003) show strong hydrogen recombination lines and low-excitation metal lines, which resemble the spectrum of a mid-B supergiant. However the variability of these features makes them entirely consistent also with emission from the accretion disk and/or accretion-driven outflow (note however that A0538-66, a Be binary with a highly eccentric 16.6 day orbit, shows extreme variability in spectral sub-class over its orbital cycle as a result of X-ray irradiation and reprocessing; Charles et al 1983).  Indirect methods of determining the mass of the donor are also indeterminate. Modeling the donor mass and eccentricity from the optical spectra (e.g. Johnston, Fender \\& Wu 1999) assumes identification of spectral components with the donor, which is clearly uncertain. The apparent radio association of Cir X-1 with the supernova remnant (SNR) G 321.9-0-3 suggested an extremely high transverse velocity of $\\sim$ 450 km $s^{-1}$ due to an asymmetric supernova kick (Stewart et al 1993). Subsequent exploration of binary parameter space with this kick velocity as a constraint suggested both an extremely high eccentricity ($0.90 < e < 0.94$) and low donor mass ($M_2$ likely below 1 $M_{\\odot}$; Tauris et al 1999). However, a recent HST study (Mignani et al 2002) finds no measurable proper motion, implying an upper limit on the transverse space velocity of $\\sim$ 200 km s$^{-1}$.  This rules out the association with SNR G 321.9-0-3, thereby leaving $M_2$ and $e$ undetermined.  The Chandra X-ray Observatory (CXO) detected a remarkable set of strong X-ray P Cygni profiles through periastron passage, corresponding to highly ionised states of Ne, Mg, Si, S and Fe (Brandt \\& Schulz 2000). The breadth of the lines ($\\sim 200-1900$ km s$^{-1}$) and their time variability suggest identification of these lines with highly ionised outflow from the inner region of the accretion disk (Schulz \\& Brandt 2002). Taken with the radio observations of an expanding jet from the source (Stewart et al 1993, Fender 2003 priv. comm), the emerging picture of Cir X-1 is of an extremely unusual neutron star analogue to the galactic black hole candidate microquasars such as SS433 (see Mirabel \\& Rodriguez 1999 for a review of stellar jet sources).  The 16.6 day cycle has long been fit with a quadratic ephemeris, based on radio flares measured at HARTRAO between 1978 and 1988 (Stewart et al 1991). This ephemeris shows a comparatively high quadratic correction, which implies a remarkably short characteristic timescale for the period evolution $(P/2\\dot{P})$ of $\\sim$ 5600 years. The longterm X-ray lightcurve of Cir X-1 over the three decades since its identification (Margon et al 1971) was recently examined (Saz Parkinson et al. 2003, hereafter SP03), using Fourier techniques to determine the first ephemeris for this system based on X-ray data alone. Period determination with this method is dominated by the X-ray maxima (section 4), whose large intrinsic phase scatter leads to systematic errors in the resulting ephemeris. Furthermore the average mass transfer rate has clearly varied by factors of at least 10 on a timescale of decades (SP03), the likely effect of which on the binary orbit leads us to doubt the validity of fitting a single ephemeris to all X-ray data based on the maxima alone.  In this paper we examine in detail the RXTE/ASM lightcurve of Cir X-1, and find the periastron X-ray dips provide a superior system clock to the X-ray maxima. Possible explanations for the phase wandering of the X-ray maxima are discussed. We use the resulting measurement of the intrinsic phase profile of the dips to extract information about the orbital eccentricity $e$ and discuss the implication of the X-ray results for the age of Cir X-1 and its future determination. In particular, the high $\\dot{P}$ term in the ephemeris suggests the system may be in an early state of post-SN evolution (see also SP03). ",
        "conclusions": "We have examined the periodic nature of the X-ray maxima and dips exhibited in the RXTE/ASM lightcurve of Circinus X-1. We find that the X-ray dips provide a more accurate system clock than the maxima, and thus identify them with the periastron approach of the binary components. The outbursts are interpreted as the viscous timescale response of the disk to perturbations from the varying mass transfer. We attempt to use the X-ray dips to constrain the nature of the donor and the eccentricity of the orbit, and find that, within the limitations of the model used, a wider range of parameters are in principle allowed than has previously been thought, although without further exploration of the effects of nonsynchronous donor rotation, quantitative limits are at this stage unreliable. The high rate of change of the orbital period adds to the growing body of evidence that Cir X-1 is in an extremely early stage of post-supernova evolution. If the system is indeed in an early stage of post-SN evolution, the observed weak external magnetic field must suggest strong screening of the internal magnetic field of the neutron star."
    },
    "0402/astro-ph0402305_arXiv.txt": {
        "abstract": "{ We present a detailed model-atmosphere analysis of  ZNG 4, a UV-bright star in the globular cluster M13. From the analysis of a high resolution ($ R \\approx 45,000 $) spectrum of the object, we derive the atmospheric parameters to be $\\rm T_{eff}  =  8500\\pm 250$ K, log g  $=$  2.5 $\\pm$ 0.5 and $\\rm [Fe/H]= -1.5$. Except for magnesium, chromium and strontium, all other even Z elements are enhanced with titanium and calcium being overabundant by a factor of 0.8 dex. Sodium is enhanced by a factor of 0.2 dex. The luminosity of ZNG 4 and its position in the color-magnitude diagram of the cluster indicate that it is a Supra Horizontal Branch (SHB) (post-HB) star. The underabundance of He and overabundances of Ca, Ti, Sc and Ba in the photosphere of ZNG 4 indicate that diffusion and radiative levitation of elements may be in operation in M 13 post-HB stars even at $\\rm T_{eff}$ of 8500K. Detailed and more accurate abundance analysis of post-HB stars in several globular clusters is needed to further understand their abundance anomalies. ",
        "introduction": "The term \"UV-bright stars\" was introduced by Zinn et al.(1972) for stars in globular clusters that lie above the horizontal branch (HB) and are bluer than red giants. The name resulted from the fact that, in the U band, these stars were brighter than all other cluster stars. Further investigations showed that this group of stars consist of blue horizontal branch (BHB) stars, supra horizontal branch stars (SHB), post asymptotic giant branch stars (post-AGB), post-early AGB (P-EAGB) stars and AGB-manque stars. (de Boer 1985, 1987, Sweigart et al.1974, Brocato et al. 1990, Dorman et al. 1993 and Gonzalez \\& Wallerstein 1994)  To derive the chemical composition of UV-bright stars in globular clusters and to understand their evolutionary stages, we started a program to obtain high resolution spectra of these objects in selected globular clusters with the High Dispersion Spectrograph (HDS, Noguchi et al. 2002) of the 8.2m Subaru Telescope. We selected a few UV-bright stars in the globular cluster M13 from the papers of Zinn et al.(1972) and Harris et al.(1983) to derive their chemical composition. In this paper we report the analysis of a high resolution spectrum of the UV-bright star ZNG 4 (RA ($16^{h}41^{m}37\\fs528$) and DEC ($+36\\degr30\\arcmin43.86\\arcsec $  ) (2000) ) (Zinn et al. 1972) in M13 as the first target of our program.  M13 (NGC 6205) is a nearby well-studied globular cluster with a distance modulus of \\mbox{$\\rm (m-M)_0$} = 14\\fm42 and metallicity of $\\rm [Fe/H] = -1.51$ (Kraft and Ivans 2003). The position of ZNG 4 in the color-magnitude diagram of M13 (Paltrinieri et al. 1998) is shown in Figure 1. Many of the globular clusters show a prominent gap in the blue tail of the HB, which is presumed to be due to differential mass loss on the Red Giant Branch (RGB). In M13 it is observed at $\\rm T_{eff} = 10000 K$ (Ferraro et al. 1997). High resolution spectroscopic studies of M13 BHB stars lying on either side of the gap were carried out by Peterson et al.(1983, 1995) and Behr et al.(1999, 2000a). They found anomalous photospheric abundances in BHB stars. These photospheric anomalies are most likely due to diffusion - the gravitational settling of helium and radiative levitation of the metal atoms in the stable atmosphere of hot stars. They found variations in the photospheric abundances and rotational velocities of BHB stars as a function of their effective temperatures.  \\begin{figure} \\centering \\includegraphics[width=8.5cm]{star2b.eps} \\caption{ Color magnitude diagram (CMD) of globular cluster M13 obtained by Paltrinieri et al.(1998). The arrow indicates the position of ZNG 4 in the CMD.} \\end{figure} ",
        "conclusions": "The chemical composition of ZNG 4 shows significant deviations in element abundances from the expected metallicity of M13. Ti, Ca, Sc and Ba are found to be relatively overabundant (Figure 5) compared to RGB stars and cool BHB stars of M13, whereas the abundances of Mg, Cr and Fe are in agreement with the cluster metallicity.  From a study of 22 M13 G-K giants, Kraft et al. (1993, 1997) found the abundances of Fe, Sc, V and Ni to be [Fe/H] = -1.49, [Sc/Fe] = -0.10, [V/Fe] = 0.00 and [Ni/Fe] = -0.04. They found Ca and Ti to be mildly overabundant ([Ca/Fe] = +0.24 and [Ti/Fe]= +0.29). Si was overabundant by +0.34 dex. Study of M13 giants by Armosky et al. (1994) yields the average abundance of Fe, Y and Ba : [Fe/H] = - 1.49 , [Y/Fe] = -0.12 and [Ba/Fe] = -0.04, whereas in ZNG 4 we find significant overabundance of Ca, Sc, Ti, Y and Ba compared to that found in M13 giants. Also, overabundance of Na [+0.2 dex] and absence of O lines support the anticorrelation of Na and O abundances found in M13 giants (Kraft et al. 1997).  Behr et al.(1999) have studied the BHB stars in M13 on either side of the HB gap. They find the photospheric compositions and stellar rotation rates to vary strongly as a function of $\\rm T_{eff}$ of the stars. Among the cooler stars in their sample, at $\\rm T_{eff}$ of 8500K, the metal abundances are in rough agreement with the canonical cluster metallicity and the hotter stars with $\\rm T_{eff}$ greater than 10000K show a deficiency of He and enhancement of Fe, Ti, Cr by a factor of 300. However, Mg remains at the canonical cluster metallicity. In ZNG 4 also, Mg abundance is in agreement with the M13 metallicity. Abundances similar to that found in BHB stars of M13 were also found in BHB stars of M15 (Behr et al. 2000b) and BHB stars of NGC 6752 (Glaspy et al. 1989, Moehler et al.1999). The abundance anomalies in these BHB stars are most likely due to diffusion - the gravitational settling of helium (Greestein et al. 1967) and radiative levitation of metal atoms (Michaud et al. 1983). Rotational velocities ({\\it v sini}) appear to have a bimodal distribution in cooler BHB stars, whereas the hotter BHB stars with $\\rm T_{eff}$ greater than 10000K are found to be slow rotators.  However, the abundances and rotational velocity of ZNG 4 are contrary to what is expected for stars on the red side of the HB gap at 11,000 K. It is a slow rotator ($v sin i =  7$ km ${sec}^{-1}$).  It shows underabundance of He and overabundances of Ti, Ca, Sc and Ba. In Table 3 and Figure 5, we have compared the abundances of ZNG 4 with the abundances of M13 BHB stars of $\\rm T_{eff}$ 7681 K (J 11) and $\\rm T_{eff}$ 12,750 K  (WF4-3485) (Behr, 2000c) and with the abundances of the RGB star L262 (Cavallo and Nagar, 2000). In Table 3, the last column shows the mean RGB abundances in M13, where elements from Na to Fe are taken from Kraft et al.(1997) and Y and Ba abundances are from Armosky et al. (1994). It is evident from the abundances listed in Table 3 and from Figure 5 that ZNG 4 shows overabundance of metals compared to that of a M13 RGB star and also when compared to that of a M13 BHB star of similar temperature. These results indicate that in ZNG4, diffusion and radiative levitation of elements may be in operation. Slowly rotating HB stars are also seen on the cooler side of HB gap, but abundance anomalies start from 11,000 K (Moehler et al. 1999, Behr et al. 2000b). This implies that ZNG 4 may have the properties of the stars on the blue side of the HB gap although it has $\\rm T_{eff}$ of 8500 K. In this regard, more accurate determination of abundances of these elements in ZNG 4 and similar stars in M13 is needed to confirm our results and conclusions.  This may be explained in two ways. One is that, for some stars in M13, the onset of diffusion seems to start at lower $\\rm T_{eff}$ ($\\approx $ 8500 K). The other argument would be that the star has evolved from the blue side of the HB gap and is moving towards the red with higher luminosity as indicated by the post-HB evolutionary tracks of Gingold (1976) and Dorman et al.(1993).  The BHB stars hotter than 11500K typically show strong photospheric helium depletions due to gravitational settling (Moehler et al. 2000 \\& 2003). The calculations of Michaud et al. (1983) indicate that helium depletion should be accompanied by photospheric enhancement of metals, since the same stable atmosphere that permits gravitational settling also permits the levitation of elements with large radiative cross sections. The depletion of helium and overabudance of some of the metals in the photosphere of ZNG 4 is in qualitative agreement with the calculations of Michaud et al. (1983).  Recently, Turcotte et al. (1998) and Richer et al. (2000) made diffusion simulations to explain the abundance patterns of chemically peculiar A and F stars. Their predicted abundance patterns are qualitatively similar to that found in ZNG 4. However, none of the recent diffusion studies treated the cases of BHB stars and post-HB stars. This phenomenon may be related to the disappearance of surface convection and hence to the formation of a stable stellar atmospheres. HB stars and post-HB stars  cooler than $\\rm T_{eff}$ $=$ 6300K have deep convective envelopes (Sweigart 2002).  Hotter than this temperature the envelope convection breaks into distinct shells associated with the ionization of H and He. Note that the surface convection disappears at 11000K ( Sweigart 2002) and BHB stars hotter than this show moderate to severe abundance anomalies.  ZNG 4 has a V magnitude of 13.964 (Paltrinieri et al. 1998). Considering the distance modulus of M13 to be 14.42 and $\\rm E(B-V)$ towards M13 to be 0.02 (Kraft and Ivans 2003), we estimated the absolute magnitude ($\\rm M_v$) of the star to be $ -0.522$. For stars with $\\rm T_{eff}$ around 8500 K, the bolometric correction (BC) is negligible (Flower 1996). Considering $\\rm BC=0$, we get the bolometric magnitude ($\\rm M_{bol}$) to be $-0.522$, which corresponds to a luminosity of 127 $\\rm L_{\\odot}$ [ $\\rm log {\\frac {L} {L_{\\odot}}} = 2.18$]. Using the equation connecting the mass, effective temperature and bolometric magnitude, we find the surface gravity, log g = 2.6 (assuming the mass of ZNG 4 to be 0.5 $\\rm M_{\\odot}$), which agrees well with the value estimated from the analysis of the spectrum of ZNG 4.  The post-AGB star Barnard 29 [ $\\rm log {\\frac {L} {L_{\\odot}}} = 3.3$] which is a member of M13 is more luminous than ZNG 4. The abundance pattern of ZNG 4 is very different from that of the post-AGB star Barnard 29 (Conlon et al. 1994, Moehler et al. 1998). The M13 BHB stars with $\\rm T_{eff}$ around 8500 K have a luminosity of about 40 $\\rm L_{\\odot}$, whereas ZNG 4 is more luminous by about a factor of 3, which indicates that ZNG 4 can be classified as a supra horizontal branch star (post-HB). Stars that lie 1.5 magnitude above the HB stars are classified as supra horizontal branch (SHB) stars in the photometric studies of M13 (Zinn 1974) and NGC 6522 (Shara et al. 1998). No detailed abundance analysis of SHB stars in globular clusters is available to compare with the abundances of ZNG 4.  Since ZNG 4 is a post-HB star and it has evolved from a hot BHB star stage and may had severe abundance anomalies similar to those found in the hot BHB stars of M13 (Table 3). The present $\\rm T_{eff}$ $=$ 8500K of ZNG 4 indicates that thin layers of subsurface convection if present may have diluted the severe abundance anomalies due to diffusion and radiative levitation that took place during its hot BHB stage of evolution. It is important to derive the chemical composition of a significant sample of post-HB stars hotter than 11000K and much cooler than 11000K to further understand the role of diffusion, radiative levitation, rotation and convection during the post-HB stage of evolution."
    },
    "0402/astro-ph0402133_arXiv.txt": {
        "abstract": "We report the EXITE2 hard X-ray imaging of the sky around 3C273.  A 2h observation on May 8, 1997, shows a $\\sim$260 mCrab source detected at $\\sim$4$\\sigma$ in each of two bands (50-70 and 70-93 keV) and located $\\sim$30' from 3C273 and consistent in position with the SIGMA source GRS1227+025. The EXITE2 spectrum is consistent with a power law with photon index 3 and large low energy absorption, as indicated by the GRANAT/SIGMA results. No source was detected in more sensitive followup EXITE2 observations in 2000 and 2001 with 3$\\sigma$ upper limits of 190 and 65 mCrab, respectively. Comparison with the flux detected by SIGMA shows the source to be highly variable, suggesting it may be non-thermal and beamed and thus the first example of a ``type 2'' (absorbed) Blazar.  Alternatively it might be (an unprecedented) very highly absorbed binary system undergoing accretion disk instability outbursts, possibly either a magnetic CV, or a black hole X-ray nova. ",
        "introduction": "GRS1227+025 was discovered serendipitously with SIGMA \\citep{Bassani91, Jourdain92} in the 40-80 keV energy band during an observation of the QSO 3C273.  Separated from 3C273 by only $\\sim$15', GRS1227+025 (hereafter GRS1227) was detected at 5.5$\\sigma$ whereas 3C273 was not detected.  The source has a significantly steeper spectrum (photon index 3$^{+1.3}_{-0.9}$, \\cite{Bassani91}) compared to 3C273, which was detected in subsequent observations, with its usual flatter photon index of $\\sim$1.5.  Upper limits (15-30 keV) with ART-P (simultaneous with SIGMA) showed that the spectrum must be heavily absorbed: N$_H \\sim$1.5$\\times$10$^{25}$ cm$^{-2}$ for a power law spectrum with photon index 3 \\citep{Bassani93}.  Followup observations done with SIGMA and OSSE (which could only measure the combined flux of 3C273 and GRS1227) showed that GRS1227 is probably variable by a factor 2-5 on $\\sim$1 day time scales \\citep{Bassani93, Johnson95}. No obvious soft X-ray counterpart was found in the SIGMA error circle (5 arcmin, 1$\\sigma$) with ROSAT (e.g. \\cite{Leach96}), although an Einstein source (1E 1227+0224) in the SIGMA error circle (1$\\sigma$) was identified with a QSO at z=0.57 that would have to be of ``type 2'' (heavily self-absorbed) if it were indeed the counterpart \\citep{Grindlay93}. ",
        "conclusions": "The GRS1227 source presents a major puzzle.  Both EXITE2 and SIGMA detect a source at approximately the same location (within 2$\\sigma$) and with similar spectral properties which are very distinct from 3C273. Although the EXITE2 detection is marginal, as was the original SIGMA detection, the combined results are more likely to be real. GRS1227 must be highly self-absorbed and yet highly variable. The SIGMA variations are a factor of $\\sim$2 and SIGMA-EXITE2 variation is nearly a factor of $\\sim$20.  This is a combination not seen before, and suggests a new class of object: a type 2 (i.e. heavily absorbed) Blazar, since Blazars (BL Lacs) are usually the only AGN which show such extreme variability.  However, it might also be a Narrow Line Seyfert 1 (NLSy) galaxy (e.g. RX J2217.9--5941, see \\cite{Grupe01}) but must still be highly self-absorbed.  It may be similar to the NLSys with significant internal absorption and significant variability described by \\cite{Fabian99}. Alternatively it might be (an unprecedented) very highly absorbed binary system, possibly a low luminosity magnetic CV undergoing outburst, since intermediate polars at high inclination can also show significant internal absorption.  From the EXITE2 vs. SIGMA source error circles (2$\\sigma$) shown in Figure~\\ref{grspos}, the combined source error box is $\\sim$15'$\\times$6', centered on the EXITE2-SIGMA combined position RA=12$^h$30$^m$14$^s (\\pm$4'), DEC= +2$^\\circ$14'53''($\\pm$4') (J2000). At this relatively uncrowded high latitude position, even a $\\sim$2' position could enable identification with an optical or radio variable. However the need for a \\la1' position is obvious: if this is, as suspected, a highly absorbed non-thermal source (BL Lac), it may not be at all conspicuous in the optical and appear as a simple elliptical galaxy.  The radio galaxies numbered (2) and (4) in Figure~\\ref{grspos}, or [SRK80]122803+023451 and 87GB[BWE91]1227+0230, are thus possible candidates.  However, in the event that the radio emission is confined to the core, GRS1227 may not even be bright at radio wavelengths since it is likely Compton-thick.  Given the space density of galaxies cataloged (SIMBAD and NED) within the combined EXITE2/SIGMA error circles (cf. Figure~\\ref{grspos}), a \\la1' source position is obviously needed for a unique galaxy identification.   Source variability measurements can partly constrain the underlying source nature. The generalized condition for maximum variability of an isotropically emitting source powered by accretion at rate \\mdot and radiating luminosity L$_x$ is \\\\  $\\Delta$L$_x$  \\la  2 $\\times$ 10$^{41}$ $\\eta_{0.1}$ $\\Delta$t ergs/s, \\\\  \\ni where $\\eta_{0.1}$ = L$_x$/\\mdot c$^2$ \\about 0.1 is the likely maximum efficiency factor (cf. \\cite{Guilbert83}). This can impose direct constraints on the source luminosity (and thus distance and nature) for an observed luminosity variation $\\Delta$L$_x$ over time interval $\\Delta$t. Therefore, a factor of 2 increase on \\about1d timescales (already suspected from SIGMA) would imply the source is beamed and thus a Type 2 Blazar if it were identified with an AGN at z \\ga 0.1 (like 3C273) with L$_x$ \\ga10$^{45}$ erg/s.  GRS1227 exhibits an unusual spectral shape.  The simultaneous detection by SIGMA (\\ga40 keV) and upper limit at \\la15 keV by the ART-P X-ray imager on GRANAT (Bassani 1991) indicated that the source is highly self-absorbed (N${_H} \\sim$10$^{24-25}$ cm$^{-2}$) and thus a Compton-thick absorber (see \\cite{Matt99} and references therein).  If the spectrum above $\\sim$40-50 keV is indeed a power law with photon index $\\sim$3, it is consistent with a Blazar or NLSy. Although NLSys are also highly variable, these have not been reported as Type 2 objects. As mentioned above, the EXITE2 spectral shape (Figure 3) is equally consistent with a flatter spectrum with photon index 1.7 but high energy cutoff at $\\sim$100 keV, as expected for a highly absorbed Seyfert 2 or Type 2 QSO. However the remarkable variability required argues against the Seyfert or Type 2 QSO interpretation and in favor of a Blazar, with highly variable non-thermal emission. In this case, large amplitude optical variability might also be expected, since \\cite{Liller75} found a $\\sim$5 magnitude brightening for the BL Lac PKS 1510-089 on the Harvard plates. Given the offset from 3C273, it is possible that significant optical variables could have been missed, and a search using the Harvard plates is in progress.  A magnetic CV interpretation is also possible even if the spectrum is indeed a (moderately steep) power law: the magnetic CV AE Aqr can be fit with a similar power law spectrum \\citep{Eracleous96}, and has been reported as a high energy gamma-ray source (cf. Meintjes and de Jager 2000 and references therein).  A magnetic CV origin for GRS1227 might also explain the extreme flare variations: the object could be a dwarf nova (like GK Per) undergoing outbursts. Although the X-ray emission in classical dwarf novae outbursts is usually soft, an outburst in a magnetic CV may contain significant non-thermal emission. Historical optical measurements of the could again provide an important test, since the optical flux (from the disk) would be expected to increase by $\\sim$5 mag.  If the EXITE2 1997 detection was the peak of an outburst like that of GK Per, the expected X-ray optical flux ratio F$_x$/F$_v$$\\sim$10 would imply an outburst magnitude V$\\sim$10, and thus quiescence magnitude V$\\sim$15-16, which is probably too bright to have been missed as a blue object in the Palomar Green survey (Green, Schmidt and Liebert 1986), which revealed no such objects near 3C273.  A high latitude black hole X-ray binary, an X-ray nova (XN) like XTE J1118+480. which showed evidence for magnetic flares and non-thermal optical emission (Merloni, DiMatteo and Fabian 2000) is conceivable.  Here the power law hard X-ray spectrum is natural; this is seen in the very high state of BH novae (e.g. Nova Muscae, \\cite{Esin97}) and is expected for the jet emission probable in a system like XTE J1118+480.  It too would have to be highly self-absorbed, but - as with the CV hypothesis - is plausible if the system were at inclination $\\sim$90$^\\circ$, and self-absorbed at the required column density of N$_H\\sim$10$^{25}$ cm$^{-2}$. Such a highly absorbed XN would be missed by the ASM on RXTE, and indeed no evidence for an ASM source is found at the time of the 1997 EXITE2 observation (or the 2000 and 2001 EXITE2 observations; cf. Figure 4). However, an optical outburst is again expected and should again be searched on historical monitoring (e.g. plate) data.  Thus the galactic binary/transient scenarios might be distinguished from the Type 2 BLazar or NLSeyfert interpretation by historical optical or radio variations, which are probably larger in the binary case given the relatively low luminosity secondary star expected. The lack of any (large amplitude) optical variable would support the Type 2 Blazar or NLSeyfert interpretation.  Definitive confirmation of the source existence, and of course a refined hard X-ray position and spectrum, are obviously needed.  Long after this paper was initially submitted, we have carried out an observation with INTEGRAL (IBIS telescope) to attempt to further measure the source position, spectrum and variability. Data have just been received and will be reported separately. An initial INTEGRAL/IBIS observation of 3C273 has been reported recently by Courvoisier et al (2003) which also provided a sensitive new upper limit ($\\sim$10 mCrab) for emission from the position of GRS1227 -- a factor of $\\sim$6 below the SIGMA  detection and 2001 EXITE2 upper limit and $\\sim$25 below the 1997 EXITE2 detection. Thus the source variability at \\ga40 keV must be at least this factor. Wide-field XMM observations of the 3C273 field should also be examined for evidence for a highly variable, highly absorbed (\\la 10 keV) X-ray source which would enable identification."
    },
    "0402/astro-ph0402419_arXiv.txt": {
        "abstract": "We consider the problem of radiative transfer in stellar atmospheres where the index of refraction departs from unity and is a function of density and temperature. We present modified Feautrier and Lambda-iteration methods to solve the equation of radiative transfer with refraction in a plane parallel atmosphere. These methods are general and can be used in any problem with 1-D geometry where the index of refraction is a monotonically varying function of vertical optical depth. We present an application to very cool white dwarf atmospheres where the index of refraction departs significantly from unity. We investigate how ray curvature and total internal reflection affect the limb darkening and the pressure-temperature structure of the atmosphere. Refraction results in a much weakened limb darkening effect. We find that through the constraint of radiative equilibrium, total internal reflection warms the white dwarf atmosphere near the surface ($\\tau \\wig< 1$). This effect may have a significant impact on studies of very cool white dwarf stars. ",
        "introduction": "Problems involving radiative transfer in refractive media appear to have received limited attention \\cite{harris65, zhelez67, pomranin68}.  In stellar atmospheres in particular, refraction is always ignored since the gas is so tenuous that the index of refraction does not depart from unity. A notable exception is found in the atmospheres of very cool white dwarfs, especially those rich in helium, where the gas (or rather the fluid) density can reach $0.1-2\\,$g/cm$^3$. Under these conditions, the refractive index of fluid, atomic helium can become as large as $\\sim 1.3$ with very large gradients (Fig. \\ref{F1}) and refraction can be expected to affect the radiation field. In this contribution, we study radiative transfer in a stellar atmosphere with refraction. While we are motivated primarily by astrophysical applications, the method developed here is general and can be directly applied to other problems with plane parallel geometry with an arbitrary, monotonic variation of the index of refraction. \\begin{figure*}[t] \\centering \\epsscale{0.5} \\plotone{f1.eps} \\caption{Variation of the index of refraction as a function of vertical optical depth $\\tau_{\\nu}$ in the nominal white dwarf atmosphere model for $\\lambda=0.948\\,\\mu$m. \\label{F1}} \\end{figure*}  We consider refraction in a stellar atmosphere that is static, plane parallel, in local thermal equilibrium, in hydrostatic equilibrium, with both radiative and convective energy transport.  The total flux is constant throughout the atmosphere and scattering is assumed isotropic.  We further assume that the refraction of ray paths follows Snell's law (i.e. geometric optics) as is appropriate in white dwarf atmospheres.  For a given run of the index of refraction through the atmosphere, this completely defines the trajectories of the ray paths and allows the derivation of a simple form of the equation of radiative transfer (ERT) in the presence of dispersive effects.  There is a number of interesting physical effects that we can expect in a stellar atmosphere where the index of refraction is greater than unity.  The index $n_\\nu$ is a monotonically increasing function of density. In a stellar atmosphere, the index decreases from the bottom toward the surface where $n_\\nu=1$ at $\\tau_\\nu=0$. All rays are refracted away from the upward vertical and some rays are internally reflected back toward the interior (Fig. \\ref{F2}).  Intuitively, this should lead to an increase in the temperature profile of the atmosphere to achieve flux conservation. Because the largest variations of the index occur in the optically thin regions of the atmosphere (Fig. \\ref{F1}) we expect refraction to be mainly a surface effect. Furthermore, the group velocity of light is reduced by a factor of $n_\\nu$ so light accelerates toward the surface where it reaches the speed of light in a vacuum.  This directly affects the flux through its dependence on the velocity of propagation.  The angular redistribution of  light due to refraction will reduce the limb darkening effect. Finally, the index of refraction of fluid helium is nearly independent of frequency from the optical to the near infrared and there should not be any significant chromatic effects. \\begin{figure*}[t] \\centering \\epsscale{0.8} \\plotone{f2.eps} \\figcaption[P2.eps]{Geometry and typical ray paths for an axially symmetric radiation field in a refractive, planar atmosphere with a monotonically varying refractive index $n_{\\nu}$. \\label{F2}} \\end{figure*}  We follow the treatment of the theory of radiative transfer in dispersive media presented by Cox \\& Giuli (1968). We develop two different numerical schemes to solve the ERT in the presence of refraction based on the Feautrier and $\\Lambda$-iteration methods. Both methods are widely used in stellar atmosphere codes and the procedures we have developed allows for a straightfroward implementation of the effects of refraction. To illustrate the effects of refraction in a stellar atmosphere, we apply both methods to solve for the radiation field in a realistic helium-rich white dwarf atmosphere model with $\\Teff=4000\\,$K, a gravity of $\\log g ({\\rm cm/s}^2)=8$  and a homogeneous composition of $n({\\rm He})/n({\\rm H})=10^6$, where $n({\\rm X})$ is the number density of element X.  This atmospheric structure, obtained with our white dwarf atmosphere code (Bergeron, Saumon \\& Wesemael 1995), is hereafter referred to as the ``nominal\" model.  In the next section, we construct the ERT for a refractive, planar atmosphere, and new expressions for the moments of the radiation field are given in Section 3. The numerical solution with modified Feautrier and $\\Lambda$-iteration methods is described in Section 4. Numerical aspects of the solution, a detailed discussion of the effects of refraction on the radiation field, and the effects on the structure of the atmosphere and the emergent flux are presented in Section 5. ",
        "conclusions": "Having realized that in the dense atmospheres of very cool white dwarf stars -- especially those with a composition dominated by helium -- the index of refraction departs significantly from unity, we have solved the problem of radiative transfer in plane parallel, LTE, static stellar atmospheres in the presence of refraction for arbitrary vertical variations of the opacity, temperature, and index of refraction. We show how to modify the Feautrier solution and the $\\Lambda$-iteration method to solve the equation of radiative transfer in the presence of refraction. The resulting methods are general and can be applied to any one-dimensional problem. For cool white dwarfs, the most important effect of refraction on the propagation of radiation in a dispersive medium is total internal reflection.  This increases the mean intensity in the optically thin region near the surface, where the dispersive effects are strongest. If radiative equilibrium is imposed, this results in higher surface temperatures than in the non-refractive case. We also find that limb darkening is dramatically reduced by refraction. This first solution of the radiative transfer equation with dispersive effects in stellar atmospheres suggest that refraction is an important effect that must be considered when modeling the atmospheres of very cool white dwarfs.  The authors wish to thank D. Mihalas and L.H. Auer for useful discussions. P. Kowalski is grateful to the Los Alamos National Laboratory for its hospitality and support. The research presented here was supported by NSF grant AST97-31438, NASA grant NAG5-8906, and by the United States Department of Energy under contract W-7405-ENG-36.  \\clearpage"
    },
    "0402/astro-ph0402127_arXiv.txt": {
        "abstract": "We report timing and spectral properties of the transient Be X-ray pulsar EXO~053109--6609.2 studied using observations made with the ASCA and Beppo-SAX observatories. Though there must have been at least one spin-down episode of the pulsar since its discovery, the new pulse period measurements show a monotonic spin-up trend since 1996. The pulse profile is found to have marginal energy dependence. There is also evidence for strong luminosity dependence of the pulse profile, a single peaked profile at low luminosity that changes to a double peaked profile at high luminosity. This suggests a change in the accretion pattern at certain luminosity level. The X-ray spectrum is found to consist of a simple power-law with photon index in the range of 0.4--0.8. At high intensity level the spectrum also shows presence of weak iron emission line. ",
        "introduction": "The two Magellanic Clouds have larger number density of High Mass X-ray Binaries (HMXB) compared to our galaxy. The Small Magellanic Cloud (SMC) with a mass of about 1\\% of our Galaxy has a total of about 70 HMXBs and HMXB candidates, almost comparable to that in our galaxy (Haberl \\& Sasaki 2000; Yokogawa et al. 2003). Though the HMXB number density in the Large Magellanic Cloud (LMC) is not as high as that in SMC, it is still significantly higher than the Galactic value (Sasaki, Haberl \\& Pietsch 2000). In-spite of a large number of HMXB pulsar objects in the two Magellanic Clouds, individual objects have not been studied in great detail except for SMC~X-1 and LMC~X-4. There are four HMXB pulsars known in the LMC, of which LMC~X-4 is the best studied. Several of these sources show a very large X-ray luminosity in their high states, close to or exceeding the Eddington limit for a 1 M$_{\\sun}$ object. Sources, in which a strong local absorption column is absent, nature of the soft spectral component can be studied in detail. Due to a low galactic absorption column density towards the LMC and SMC, these objects are also suitable for study of the local absorbing material. A soft excess in the X-ray spectrum, now detected in several accreting pulsars can be a common feature of the HMXB pulsars. But since most of the HMXB pulsars are in the Galactic plane, large line of sight absorption usually makes it difficult to detect it. It is therefore, important to study details of the individual HMXBs in the LMC and SMC.  The transient Be X-ray binary source EXO~053109--6609.2 was discovered with the EXOSAT observatory in 1983 (Pietsch, Rosso \\& Dennerl 1989). Hard X-ray emission from this source was detected with the SL2 XRT experiment (Hanson et al. 1989) and its intensity variations were studied in detail with the ROSAT (Haberl, Dennerl, \\& Pietsch 1995). X-ray pulsations were also discovered with the ROSAT (Dennerl, Haberl \\& Pietsch 1995). This pulsar is located at a distance of 17$\\arcmin$ from the luminous binary X-ray pulsar LMC~X-4, and has a pulse period of 13.7 s, which is also close to the 13.5 s pulse period of LMC~X-4. During a Beppo-SAX observation of LMC~X-4, EXO~053109--6609.2 was in the field of view of the narrow field instruments of Beppo-SAX, and pulsations were detected along with a positive local period derivative (Burderi et al. 1998). The source was also observed with the XMM in 2000 during which it was in a relatively bright state (Haberl, Dennerl, \\& Pietsch 2003). The optical counterpart is one of the two close double Be-stars (Stevens, Coe and Buckley 1999)  In the subsequent sections we present the archival observations that we have used to study this source (\\S 2), results from temporal and spectral analysis of the data (\\S 3) and follow up this with a discussion on the nature of this pulsar (\\S 4). ",
        "conclusions": "The pulsar EXO~053109--6609.2, like many other HMXB pulsars with Be star companions, is believed to be a transient source. However, we have found this pulsar to be bright, at a luminosity level of $3 \\times 10^{35}$ erg s$^{-1}$ to $6 \\times 10^{36}$ erg s$^{-1}$ in most of the X-ray observations made since 1993 with the low background imaging instruments. Several ROSAT observations during March to May 1993 also found the source to have significant X-ray intensity (Haberl et al. 1995). In this respect, this pulsar is more akin to 4U~1907+09 (In't Zand 1999, Mukerjee et al. 2001), which also shows intensity variations of similar scale but without any obvious transient phenomena. Its long term X-ray intensity behaviour is different from that of the Be-Star transients like EXO~2030+375 (Parmar, White \\& Stella 1989). The peak X-ray luminosity of $6 \\times 10^{36}$ erg s$^{-1}$ observed in 1997 is however, much smaller than that of the pulsars SMC~X-1 and LMC~X-4 (Paul et al. 2002), which have high mass companion stars in very close orbits. This indicates a large size orbit and hence orbital period. The archival X-ray observations are not sufficiently long to search for pulse arrival time delay that can help to determine the orbital parameters. Analysing the RXTE-ASM light curve of EXO~053109--6609.2 we have found a peak in the power spectrum at 183 days and its integer multiples. Intensity variations of the bright nearby source LMC~X-4 is unlikely to create this because the 30.5 day superorbital intensity variation of LMC~X-4 is not reflected in the power spectrum of EXO~053109--6609.2. However, since this is very close to a half-year, yearly observational effects cannot be ruled out.  Combining the present observations with the earlier measurements of pulse periods with the ROSAT and the Beppo-SAX, we have obtained a period history of the pulsar, shown in Figure 1. It shows a clear monotonic spin-up trend since the ASCA observation in 1996, though not with a constant period derivative. However, certainly there were spin-down episodes between the 1991 ROSAT observation and the 1996 ASCA observation. A spin-down was also observed during the 1997 Beppo-SAX observation itself, which is a surprise considering the high intensity level during this observation. The last 5 years of pulse period measurements presented here shows a spin-up of the neutron star. This suggests that accretion onto the neutron star was in progress for at least a significant part of this period. However, we cannot rule out a significant binary motion related component in the observed pulse period variations.  The X-ray emission from the polar cap regions can have different geometric pattern at different mass accretion rates. In the transient X-ray pulsar EXO 2030+375, a change in beaming pattern, from a fan beam in high luminosity state to a pencil-beam in low luminosity state was observed with the EXOSAT (Parmar et al. 1989). The pulse profiles of EXO~053109--6609.2 at different luminosity levels presented here and elsewhere (Burderi et al. 1998, Haberl et al. 2003) clearly indicate a dependence on the mass accretion rate. The pulse profile is single peaked and nearly sinusoidal in low intensity state (Beppo-SAX : 1998), and it has two distinct well separated peaks or an admixture of the above two during other observations. It is possible that in EXO~053109--6609.2 also, the X-ray beam pattern changes from a pencil beam at low luminosity to a fan beam at high luminosity. The change in beam pattern probably takes place at a luminosity level of $\\sim10^{36}$ erg s$^{-1}$.  During the ASCA observation in 1996 and the Beppo-SAX observation in 1998, the source was in low intensity state and the spectral parameters are not constrained well from these observations. However, the spectral characteristics reported here from the other three observations indicate intrinsic spectral variability in the pulsar. The power-law component of the spectrum is softer at higher luminosity.  The XMM spectrum of this source unambiguously showed the presence of a soft spectral component in excess of a partially absorbed hard power-law extended to lower energies (Haberl et al. 2003). The same observation also showed that the pulse profile is different in the soft and the hard X-ray bands. This indicates that the soft component may have a different origin, or at least the geometry of emission is different in different energy ranges. Presence of narrow features in the soft excess prompted Haberl et al. (2003) to use a hot plasma model that yields a temperature of 0.1 keV and emission measure of 4 $\\times$ 10$^{61}$ cm$^{-3}$. Such a plasma is unlikely to show pulsations due to its large physical size and large cooling time scale, contrary to what is seen in the 0.15-0.4 keV band. A blackbody type component with separate low energy emission lines as seen in Her X-1 (Endo et al. 2000) may be a better description of the soft excess of this pulsar. Several other accreting X-ray pulsars with relatively low line of sight absorption show soft excess (Paul et al. 2002 and references therein). In some sources the soft excess probably originates in a part of the X-ray irradiated inner accretion disk (Endo et al. 2000; Naik \\& Paul 2004). From independent analysis of the XMM spectra we have found that in EXO~053109--6609.2, the total luminosity of unabsorbed X-ray emission in the soft excess is $1.1 \\times 10^{36}$ erg s$^{-1}$, which is only about 1\\% of the total emission in the 1--10 keV energy band. The ratio of the soft excess flux to the total flux is smaller than that of SMC~X-1 (3.6\\%), LMC~X-4 (6.4\\%) and Her~X-1 (10\\%). In this pulsar, the size of the black body emission region is required to be only about $10^{7}$ cm and, therefore, can easily be accommodated in inner part of the accretion disk heated by the hard X-rays."
    },
    "0402/astro-ph0402311_arXiv.txt": {
        "abstract": "\\large The possibility to avoid the cosmic initial singularity as a consequence of nonlinear effects on the Maxwell eletromagnetic theory is discussed. For a flat FRW geometry we derive the general nonsingular solution supported by a magnetic field plus a cosmic fluid and a nonvanishing vacuum energy density. The nonsingular behavior of solutions with a time-dependent $\\Lambda(t)$-term are also examined. As a general result, it is found that the functional dependence of $\\Lambda(t)$ can uniquely be determined only if the magnetic field remains constant. All these models are examples of bouncing universes which may exhibit an inflationary dynamics driven by the nonlinear corrections of the magnetic field. ",
        "introduction": "\\label{s1}  A fundamental difficulty underlying the standard Friedmann-Robertson-Walker (FRW) cosmology is the prediction of an initial singular state where all curvature invariants and some material quantities like pressure, energy density and temperature become infinite\\cite{Wheeler}. Generically, ``the break down of the laws of physics at a singularity\" is a clear manifestation of mathematical inconsistency and physical incompleteness of any cosmological model. In this way, although strongly supported by the recent observations at low and intermediate redshifts, the present big-bang picture with dark energy must be improved at very high redshifts.  In order to solve such a problem several attempts based on many disparate mechanisms have been proposed in the literature. Some earlier approaches trying to develop a well behaved and more complete cosmological description include: quadratic Lagrangians and other alternatives theories for the gravitational field\\cite{QL}, a creation-field cosmology\\cite{pad}, a huge vacuum energy density at very early times\\cite{VED}, nonminimal couplings\\cite{NMC}, nonequilibrium thermodynamic effects\\cite{NET}, and quantum-gravitational phenomena closely related to a possible spontaneous birth of the universe\\cite{Kandrup}.  More recently, a new interesting mechanism aiming to avoid the cosmic singularity has been discussed by De Lorenci et al.\\cite{novello98a} through a nonlinear extension of the Maxwell eletromagnetic theory. The associated Lagrangian and the resulting electrodynamics can theoretically be justified with basis on different arguments. For example, the nonlinear terms can be added to the standard Maxwell Lagrangian by imposing the existence of symmetries such as parity conservation, gauge invariance, Lorentz invariance, etc\\cite{novello96,Munoz96}, as well as by the introduction of first order quantum corrections to the Maxwell electrodynamics\\cite{heisenberg36,schwinger51}. It is worth notice that nonlinear corrections may also be important to avoid the black hole singularity. Actually, an exact regular black hole solution has been recently obtained with basis on the Einstein-dual nonlinear electrodynamics as proposed by Salazar, Garcia and Plebansky \\cite{SP87,BG98}. Note, however, that the purpose of the present work is not to give a detailed description of the possible nonlinear theories and their physical effects. Our basic aim is rather modest - we will try to gain new insights into the possibilities beyond Maxwell theory based on the simplest nonlinear electromagnetic lagrangian and its connection to the cosmic singularity problem.  In this concern, it has been found \\cite{novello98a} that the primordial singularity can be removed because the nonlinear corrections reinforce the negative pressure at the early stages of the universe. The authors also argued that the nonsingular behaviour is unaffected by the presence of other ultrarelativistic components obeying the equation of state $p_{(ur)} = \\rho_{(ur)} /3$.  In the present work, we extend the analysis by De Lorenci et al.\\cite{novello98a} using different ingredients. We first obtain the behavior of the scale function for a spatially flat FRW geometry ($\\Lambda=0$), thereby showing that their results is a particular case of the general solution. It is also analyzed how this solution is modified by the presence of a $\\Lambda$-term (vacuum energy density) which may be constant or a time varying quantity. Solutions with constant vacuum energy density also leads to a non-singular universe, and the same happens with a time-dependent $\\Lambda(t)$-term. However, there are singular solutions where the decaying vacuum supplies the energy to a constant cosmological magnetic field. In this case, the time dependence of the cosmological term can uniquely be determined and corresponds to a slightly modification of the more frequent forms suggested in the literature\\cite{OT,CLW92,overduin98}. The main physical restrictions, including the time interval where the nonlinear corrections must be important are also discussed.  The article is organized as follows. In section \\ref{s2}, we write down the Einstein field equations (EFE) for a flat FRW geometry supported by an energy momentum derived from the extended (nonlinear) eletromagnetic Lagrangian. In section \\ref{s3}, we generalize the solution derived in \\cite{novello98a} which assumes a vanishing cosmological term, a spatially flat geometry and a time-dependent magnetic field. In section \\ref{s4}, we obtain a new solution that takes into account the presence of a constant $\\Lambda$. The behavior involving a time-dependent $\\Lambda$ is discussed in section \\ref{s5}. Finally, in the conclusion section, we summarize the basic results and also present some suggestions for future work. In what follows, Greek indices run from 0 to 3, Latin indices run from 1 to 3. Unlike of Ref.\\cite{novello98a}, we adopt the International System of Units (see the Appendix on units and dimensions of Ref. \\cite{jackson} for further details.) ",
        "conclusions": "\\label{s6}   We have examined whether nonlinear corrections to the Maxwell electrodynamics may avoid the cosmic singularity occurring in flat FRW universes. In brief, the answer is positive. We show that by discussing a large class of analytical cosmological models under three different assumptions. In the first case, the cosmological $\\Lambda$-term is identically zero and the dynamics is driven by a time dependent magnetic field. This class generalizes the particular solution previously found by De Lorenci {\\em et al.}\\cite{novello98a}, and confirms their statement concerning the avoidance of the initial singularity. In principle, since the solutions are  non-singular, they potentially  solve the horizon problem. We have also examined if the basic features of such models remain true if new ingredients are introduced in the matter content. In this concern, models with a constant and time dependent $\\Lambda$-term were studied with some detail. Again, for both cases, the universe is also non-singular, bouncing at a critical time when the scale factor reaches its minimum value.  For a decaying vacuum energy density we discuss two different scenarios. In the first one, it was phenomenologically described by $\\Lambda(t) \\sim H^{2}$ as assumed by several authors \\cite{CLW92,overduin98}. These models are nonsingular and resemble the solutions with no $\\Lambda$. The second scenario is a rather curious solution which describes a universe driven by a constant magnetic field. The time behaviour of the cosmological term is now uniquely determined by the EFE as a logarithm of the scale factor. It should be interesting to examine if such results are maintained in the presence of other matter fields, as well as for universes with non-zero curvature.  Finally, in analogy with the cosmological case, one may ask if nonlinear terms in the Maxwell Lagrangian may remove the physical singularity present in a charged black hole (Reissner-Nordstrom solution). This problem will be discussed in a forthcoming communication."
    },
    "0402/gr-qc0402026_arXiv.txt": {
        "abstract": "In this work we study the magnitude-redshift relation of a non-standard cosmological model. The model under consideration was firstly investigated within a special case of metric-affine gravity (MAG) and was recently recovered via different approaches by two other groups. Apart from the usual cosmological parameters for pressure-less matter $\\Omega_{\\rm m}$, cosmological constant/dark energy $\\Omega_{\\lambda}$, and radiation $\\Omega_{\\rm r}$ a new density parameter $\\Omega_\\psi$ emerges. The field equations of the model reduce to a system which is effectively given by the usual Friedmann equations of general relativity, supplied by a correction to the energy density and pressure in form of $\\Omega_\\psi$, which is related to the non-Riemannian structure of the underlying spacetime. We search for the best-fit parameters by using recent SN Ia data sets and constrain the possible contribution of a new dark-energy like component at low redshifts, thereby we put an upper limit on the presence of non-Riemannian quantities in the late stages of the universe. In addition the impact of placing the data in redshift bins of variable size is studied. The numerical results of this work also apply to several anisotropic cosmological models which, on the level of the field equations, exhibit a similar scaling behavior of the density parameters like our non-Riemannian model. ",
        "introduction": "}  Today the observational situation in cosmology is very promising. With several independent cosmological tests at hand the so-called standard model of cosmology \\cite{KolbTurner,PadmanabhanAstro3} has emerged and passed most of these tests. While one of the main benefits of the FLRW (Friedmann-Lema\\^{\\i}tre-Robertson-Walker) model is given by its simplicity, recent observations of type Ia supernovae \\cite{Perlmutter}-\\cite{Riess2} made clear that we only know little about the dominating energy density component of the universe, which enters the general relativistic description in form of a cosmological constant. Such a component is nowadays usually termed dark energy.  Since the cosmological constant seems to be inevitable for the description of the SN Ia data within the FLRW model it is interesting to figure out whether such a concept is still needed within an alternative cosmological model. From the viewpoint of a non-Riemannian gravity theory it would be very satisfactory if such a dark energy component could be ascribed to a deviation from the usual Riemannian geometry.  In this work we analyze the SN Ia data within a cosmological model \\cite{Puetzfeld1,Puetzfeld2} which is based on a non-Riemannian gravity theory namely metric-affine gravity (MAG) \\cite{PhysRep}. As we show in the following the field equations of the model turn out to be very similar to the usual Friedmann equations. We carry out a numerical analysis on the basis of the new field equations. Thereby we put a quantitative constraint on the contribution of the new density parameter $\\Omega_{\\psi}$, which is linked to the non-Riemannian structures of the underlying gravity theory and which contributes to the total energy density in the universe.  The plan of the paper is as follows. In section \\ref{EXTWEYL_magnitude_redshift_relation_section} we derive the magnitude-redshift relation from the field equations found by Obukhov \\etal in \\cite{Obukhov}. With this relation at hand we perform fits to the combined SN Ia data sets of Wang \\cite{Wang} and Tonry \\etal \\cite{Tonry} in section \\ref{Numerical_results_section}. In addition we study the impact of placing the data sets in redshift bins of variable sizes. We compare our findings with the results from independent age estimates  of the universe and work out the deceleration history of our model. Finally, we draw our conclusion in section \\ref{CONCLUSION_section}. In \\ref{TRIPLET_KAPITEL} we briefly introduce the field equations and geometrical quantities of metric-affine gravity (MAG) and provide an introduction to the so-called triplet ansatz of MAG. Readers who are not familiar with MAG might also want to consult \\cite{PhysRep} for a comprehensive review. In \\ref{OPERATIONS_SECTION} we fix our conventions and in \\ref{NATURAL_UNITS} we provide an overview over the units used throughout the preceding sections. ",
        "conclusions": "\\paragraph{Age} For the best-fit parameters listed in table \\ref{TABLE_triplet_best_unbinned} and table \\ref{TABLE_triplet_best_binned} we obtain the following values for the age of the universe: (WC92, $11.77$ Gyrs), (WC230, $11.78$ Gyrs), (WC114, $9.48$ Gyrs), (WC25, $7.52$ Gyrs). Our best-fit parameters yield cosmological models with younger\\footnote{Younger with respect to the standard FLRW model with $\\Omega_{{\\rm m} 0}=0.3$, $\\Omega_{\\lambda 0}= 0.7$, and $H_0=65 \\, {\\rm km}\\, {\\rm s}^{-1} {\\rm Mpc}^{-1}$, for which one obtains an age of $14.5$ Gyrs.} age than the estimates from globular clusters and nuclear cosmochronology which range from 11 to 15 Gyrs \\cite{Chaboyer,Truran}, but we should keep in mind that the uncertainties in the parameters allow a broad range of age of the Universe. When one considers the whole allowed interval as displayed in figures \\ref{FIG_triplet_contours_lambda}, \\ref{FIG_triplet_contours_lambda230}, and \\ref{FIG_triplet_contours_lambda230b}, the age of the Universe is not a problem.  In figure \\ref{FIG_age} we plotted the variation the age for varying $\\Omega_{\\psi 0}$. Regarding to this plot there should be no conflict with the age of the oldest objects in the universe as long as $\\Omega_{\\psi 0}$ is smaller than $0.01$.  \\begin{figure} \\begin{center} \\epsfig{file=age230.ps,width=9.6cm,angle=-90} \\end{center} \\caption[Age function]{Age of the universe for varying $\\Omega_{\\psi 0}$ parameter. The other parameters are fixed to the values as listed in table \\ref{TABLE_triplet_best_unbinned} for the WC230 fit. The shaded region corresponds to the age of the universe as inferred from globular clusters and nuclear cosmochronology \\cite{Chaboyer,Truran}.} \\label{FIG_age} \\end{figure}  \\paragraph{Deceleration parameter} For a model with pressure-less matter, radiation, cosmological constant, and the additional contribution from $\\Omega_\\psi$, the deceleration parameter is given by \\begin{equation} \\hspace{-2.5cm}q=\\frac{\\Omega _{{\\rm m}0}\\left( 1+z\\right) ^{3}-2\\Omega _{\\lambda 0}+4\\Omega _{\\psi 0}\\left( 1+z\\right)^{6}+2 \\Omega _{{\\rm r}0}\\left( 1+z\\right) ^{4}}{2 \\left[ z \\Omega _{{\\rm m}0}\\left( 1+z\\right) ^{2}-z \\Omega _{\\lambda 0 }\\left( 2+z\\right) +z\\Omega _{\\psi 0} \\sum\\limits_{i=2}^{5}\\left( 1+z\\right) ^{i}+z\\Omega _{{\\rm r} 0} \\sum\\limits_{i=2}^{3}\\left( 1+z\\right) ^{i}+ \\left( 1+z\\right) ^{2}\\right]}. \\label{EXT_deceleration_parameter} \\end{equation} The deceleration factor versus the redshift is plotted in figure \\ref{FIG_deceleration_parameter}. Therefrom we can infer that the presence of an additional (positive) $\\Omega _{\\psi 0}$ contribution would lead to a later transition to the accelerating epoch than in the standard model with $\\Omega _{{\\rm m}0}\\approx0.3$ and $\\Omega _{{\\lambda} 0}\\approx 0.7$.  \\begin{figure} \\begin{center} \\epsfig{file=deceleration_parameter.ps, width=9.6cm,angle=-90} \\end{center} \\caption[Deceleration parameter]{Deceleration parameter versus redshift for different parameter choices. We make use of the notation $(\\Omega_{{\\rm m} 0},\\Omega_{{\\lambda} 0},\\Omega_{{\\psi} 0},\\Omega_{{\\rm r} 0})$.} \\label{FIG_deceleration_parameter} \\end{figure}   \\paragraph{Future measurements}  In figure \\ref{FIG_mu_z_both_models} we plotted the distance modulus versus the redshift for the FLRW model as well as for our new model. As becomes clear from this plot the new model would not be distinguishable from the FLRW model if we fix the new $\\Omega_{\\psi 0}$ density parameter to very small values. If one sticks to the upper limit for $\\Omega_{\\psi 0}$ of about $0.3$, which we obtained by fitting to the data set of Wang, it should be possible to distinguish our model from the FLRW scenario as soon as more data at redshifts between $z=1$ and $z=2$ becomes available. In order to quantify the difference in $\\mu$ between the FLRW and our model, we displayed $|\\Delta \\mu|$ in figure \\ref{FIG_mu_z_abs_both_models} for several values of the $\\Omega_{\\psi 0}$ density parameter.  \\begin{figure} \\begin{center} \\epsfig{file=muz_diff.ps,width=9.6cm,angle=-90} \\end{center} \\caption[$\\mu - z $ relation]{Distance modulus versus redshift for the FLRW and our non-standard model. We make use of the notation $(\\Omega_{{\\rm m} 0},\\Omega_{{\\lambda} 0},\\Omega_{{\\psi} 0},\\Omega_{{\\rm r} 0},H_0)$. For very small values of the $\\Omega_{\\psi 0}$ density parameter our model is virtually indistinguishable from the FLRW at low redshifts. The red line corresponds to the upper limit on $\\Omega_{\\psi 0}$ of about $0.3$ as inferred from our previous fits.} \\label{FIG_mu_z_both_models} \\end{figure}  \\begin{figure} \\begin{center} \\epsfig{file=muz_diff_abs.ps,width=9.6cm,angle=-90} \\end{center} \\caption[$\\mu - z $ relation]{Absolute difference between the distance modulus of the FLRW and our new model for different values of the $\\Omega_{\\psi 0}$ parameter. The rest of the cosmological parameters is fixed to the following values: $\\Omega_{{\\rm m} 0}=0.3,$ $\\Omega_{{\\lambda} 0}=0.7,$ $\\Omega_{{\\rm r} 0}=0.3,$ $H_0=65$ ${\\rm km \\, s}^{-1}{\\rm Mpc}^{-1}$.} \\label{FIG_mu_z_abs_both_models} \\end{figure}  \\paragraph{Summary}  Most interestingly there exist Weyl-Cartan based scenarios in which the usual factorization of the density factors is applicable, in contrast to the more general scenario discussed in \\cite{Puetzfeld1,Puetzfeld2} the model presented in this work allows for such a simple relation between the density parameters. We worked out the magnitude-redshift relation and managed to pin-down the cosmological parameters by using recent SN Ia data. Moreover we have shown that our results are also applicable to the models of two other groups. The model discussed here turns out to be a viable cosmological model from the viewpoint of the SN Ia observations.  The best-fit parameter sets WC92, WC230, and WC114 reveal a clear trend to low values for the pressure-less matter component $\\Omega_{{\\rm m} 0}$. This is a rather desirable behavior since one would prefer a model which predicts rather small values for $\\Omega_{{\\rm m} 0}$ and thereby lies closer to the results for the baryonic matter component as inferred from nucleosynthesis \\cite{Olive2}.  Unfortunately our model does also favor a non-vanishing cosmological constant, as it is also favored by the standard FLRW model. Hence, in order to get rid of the cosmological constant one would have to change either the Lagrangian in (\\ref{Obukhov_Lagrangian}) or the ansatz for non-Riemannian field strengths, i.e.\\ torsion and nonmetricity, which in this model are determined by the one form in (\\ref{EXT_obukhov_triplet_1_form}). Nevertheless it should be possible to work out a modified model which removes the need for the cosmological constant and thereby the need for a rather mysterious dark energy component.  We have to stress at this point that our goal was to perform fits to the available SN Ia data sets without any prerequisites. Future work will concentrate on the determination of the cosmological parameters \\cite{Peebles1,Lineweaver} with the help of other independent cosmological tests like nucleosynthesis \\cite{Olive2} and the fluctuations within the cosmic microwave background \\cite{Jaffe,WMAP1}.  In addition one could investigate the impact of the recently released SNe data set from the Deep Survey \\cite{Barris} on the fit result.  Most interestingly Daly \\etal suggested a test on the basis of recent radio galaxy measurements reaching out to $z \\approx 1.8$ which is currently not covered by the supernova data. As we have shown before the availability of data at higher redshifts should increase our capability of discriminating between our model and the usual FLRW scenario. Hence the analysis of this data set might lead to some interesting new insights.   \\ack The authors are grateful to Y.\\ Wang (Oklahoma) for providing her SNe data set. Additionally, one of the authors (DP) wants to thank C.\\ Raas (Cologne) for his patient and enduring advice on several computer related issues. XC is supported by the NSF grant PHY99-07949.\\bigskip  \\appendix"
    },
    "0402/astro-ph0402061_arXiv.txt": {
        "abstract": "New high-resolution, wide-field 90~cm VLA\\footnote{% The National Radio Astronomy Observatory (NRAO) is a facility of the National Science Foundation operated under a cooperative agreement by Associated Universities, Inc.} observations of the Galactic Center region by Nord et al.\\ have revealed twenty nonthermal filament (NTF) candidates.  We report 6~cm polarization observations of six of these.  All of the candidates have the expected NTF morphology, and two show extended polarization confirming their identification as NTFs.  One of the new NTFs appears to be part of a system of NTFs located in the \\objectname[]{Sgr~B} region, 64~pc in projection north of \\objectname[]{Sgr~A}. These filaments cross the Galactic plane with an orientation similar to the filaments in the Galactic Center Radio Arc.  They extend the scale over which the NTF phenomena is known to occur to almost 300~pc along the Galactic plane. Another NTF was found in the Galactic plane south of the \\objectname[]{Sgr~C} filament but with an orientation of~45\\arcdeg\\ to the Galactic plane.  This is only the second of~12 confirmed NTFs that is not oriented perpendicular to the Galactic plane.  An additional candidate in the Sgr~C region was resolved into multiple filamentary structures.  Polarization was detected only at the brightness peak of one of the filaments.  Several of these filaments run parallel to the Galactic plane and can be considered additional evidence for non-poloidal magnetic fields at the GC. Together the 90 and~6~cm observations indicate that the Galactic center magnetic field may be more complex than a simple globally ordered dipole field. ",
        "introduction": "\\label{sec:intro}  Understanding the origin, evolution, topology and strength of the Galactic center (GC) magnetic field(s) is a necessary prerequisite for constructing a coherent picture of activity at the \\hbox{GC}.  The discovery of the \\objectname[]{Galactic Center Radio Arc} (GCRA) \\citep{y-zmc84} demonstrated that magnetic structures with ordering on the scale of $\\sim 40$~pc exists in the \\hbox{GC}.  The GCRA may also be part of an even a larger scale structure known as the \\objectname[]{Omega lobe} \\citep{sh84,s85}, which is a loop-like structure extending several hundred parsecs from the GCRA and may connect with the Sgr~C filament at its other footpoint. Ejection or expansion of magnetic fields from an accretion disk \\citep{uss85,hnp98} or a shocked Galactic wind \\citep{c92,b-hc03} are possibilities for the origin of these structures.  In addition, the discovery of isolated nonthermal filaments throughout the inner few hundred parsecs \\citep[e.g.,][]{y-zb89,m94,lklh00} with orientations largely perpendicular to the Galactic plane have been interpreted as evidence for a space-filling poloidal field in the ionized medium \\citep{ms96,m98}.  A pervasive poloidal field could result from the inflow of a primordial halo field toward the high concentration of mass in the GC \\citep{sf87,ccm00,c01} or from a nuclear star burst as stellar winds and supernovae inject stellar fields into the surrounding medium \\citep{w02}.  Black hole accretion disks can also generate magnetic fields \\citep[e.g.,][]{crv94} that subsequently spread throughout the galaxy \\citep[and references therein]{w02,dl90,k02}.  Thus, both ejection of dynamo generated fields and inflow and concentration of primordial fields could be acting independently to create the observed structures.  Although these theoretical ideas are viable and reasonable they have not been developed with sufficient rigor and detail to establish any definitive conclusions on the nature of magnetism at the \\hbox{GC}.  Observationally the nonthermal filaments (NTFs) are unique to the GC and are defined and identified by (1)~Extreme length to width ratios (tens of parsecs long and only a few tenths of a parsec wide); (2)~Nonthermal spectral indices; and (3)~High intrinsic polarization \\citep[for recent reviews see][]{bl01,lnlk03}.  Faraday rotation measurements indicate that the NTF magnetic fields are aligned longitudinally \\citep[e.g.,][]{tihtksk86,y-zwp97,lme99}.  Of the nine previously identified isolated NTFs---the most recent NTF discovery is described by \\cite{r03}---eight of them, along with the bundled NTFs in the GCRA, are within 20\\arcdeg\\ to the perpendicular to the Galactic plane.   From the distribution and lengths of NTFs, it has been suggested that the magnetic field in the ionized medium is dominated by a poloidal field that extends many tens of parsecs from the Galactic plane \\citep[e.g.,][]{ms96,m98}.  [The lone NTF parallel to the plane, assuming a GC distance of 8 Kpc \\citep{r93}, is located southwest about~75~pc in projection from the Galactic plane and the farthest from \\objectname[]{Sgr~A}, approximately 225~pc in projection. It therefore could be in a region where the field lines are merging with the toroidal field in the disk \\citep{laklg99}.]  Several estimates for the strength of NTF magnetic fields can be made. Estimates based on radio continuum measurements and assuming equipartition field strengths are of the order of a few tenths of a milliGauss for several filaments (the \\objectname[]{Northern Thread}, \\citealt{lme99}; the \\objectname[]{Sgr~C} filament, \\citealt{lklh00}; the \\objectname[]{Snake filament}, \\citealt{gnec95}).  The absence of any bending or distortion of the NTFs against the strong ram pressure of the GC molecular clouds indicates a field strength of about~1~mG (\\citealt{y-zm87a,y-zm87b}; see \\citealt{c01} for further discussion of this point).  Such a field strength is considerably stronger than the general interstellar magnetic field, which is typically no more than a few tens of microGauss \\citep[e.g.,][]{c99}.  The implied internal magnetic pressures of the NTFs are so large that unless they are confined by a comparable pressure, their lateral dimensions would expand well beyond the observed few tenths of a parsec on timescales short compared to their synchrotron lifetime.  Confinement is not an issue if there is a $\\sim 1$~mG field filling the entire region.  In that case the NTFs are those flux tubes that happen to be illuminated with high energy electrons that are energized by some local interaction.  The energy density of such a space-filling field is quite high, of the order of $4 \\times 10^{-8}$~ergs~cm${}^{-3}$ and the total magnetic energy is $\\sim 6 \\times 10^{54} (R/75)^2(L/300)$~ergs, where $R$ and $L$ are the radius and length (in pc) of the cylindrical region over which the NTF phenomenon is found.  However, alternatives to the pervasive field interpretation are also viable (e.g., Yusef$-$Zadeh 2003; \\cite{sl99}). In the cometary model of \\cite{sl99}, the NTFs are magnetized wakes generated by the interaction of molecular clouds with a GC wind.  Considerable evidence for a Galactic wind is derived from X-ray observations of hot expanding gas \\citep{kmstty96}.  The existence of a GC wind and related outflows is discussed in the context of recent high resolution infrared observations by \\cite{b-hc03}.  In the cometary model the wind is assumed to advect a weak $\\sim 10$~$\\mu$G magnetic field that is amplified by a factor of roughly 100 as the field is stretched by the flow and wraps around the obstacle, a molecular cloud, forming an elongated magnetotail \\citep[for simulations, see][]{gmrj00,dels02}).  In this scenario the total required magnetic energy is reduced by several orders of magnitude and presumably NTFs with every orientation could be observed since they are viewed in projection.  The above discussion indicates that additional observations are required to to establish a consensus interpretation of the NTF phenomenon and it's relationship to galactic center magnetism. This paper reports new total and polarized intensity observations of a number of NTF candidates that challenge the interpretation of a space-filling globally ordered poloidal field.  In \\S\\ref{sec:observe} we describe the observations and analysis and in \\S\\ref{sec:discuss} we discuss the implications of our observations. ",
        "conclusions": "\\label{sec:discuss}  The 6~cm observations show that all the observed NTF candidates are filamentary in nature, with a several of them showing polarization.  Although only 3 of the 6 candidates exhibited the expected polarization, the low signal-to-noise would have precluded detecting polarized emission from the remainder, even if they were polarized at the level typical of NTFs.  Generalizing these results, we conclude that many of the 90~cm NTF candidates probably are true NTFs.  Consequently there may be a large population of shorter, low surface brightness NTFs.  Are these shorter, low surface brightness NTFs the same class of phenomenon as the longer, brighter NTFs that are already known? \\cite{nbhlklad03} shows that the typical surface brightness of the new filaments is approximately 20~\\mjybm\\ at~90~cm; the surface brightness of the previously known filaments is 80~\\mjybm.  The flux density $S_\\nu$ as a function of frequency~$\\nu$ from a synchrotron source of volume~$V$ is $S_\\nu \\propto N_0 VB^{-\\alpha+1} \\nu^{\\alpha}$ where $N_0$ is a number density of synchrotron-emitting electrons and~$B$ is the magnetic field within the volume.  The widths of the new filaments are comparable to those of the previously known, prominent NTFs, of order a few tenths of a parsec.  Assuming that these are approximately cylindrical objects, then the volume element per unit length is the same for both the previously known and newly identified filaments. For a spectral index $\\alpha \\approx -1$, the luminosity of a synchrotron source scales with magnetic field as $B^2$.  Therefore, in order to reproduce the observed range in surface brightnesses requires that the product $N_0B^2$ vary by only a factor of~4.  Either a factor of~4 variation in the synchrotron-emitting electron number density or a factor of~2 in the magnetic field strength is sufficient.  Of course, if both $N_0$ and~$B$ vary, even smaller variations in each are required in order to obtain the required variation in $N_0B^2$. We conclude that the same physical process could be responsible for the entire population with small local variations in magnetic field strength and/or energetic particle density accounting for any differences.  A larger number of NTFs offers additional insights into the structure of any large$-$scale magnetic field.  \\objectname[NTF]{NTF~359.32$-$0.16} is oriented at an angle of~45\\arcdeg\\ to the Galactic plane and is only the second NTF that is not oriented more-or-less perpendicularly with respect to the Galactic plane.  It is also the considerably closer to \\objectname[]{Sgr~A} (87~pc in projection) than the other parallel NTF, \\objectname[NTF]{NTF~358.85$+$0.47} (the Pelican).  There is considerable evidence that the magnetic field in the neutral medium along the Galactic plane in the \\objectname[]{Sgr~A} and \\objectname[]{Sgr~B} regions is parallel to the plane \\citep[e.g.,][]{novaketal03}\\footnote{% Additional results on the GC magnetic field are derived from submillimeter and far infrared observations of the neutral medium.  The polarized infrared emission from dust grains embedded within the molecular clouds along the Galactic plane suggest that the magnetic field is parallel to Galactic plane and therefore toroidal \\citep{novaketal03}.  Estimates for the strength of this field also indicate a few mG \\citep{cnhdvdd03}.} suggesting that there is a large-scale toroidal field in the neutral medium.  Given its 45\\arcdeg\\ orientation, it seems unlikely that this NTF is related to a toroidal field.  Furthermore SCUBA images \\citep{pierce-priceetal00} of the Galactic center indicate that the thermal emission from the \\objectname[]{Sgr~C} environment is much less than in the \\objectname[]{Sgr~A} or \\objectname[]{Sgr~B} regions.  Thus, the region near \\objectname[]{Sgr~C}, including \\objectname[NTF]{NTF~359.32$-$0.16} may not be dominated by the neutral medium.  Figure~\\ref{fig:fig11.eps} is schematic of the GC region based on the wide$-$field image of Nord et al (2003a,b).  The confirmed NTFs (including those from this work) are shown in the heavy dark lines while the NTF candidates are denoted by the thin lines.  Confirmation of additional candidates would definitively rule out a global dipole field model.  However, the present data is already difficult to reconcile with global ordering.   Figure 10 shows that there are several cases where two NTFs, which are close together in projection, have very different orientations.  In order to preserve $\\nabla \\cdot B =0$ in a region of space, distinct field lines must be more or less parallel on scales smaller than any gradient scale.  The northern end of \\objectname[NTF]{NTF~359.32$-$0.16} is 15~pc in projection from the \\objectname[]{Sgr~C} filament and has an orientation $\\sim 45\\arcdeg$ to it.  The \\objectname[]{Sgr~C} filament is about~30~pc long and, although it has a nonuniform brightness, it exhibits a constant spectral index with length, suggesting that it is tracing a uniform magnetic field.  For a dipole magnetic field $B/\\nabla B = r/3$, where $r$ is the radial distance from the dipole.  For a dipolar field centered on \\objectname[]{Sgr~A} at the location of \\objectname[NTF]{NTF~359.32$-$0.16} (87~pc in projection) the gradient scale length is almost 30~pc, consistent with the \\objectname[]{Sgr~C} \\hbox{NTF}. However, if \\objectname[NTF]{NTF~359.32$-$0.16} is at the same distance with a $45\\arcdeg$ difference in orientation, it could not be tracing the same global, divergence-free field.  Another example is the northern thread \\objectname[NTF]{NTF~0.08$+$0.15} and the NTF candidate \\objectname[]{G359.90$+$0.19}.  These two are separated by only about~5~pc in projection and both are about~30~pc in projection from \\objectname[]{Sgr~A}, yet their orientations differ by~25\\arcdeg.  These differences in orientation cannot be explained as due to differences in their relative distances.  While little is known about their relative separations, as noted earlier the NTF phenomenon occurs over a linear distance of roughly 300~pc along the Galactic plane ($-1\\arcdeg \\lesssim \\ell \\lesssim 0\\fdg4$).  It is reasonable to assume that the NTFs are also distributed over a comparable distance along the line of sight.  If this is the case, large apparent changes in orientation could be obtained in a dipolar field by filaments well above the Galactic plane and located at different distances along the line of sight (e.g., the northern thread \\objectname[NTF]{NTF~0.08$+$0.15} and the NTF candidate \\objectname[]{G359.90$+$0.19}).  Even so, if the GC magnetic field is poloidal and dominated by a dipolar component, all filaments crossing the Galactic plane should be nearly perpendicular to the Galactic plane, regardless of their distance along the line of sight.  The different orientations of the \\objectname[]{Sgr~C} filament and \\objectname[NTF]{NTF~359.32$-$0.16} demonstrate that the GC magnetic field cannot be dominated by a largely dipolar, poloidal component.  The new 90~cm survey coupled with the 6~cm observations described here indicate that the population of galactic center NTFs is significantly larger than previously thought.  These new NTFs can not be interpreted as components of a globally ordered poloidal magnetic field and suggest a more complicated field structure. Further, high sensitivity observations, especially with new instrumentation such as the expanded VLA and LOFAR (see Kassim et al 2003; White, Kassim \\& Erickson 2003) will be required to discover and characterize the entire population of Galactic center NTFs."
    },
    "0402/astro-ph0402582_arXiv.txt": {
        "abstract": "We present chromospheric \\caii\\ activity measurements, rotation periods and ages for $\\sim$1200 F-, G-, K-, and M- type main-sequence stars from $\\sim$18,000 archival spectra taken at Keck and Lick Observatories as a part of the California and Carnegie Planet Search Project.  We have calibrated our chromospheric S values against the Mount Wilson chromospheric activity data.  From these measurements we have calculated median activity levels and derived \\rphk, stellar ages, and rotation periods from general parameterizations for 1228 stars, $\\sim$1000 of which have no previously published S values.  We also present precise time series of activity measurements for these stars. ",
        "introduction": "The California \\& Carnegie Planet Search Program has included observations of $\\sim$2000 late-type main-sequence stars at high spectral resolution as the core of its ongoing survey of bright, nearby stars to find extrasolar planets through precision radial velocity measurements \\citep[e.g.,][]{Cumming1999,Butler2003}.  One source of error in the measured velocities is that due to ``photospheric jitter'':  flows and inhomogeneities on the stellar surface can produce variations in the measured radial velocity of a star and may even mimic the signature of planetary companions (\\citet{Henry2002,Queloz2001,Santos2003}). In addition to providing a precision radial velocity, each of our radial velocity observations provides a measurement of the strength of the stellar chromospheric \\caii\\ emission cores.  These measurements are an indicator of stellar magnetic activity and can provide an estimate of the photospheric jitter and rotation period of a star, both critical values for understanding and interpreting the noise present in radial velocity measurements \\citep{Noyes1984,Saar2000,Santos2000}.  The largest campaign to measure and monitor \\caii\\ emission has been the Mount Wilson program begun by O. C. \\citet{Wilson1968} and continued and improved since then by \\citet{Vaughan1978} and others \\citep{Baliunas1998}. From 1966-1977 this program used the ``HKP-1'' photometer which employed a photoelectric scanner at the coud\\'e focus of the 100 inch telescope.  Since 1977 the ``HKP-2'' photometer has been used, which is a new, specially designed photomultiplier mounted at the Cassegrain focus of the 60 inch telescope \\citep[e.g.,][]{Baliunas1995}.  \\citet{Duncan1991} published data from this program in the form of ``season averages'' of H \\& K line strengths from 65,263 observations of 1296 stars (of all luminosity classes) in the Northern Hemisphere, and later as detailed analyses of 171,300 observations of 111 stars characterizing the varieties and evolution of stellar activity in dwarf stars.  This program defined the Mount Wilson ``S value'' which has become the standard metric of chromospheric activity.  \\citet{Henry1996} published data from a survey of stars in the southern hemisphere, providing S values from 961 observations of 815 stars. Other surveys include the Vienna-KPNO survey \\citep{Strassmeier2000}, whose motivation was to find Doppler-imaging candidates by using \\caii\\ as a tracer of rotation period in 1058 late-type stars, and that of our Anglo-Australian Planet Search \\citep{Tinney2002}, which reported S values for 59 planet search stars not observed by previous surveys. ",
        "conclusions": "We have measured chromospheric activity as S values from over 15,000 archival spectra taken over the course  of the California and Carnegie Planet Search Program. These spectra were taken with the HIRES spectrograph at Keck Observatory and the two detectors at the Hamilton spectrograph at Lick observatory and contain both precision velocity information and the \\caii\\ lines from which S values were derived.  Extraction of activity measurements from the Keck spectra was successful, with over 95\\% of all Keck spectra used yielding useful S values.  The Lick spectra were more problematic: only $\\sim 50\\%$ of the spectra proved useful due to low signal to noise ratios and poor extraction in this very blue region by the automated extraction pipeline.  Nonetheless, the 1400 good Lick measurements and the 14514 good Keck measurements combine to give a record of the chromospheric activity for over 1000 late-type main-sequence stars.  Analysis of the measured activity level of $\\tau$ Ceti demonstrates a typical per-observation random error of 6\\% at Keck and 6.5\\% and 10.5\\% for the two detectors at Lick.  We combined the Keck data with the good Lick data to create median S values in 30-day bins.  We have generated ``grand S'' by taking the median of these monthly values.  These grand S's represent median activity levels for our program stars for the periods we observed them.  For stars with $0.5<B-V<0.9$, there are well-calibrated relationships between mean activity level, age, and rotation period, allowing the determination of those quantities for our stellar sample.  We present our grand S values and derived ages and rotation periods in Table~\\ref{punchline}.  We also have measured differential S values for each stars which are more precise.  We present these data electronically as Table~\\ref{seasonsample}.  For each star, these measurements of activity, $S$ and \\rphk\\  represent median activity levels over the duration of the observations, which may be significantly shorter than a stellar activity cycle.  This represents a source of error when deriving stellar properties such as age from these measurements.  The differential S values, \\sdiff, measured here are much more precise than $S$ but no more accurate: we scaled them to the median of the S values, so while they are excellent measurements of temporal variations in activity for a given star they contain no additional information regarding differences in the overall activity among stars.  We have also measured metallicities for many of these stars and in a future work \\citep{Wright2004} will present an analysis of these same stars for stellar characteristics, chromospheric periodicities, and evolutionary status, including the effects of metallicity from \\citet{Valenti2004} and \\citet{Fischer2004}."
    },
    "0402/astro-ph0402257_arXiv.txt": {
        "abstract": "We investigate the stellar populations of a sample of Tidal Dwarf Galaxies, combining observations and evolutionary synthesis models to try and reveal their formation mechanism. On optical images we select a first sample of TDGs for which optical spectroscopy is used to measure metallicities and velocity structure. Finally, we estimate ages, burst strengths, and stellar masses from near-infrared imaging in comparison with a dedicated grid of evolutionary synthesis models, to assess if Tidal Dwarfs are formed out of collapsing gas clouds or by an accumulation of old stars from the parent galaxy or by a combination of both. ",
        "introduction": "The old idea of Zwicky(1956) that (dwarf) galaxies could be formed during encounters of giant galaxies has received a lot of observational support within the last decade (Duc \\& Mirabel, 1994, Duc \\& Mirabel, 1998, Hibbard et~al., 1994). This data revealed the so called {\\it Tidal Dwarf Galaxies} ({\\bf TDG}s) to be gas rich, dwarf galaxy sized knots with stellar masses of $10^6\\dots10^8$\\,M$_\\odot$ made of recycled material, still embedded in the tidal tail of their parent galaxy. Often, a young starburst on top of an older stellar population is observed. This is widely interpreted as formation of a new galaxy ``in situ'' in the tidal tails.  To tell the phenomenon of TDGs apart from e.g.~super star clusters or chance alignments of matter, the following definition of TDGs was introduced (Duc et~al., 2000; Weilbacher \\& Duc, 2001):\\\\ A Tidal Dwarf Galaxy is {\\it a self-gravitating object with the mass of a dwarf galaxy, formed from recycled material that was expelled from a big galaxy in the course of tidal interaction}.  Why are TDGs thought to be significant? First, they represent a special mode of galaxy formation which can be observed in the local universe, i.e.~much more detailed than high-redshift galaxies forming in the early universe. Second, part of the dwarf galaxy populations of the Local Group or nearby galaxy clusters may have been formed as TDGs (Okazaki \\& Taniguchi, 2000). And finally, they may explain part of the faint blue galaxy population at intermediate redshifts. At those times, the merger rate was higher than today (Conselice et al., 2003) and the galaxies were more gas rich, so that a large number of bright TDGs could have been formed.  But how do TDGs form? Are they preferentially built as stellar knots of which there may be several per merger as seen in the dynamical models of Barnes \\& Hernquist (1992), or from giant gas clouds at the tips of the tidal tails which then afterwards may also attract stars into their potential as in the simulations of Elmegreen et~al. (1993)? This question was the original starting point of our project. We therefore investigate the stellar populations of a sample of TDGs in order to find out which of these models is realized or more common in nature. ",
        "conclusions": "We used a three stage approach of optical photometry, optical spectroscopy, and near-infrared photometry combined with dedicated evolutionary synthesis models to select a sample of TDG (candidates), and study their gaseous properties and stellar populations.  We selected 44 TDG candidates from optical imaging in comparison with our models from over 100 knots in the tidal tails of our sample of 14 interacting and merging galaxies. From optical spectroscopy we confirm the association of the TDG candidates with the parent galaxies. The metallicity measurement shows them to have a mean oxygen abundance close to 1/4\\,$Z_\\odot$, while three TDG candidates were rejected due to their too low metallicities. Low resolution spectroscopy also revealed 13 TDGs to show kinematical signatures independent of their surrounding tidal tail. We therefore roughly find at least one true TDG per interacting system.  From our near-infrared analysis and the detailed comparison of their spectral energy distributions with our evolutionary synthesis models we confirm that the TDGs in our sample indeed have the stellar masses of dwarf galaxies and show strong star-formation activity. Their burst strengths show most of them to have primarily formed from stellar condensations and only a smaller percentage from the collapse of giant gas clouds."
    },
    "0402/astro-ph0402531_arXiv.txt": {
        "abstract": "The importance of far-infrared observations in our understanding of extreme activity in interacting and merging galaxies has been illustrated by many studies.  Even though two decades have passed since its launch, the most complete all-sky survey to date from which far-infrared selected galaxy samples can be chosen is still that of the Infrared Astronomical Satellite (IRAS).  However, the spatial resolution of the IRAS all-sky survey is insufficient to resolve the emission from individual galaxies in most interacting galaxy pairs, and hence previous studies of their far-IR properties have had to concentrate either on global system properties or on the properties of very widely separated and weakly interacting pairs.  Using the HIRES image reconstruction technique, it is possible to achieve a spatial resolution ranging from 30\\arcsec \\ to 1.5\\arcmin \\ (depending on wavelength and detector coverage), which is a fourfold improvement over the normal resolution of IRAS. This is sufficient to resolve the far-infrared emission from the individual galaxies in many interacting systems detected by IRAS, which is very important for meaningful comparisons with single, isolated galaxies.  We present high-resolution 12, 25, 60, and 100 $\\mu$m images of 106 interacting galaxy systems contained in the IRAS Revised Bright Galaxy Sample (RBGS, Sanders et al. 2003), a complete sample of all galaxies having a 60$\\mu$m flux density greater than 5.24 Jy.  These systems were selected to have at least two distinguishable galaxies separated by less than three average galactic diameters, and thus we have excluded very widely separated systems and very advanced mergers.  Additionally, some systems have been included which are more than three galactic diameters apart, yet have separations less than 4\\arcmin, and which are thus likely to suffer from confusion in the RBGS.  The new complete survey has the same properties as the prototype survey of \\citet{hires}.  We find no increased tendency for infrared-bright galaxies to be associated with other infrared bright galaxies among the widely separated pairs studied here.  We find small enhancements in far- infrared activity in multiple galaxy systems relative to RBGS non-interacting galaxies with the same blue luminosity distribution.  We also find no differences in infrared activity (as measured by infrared color and luminosity) between late and early-type spiral galaxies. ",
        "introduction": "In the last two decades it has become apparent that interactions between galaxies can play a significant role in their evolution.  From the early dynamical simulations of \\citet{toomre} to more modern works by Barnes, Hernquist, and others (Barnes \\& Hernquist 1992, and references therein) it has become apparent that interactions and mergers between galaxies can radically alter their morphology by inducing shells, bars, tails, and other tidal features.  Perhaps more importantly, cancellation of angular momentum during the merger process can lead to a radical redistribution of the gas content of the galaxies, with very rapid gas inflow into the galaxy cores.  This supply of fresh material could possibly fuel an active galactic nucleus, or provide the high gas densities needed to lead to a sudden burst of star formation.  There is considerable evidence that enhanced star formation is associated with interacting galaxies (Sulentic 1988, and references therein). The young OB stars that dominate the starburst radiate primarily in the optical and ultraviolet, but surrounding gas and dust reprocesses this radiation and thus strongly radiates at thermal wavelengths in the far-infrared.  Far-infrared luminosity is thus indicative of the magnitude of recent star formation activity \\citep{telesco,lonsdale}. Additionally, due to the increased temperature of the heated dust, we expect the far-infrared colors to be a good diagnostic of enhanced star formation.  Therefore, many studies have therefore concentrated on the far-infrared properties of interacting galaxies.  Several studies have also discussed the incidence of multiple bright galaxies being found within a given interacting galaxy system. \\citet{haynes} found that for galaxy pairs separated by 2-10\\arcmin, approximately 10\\% have multiple components brighter than 0.5 Jy at 60$\\mu$m and 1 Jy at 100$\\mu$m.  \\citet{xu} also concluded that in the majority of interacting systems, only one galaxy is infrared bright. These results agree with an earlier work by \\citet{bob} which concluded, based on near- IR colors, that most often only one galaxy in a pair showed signs of unusual activity.  This is an interesting result, because it suggests that specific properties of the interacting galaxies may determine whether or not they become emitters in the far-IR as well.  Testing this hypothesis requires resolution of the individual galaxies in the far-IR, which is the goal of this IRAS study.    The canonical figure used by many authors to delineate interacting versus non-interacting systems is a projected separation of three average galactic diameters, as presumably galaxies this close to one another are also close enough to exert a considerable gravitational effect \\citep{dahari,byrd,hires}.  However, for most galaxies detected by IRAS this typically corresponds to an angular separation of a few arcminutes.  This is less than the resolution normally achieved by IRAS using the 1-d coadder ADDSCAN or the 2-d FRESCO imaging process. As a result, it has been impossible to study the far-infrared properties of the individual galaxies, and most studies have either made assumptions about the distribution of flux between galaxies within the interacting system \\citep{bushouse} or have concentrated on widely separated pairs \\citep{haynes,xu}.  Since previous studies of very widely separated galaxy pairs indicate that in a substantial fraction of interacting systems only one galaxy is unusually active in the far-IR \\citep{xu}, it is necessary to resolve these galaxies in order to properly study those properties such as morphology which are unique to the individual galaxies.  Additionally, \\citet{xu} found evidence that at smaller separations (and hence greater interaction strengths) there was a greater enhancement of far-IR activity.  Therefore it would be valuable if these studies could be extended to smaller separations where more observable changes are taking place.  Development of the Maximum Correlation Method algorithm (MCM, Aumann et al. 1990) for use in IRAS image reconstruction significantly increased the resolution of IRAS observations.  As implemented in the HIRES process, MCM is an iterative image reconstruction technique that involves using the known response functions of the IRAS detectors to scan simulated image estimates which are then compared to the actual detector data.  In this way, a high resolution image estimate is formed.  The result is typically a five-fold increase in resolution, varying from roughly 30\\arcsec$\\times$45\\arcsec \\ at 12$\\mu$m to 72\\arcsec$\\times$130\\arcsec \\ at 100$\\mu$m, with the actual achieved resolution being highly dependent on the geometry of the detector coverage \\citep{hires}. Unfortunately, the HIRES process is extremely computer intensive. When developed, a single field typically took a day or more to process.  As a result, the earlier work by \\citet{hires} was rather limited in scope, with only 23 systems being resolvable.  On a modern computer, this computing time is reduced to approximately 15 minutes, thus making feasible the processing of a substantially larger sample \\footnote{HIRES processing is available from the Infrared Processing and Analysis Center Infrared Science Archive, Jet Propulsion Laboratory, California Institute of Technology {\\it http://irsa.ipac.caltech.edu/IRASdocs/hires\\_over.html}}.  In Section 2 of this paper we present the sample selection criteria for the objects examined here, and present the data reduction techniques used for reconstructing the IRAS images and measuring the galaxy fluxes. The fluxes at each {\\it IRAS} wavelength are presented in tabular form, and contours of the infrared emission are shown overlaid on optical images of the galaxies. Section 3 presents properties of the catalog, and some results derived from them. Appendix A presents additional notes for selected galaxy systems. Finally, in Appendix B we include data for galaxy systems that were originally included in the BGS, but were subsequently dropped from the RBGS after a reanalysis of their fluxes. These objects are provided for the interest of the reader, but do not bear on the analysis of the catalog. ",
        "conclusions": "We have presented an atlas of high resolution IRAS observations of all 106 of the paired (and in many cases interacting systems) in the IRAS Revised Bright Galaxy Sample  with a 60$\\mu$m flux density greater than 5.24 Jy. The atlas contains infrared contours overlaid on optical images and a catalog of fluxes or upper limits in all four IRAS wavebands.  We have presented the infrared luminosities and colors of the paired galaxy sample, and compared them to a sample of isolated galaxies. We find substantially the same results as \\citet{hires}, namely that the paired galaxies have a measurably different distribution of infrared properties than isolated galaxies.  Using morphological optical classifications for the galaxies we conclude that there is no difference between late and early-type spirals in terms of their far-IR properties.  In particular, no significant enhancement is seen in the far-infrared luminosity or color of the late-type spirals as compared to the early-type spirals."
    },
    "0402/astro-ph0402477_arXiv.txt": {
        "abstract": "The background field of the Tadpole galaxy image taken with the Hubble Space Telescope ACS contains 87 clump-clusters that are probably low-inclination versions of the 69 chain galaxies we found in a previous paper. This conclusion follows from the similarity in colors and apparent magnitudes of the brightest clumps in the clump-clusters and chain galaxies, and from the similarity between the distribution of axial ratios for the combined sample and the distribution for normal disk galaxies. These distributions are consistent with chains and clump-clusters being disks viewed at random angles with intrinsic axial ratios of $0.1-0.2$ or less. The disks appear to be gas-rich galaxies at z=1 to 2 that are forming an early generation of stars in a small number of very large star complexes. ",
        "introduction": "Deep high resolution images of distant galaxy fields, such as the GOODS fields (Giavalisco et al. 2004) which include the Hubble Deep Field North (Williams et al. 1996) and the Chandra Deep Field South (Volonteri, Saracco, \\& Chincarini 2000), and the Hubble Deep Field South (Casertano et al. 2000), provide direct evidence for processes related to galaxy formation and evolution. A variety of unusual morphologies is evident, including linear structures called chain galaxies (Cowie, Hu, \\& Songaila 1995; van den Bergh et al. 1996). In the GOODS-South field, Moustakas et al. (2004) identified 275 extremely red objects and classified 41 of them as ``others,'' which included chain and low surface brightness galaxies. The spectral energy distributions of these ``others'' from U through K band did not distinguish between elliptical, disk, and irregular types even though their visual appearances varied greatly.  To study chain galaxies in more detail, we surveyed faint linear features in the Hubble Space Telescope Advanced Camera for Surveys (HST ACS) field of the Tadpole galaxy, UGC 10214, finding 69 chain galaxies, 37 double-core galaxies, 21 tadpole-shaped galaxies, and 32 edge-on normal galaxies distinguished by a central bulge and approximately exponential disk (Elmegreen et al. 2004; hereafter Paper I). All of the chain, double, and tadpole systems were shown to have similarly bright clumps, probably the result of star formation, that are bluer and fainter than the centralized bulges of most normal galaxies. The chain galaxies were also shown to dominate all linear types at apparent I-band AB magnitudes fainter then 24th mag. We suggested that chains are edge-on irregular galaxies that have no bulges or exponential disks (see also Dalcanton \\& Schectman 1996, Reshetnikov, Dettmar, \\& Combes 2003). Two possible examples of faint face-on chains were shown.  We have now surveyed the ACS field of the Tadpole Galaxy for all faint objects larger than 0.5 arcsec, not only those with linear structures. These include 87 possible face-on versions of chain galaxies, 191 normal spiral galaxies of all inclinations, and 34 round, centrally-concentrated galaxies that could be ellipticals. The face-on versions of chain galaxies appear mostly as small collections of blue clumps with very faint emission between. To avoid the presumption that they are complete galaxies before spectra or deeper images are available, we call them ``clump-clusters,\" as in Paper I. The normal galaxies are distinguished by their bright central concentrations of light, which are presumably bulges, and by their approximately exponential disks.  We show here that the clumps in clump-clusters have the same colors and magnitudes as the clumps in linear chain galaxies, that both types of clumps are distinct from the bulges in normal galaxies (which are apparently brighter and extend to redder colors), and that clump-clusters are slightly brighter than chain galaxies -- presumably as a result of projection -- even though their average surface brightnesses are comparable or slightly fainter. We also show that the distribution of axial ratios for chain galaxies and clump-clusters is similar to the distribution for local disk galaxies. These observations demonstrate that chain galaxies are edge-on irregular disks, and their face-on counterparts are visible as clump-clusters. The properties of the chain/clump-cluster systems are then reviewed. ",
        "conclusions": "Chain galaxies appear to be edge-on versions of clumpy, disky, nearly-circular irregular galaxies. They have no prominent bulges or exponential disks like the normal galaxies in the same deep field (Paper I), and there are no obvious spiral arms, bars, or rings as in modern disk galaxies. They contain 3 to 5 enormous clumps, which appear to be star-forming regions $\\sim500$ pc in size given the blue colors and likely range of cosmological $z$ between 1 and 2. These redshifts are consistent with photometric and spectroscopic redshifts of chain galaxies with similar angular sizes measured in the HDF, as noted by comparing our objects with objects in the online interactive redshift images of the HDF North (Fernandez-Soto, Lanzetta, \\& Yahil 1999) and South (Yahata et al. 2000).  At these redshifts, the I passband corresponds to rest frame ultraviolet where local late-type galaxies also have prominent star formation (e.g., Windhorst et al. 2002), although generally with smaller regions.  The disk-like structure of chains and clump-clusters inferred from the distribution of axial ratios suggests that most of them did not form by dissipationless mergers of pre-existing clumps. Such mergers would make more spheroidal systems. On the other hand, numerical simulations of young galaxy disks show gravitational instabilities forming $10^9$ M$_\\odot$ clumps like those observed here; the intermediate stages resemble our clump-clusters and the clumps eventually merge into bulges (Noguchi 1999). Our observations are consistent with this picture. The large sizes of our star-forming regions suggest that gaseous velocity dispersions are high during disk instabilities, perhaps 20-40\\% of the rotation speeds (the ratio of the velocity dispersion to the rotation speed is approximately equal to the square root of the intrinsic ratio of axes). The prominence of clumps rather than spiral arms suggests further that gaseous densities are much higher than the critical tidal density from the rotation curve, i.e., background shear is slow compared to unstable growth. Most likely, the disks are nearly pure gas when the instabilities begin (Noguchi 1999). If the observed clumps do not merge into bulges, then they could form the old stellar disks of today's galaxies. With high velocity dispersions, these old disk components should be thick.  We are grateful to J. Blakeslee, H. Ford, and J. Mack for providing reduced images of the Tadpole field, and to an anonymous referee for helpful comments. B.G.E. is grateful to NSF for support under the grant AST-0205097.  \\clearpage"
    },
    "0402/astro-ph0402194_arXiv.txt": {
        "abstract": "The effective compton cross section in a non-degenerate plasma($n\\ll\\{\\frac{(kT/c)^{2}+2mkT}{h^{2}}\\}^{^{3/2}}$) is investigated in a wide range of temperatures. The results show a decreasing behavior with temperature especially for $kT\\gg m_{e}c^{2}$. The results may be important in phenomena like accretion discs or ultra-relativistic blast waves in GRB models, where the emitted radiation has to pass through a medium containing high energy electrons.\\\\  {\\it keywords}:Compton scattering, high energy physics, gamma-ray bursts, accretion ",
        "introduction": "Compton scattering is one of the most important phenomena in astrophysics that affects the received flux from terrestrial objects. In a non-degenerate low temperature plasma in which $\\frac{n^{2/3}h^{2}}{2\\pi\\:m_{e}}\\ll kT\\ll m_{e}c^{2}$ \\cite{Phillips}, the free electrons can be considered effectively at rest, provided that the photon energy $E=h\\nu$ to be much greater than $kT$. In such a situation the Klein-Nishina formula for Compton scattering accurately works. The Klein-Nishina formula is only dependent on the energy of the incoming photon as measured in the rest frame of the electron \\cite{Stecker}, \\cite{Harwit} : \\begin{equation}\\label{eq:a2} \\sigma_{_{c}}=\\sigma_{_{T}}.\\frac{3}{4} \\{\\frac{1+\\epsilon}{\\epsilon^{2}}[\\frac{2(1+\\epsilon)}{1+2\\epsilon} -\\frac{1}{\\epsilon}\\ln(1+2\\epsilon)]+\\frac{1}{2\\epsilon}\\ln(1+2\\epsilon)- \\frac{1+3\\epsilon}{(1+2\\epsilon)^{^{2}}} \\end{equation} in which $\\epsilon$ is the ratio of the photon energy $E$ to the electron rest energy $m_{e}c^{2}$ (Fig.(1)). But at high temperatures where the electron mean energy is greater than or comparable to $m_{e}c^{2}$, the energy of the photon is highly different relative to different individual electrons. So, since the total Compton cross section is a function of the photon energy in the rest frame of the electrons, one would expect an effective cross section due to the thermal distribution of electrons. The aim of this paper is to obtain an effective cross section for Compton scattering in a wide range of plasma temperatures and photon energies. In Sec. 2 the averaging method is introduced. In Sec. 3 the momentum distribution of electrons in a wide range of temperatures is discussed and the normalization constant of the distribution is evaluated as a function of plasma temperature, and then, a numerically computable form for the effective cross section of Compton scattering is obtained. The results are presented in Sec. 4, as well as a discussion on their importance in GRBs. ",
        "conclusions": "As it can be seen in the three-dimensional plot of Fig.(4), and in the contour plot of Fig.(5), the effective cross section is closely the same as the $\\sigma_{c}$ (Fig.(1)) at temperatures $kT<m_{e}c^{2}$. This is an anticipated result. The deviation begins when $kT>m_{e}c^{2}$, with an overall shift to left at higher and higher temperatures. This behavior could have been expected too. The Compton cross section decreases by increasing of photon energy (Fig.(1)), so, for a photon with energy $E=h\\nu$, the more the temperature is, the more energy in rest frame of individual electrons it has, and furthermore, the less the effective cross section is expected. This simple reasoning may seem wrong of course. Because though it is correct for the electrons moving toward the photon beam but in the rest frame of the electrons moving along the beam the energy of a photon would be less than the corresponding value as measured in laboratory frame and it decreases more and more in higher and higher temperatures, and so, in contrary one may expect the effective cross section to be less and less, as long as these electrons are concerned. Really, the final answer can be found in Eqs.(5) and (8).For ultra-relativistic electrons the distance $dl^{\\prime}$ as seen by the electrons moving in the same direction as the beam one, is $4\\gamma^{2}$ times less than the corresponding value for ones moving in the opposite direction. So, as can be seen in Eqn. (5), despite the greater cross section they have individually, their share in scattering happens to be less than that other electrons, so that in high temperatures the effective cross section as a whole is determined by the electrons moving toward the photon beam. In Fig.(6) the portion of these two groups of electrons in Compton scattering is compared with each other.  In a GRB model prepared by authors \\cite{grb}, a collimated ultra-relativistic ejecta with a Lorentz factor $\\Gamma \\sim 1000$ collides with a dense cloud surrounding the stellar engine. The photons emitted from the shocked medium has to pass through the cloud before entering free space, but the opacity of the cloud is extremely too high to let them escape. Really, in the model the cloud thickness $L$, and its density $n$, are of order $10^{13} cm$ and $10^{17} cm^{-3}$ respectively, which result in an opacity high to $\\sigma_{_{T}} n L \\sim 10^{6}$ which obviously prevents any radiation to escape. But, since in external shock models for GRBs (see \\cite{piran} and \\cite{Katz} for a review) the particles in the shocked medium are expected to have mean energies of order $\\frac{u}{4\\:\\Gamma\\:n}=\\Gamma\\:m_{p}\\:c^{2}\\sim10^{6}\\:m_{e}\\:c^{2}$ , the real optical depth of the medium for $Mev$ photons would be $10^{-6}$ times less than what might be roughly expected in beginning (Table (1) ), and consequently a photon radiated from the shocked matter may succeed to go out of it without being scattered, provided that its passage make an angle larger than $(\\sqrt{\\frac{5}{3}}\\Gamma)^{-1}$ with the velocity vector of shocked matter so that it remain in the shocked medium by the time it cross the opaque cloud. So, we concluded that if the shock front succeed to cross the cloud it might be seen and so might really make a GRB, and if it do not and stop in the cloud, its photons would not be able to cross the cool dense cloud and the phenomenon must be considered a $\\textit{failed}$ GRB. Though the presented reasoning may seem to be restricted only to a medium in thermal equilibrium but it must be correct for all media in which the mean energy of electrons is of order of $10^{6}\\:m_{e}\\:c^{2}$. for example in GRB models, the electrons in the shocked matter are assumed to have a power law distribution: \\begin{equation}\\label{eq:a2} N(\\gamma_{e})\\propto\\gamma_{e}^{2}\\:\\:\\:\\:\\:\\:\\:\\:\\:\\:\\:\\:\\:for :\\:\\:\\:\\:\\:\\:\\:\\:\\:\\:\\:\\gamma_{e}>\\gamma_{e,min} \\end{equation} By repeating the averaging procedure introduced in this paper the effective Compton cross section for such a distribution can be found \\cite{momeni2}. The obtained results are comparable in orders of magnitude to ones presented in Table (1) and shown in Figs.(4) and (5), replacing $\\tau$ with mean Lorentz factor of electrons."
    },
    "0402/hep-ph0402142_arXiv.txt": {
        "abstract": "s{ In the brane-world scenario with low tension, brane fluctuations (branons) together with the Standard Model particles are the only relevant degrees of freedom at low energies. Branons are stable, weakly interacting, massive particles and their relic abundance can account for the dark matter of the universe. In a certain range of the parameter space, they could be detectable by future direct search experiments.} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0402/hep-ph0402232_arXiv.txt": {
        "abstract": "\\hspace*{\\parindent} We study cosmological formation of D-term strings, axionic strings, domain walls and Q-balls in braneworld models of the Hanany-Witten type. For the D-term strings, we show that the strings are the daughter branes extended between mother branes. We show that the domain walls can be produced by conventional cosmological phase transitions. In this case, the formation of the domain walls is induced by the continuous deformation of the branes, which means that they are not created as daughter branes. First we consider classical configurations of the axionic strings and the domain walls, then we investigate the quantum effect of the brane dynamics. We also study brane Q-balls and show how they can be distinguished from conventional Q-balls. ",
        "introduction": "Although there were great successes in quantum field theory, we still have no consistent scenario in which the quantum gravity is included. The most promising scenario in this direction will be string theory where the consistency is ensured by the requirement of additional dimensions and supersymmetry. The idea of large extra dimension\\cite{Extra_1} may solve or weaken the hierarchy problem. In this case, denoting the volume of the $n$-dimensional compact space by $V_n$, the observed Planck mass is obtained by the relation $M_p^2=M^{n+2}_{*}V_n$, where $M_{*}$ denotes the fundamental scale of gravity. The standard model fields are expected to be localized on a wall-like structure, and the graviton propagates in the bulk. The natural embedding of this picture in the string theory context is realized by a brane construction. Inflation with such a low fundamental scale is still an interesting topic\\cite{low_inflation, matsuda_nontach, matsuda_defectinfla}. Other cosmological issue such as baryogenesis with low fundamental scale is discussed in ref.\\cite{low_baryo, Defect-baryo-largeextra, Defect-baryo-4D}, where cosmological defects play important roles. Constructing models for the particle cosmology where non-static brane configurations (such as brane defects and Q-balls\\cite{BraneQball}) are very important. We are expecting that future cosmological observations will reveal the evolution of the Universe, which might also reveal the physics beyond the standard model. To know what kind of brane defects are allowed in the evolution of the Universe, we need to understand how they are formed (and disappeared) in the history of the Universe. In the original scenario for brane inflation\\cite{brane-inflation0}, the inflationary expansion is driven by the potential between branes and anti-branes evolving in the bulk space of the compactified dimensions. The end of inflation is induced by the brane collision where the brane annihilation proceeds through tachyon condensation\\cite{tachyon0}. During brane inflation, tachyon is trapped in the false vacuum. Then the tachyon starts to condensate after inflation, which may result in the formation of the daughter branes.  The production of cosmological defects after brane inflation is discussed in ref.\\cite{Brane-defects, angled-defect}, where it is concluded that cosmic strings are copiously produced but the domain walls are negligible. In ref.\\cite{Majumdar_Davis}, however, it is discussed that all kinds of defects can be produced and the conventional problems of cosmic domain walls and monopoles should arise. Later in ref.\\cite{D-brane-strings, Halyo, BDKP-FI}, the brane production is reexamined and the conclusion was different from \\cite{Brane-defects,angled-defect} and \\cite{Majumdar_Davis}. In ref.\\cite{Brane-defects, angled-defect}, it is discussed that the effect of compactification is significant for the defect formation due to tachyon condensation. It must be useful to make a brief review of the previous arguments about the cosmological formation of brane defects. Their argument is that since the compactification radius is small compared to the horizon size during inflation, any variation of a field in the compactified direction is suppressed. Then the daughter brane wraps the same compactified dimensions as the mother brane. As a result, the codimensions of the daughter branes lie within the uncompactified space. Since the number of the codimension must be even, the defect is inevitably a cosmic string. Moreover, in ref.\\cite{D-brane-strings}, it is pointed that the analysis does not fully account for the effect of compactification, since the directions transverse to the mother brane is not considered. The effect of the RR fields extended to the compactified dimensions is discussed in ref.\\cite{D-brane-strings}. The result is that the creation of the gradients of the RR fields in the bulk of the compactified space is costly in energy, so that the creation of the daughter brane is suppressed if it does not fill all the compactified dimensions. In this case, it was concluded that the production of cosmic strings requires efficient mechanisms, and monopoles and domain walls are not produced after brane inflation.\\footnote{Another kind of defects, which are parameterized by the positions of the branes, were constructed in ref.\\cite{Alice-string} and later in ref.\\cite{incidental}. In ref.\\cite{Alice-string}, brane is replaced by a domain wall that is embedded in the higher-dimensional spacetime, so that one can see what happens in the core. Then the position of a brane in the fifth dimension is used to parameterize the cosmic string in the effective four-dimensional spacetime. The brane is shown to be smeared in the core, so that it resolves the anticipated singularity. Then in ref.\\cite{incidental}, the {\\bf relative} positions between branes are used. In ref.\\cite{incidental}, these defects are called incidental brane defects. In these field-theoretical constructions, branes are replaced by domain walls or vortices embedded in the higher-dimensional spacetime.}  In this paper, however, we show explicit examples where the production of cosmic strings is realized by the formation of daughter branes that are extended between splitting mother branes. It should be noted that we are not considering a counter example of the mechanism of tachyon condensation. The problem of the RR field is avoided, since the length of the extended daughter brane vanishes when it is formed. The tension of the fully extended daughter brane matches to the tension of the D-term string in the effective Lagrangian.\\footnote{In our next paper\\cite{matsuda_angleddefect}, we consider another type of angled brane inflation and solve the problem of the $\\theta$-dependence of the string tension.} Moreover, we also show that other cosmological defects, such as domain walls and Q-balls, can be produced after brane inflation. In our model, it is natural to think that the domain walls and the Q-balls are not produced by the brane creation.\\footnote{Of course it is not impossible to think that such domain walls are the daughter branes being extended between vacuum branes. In our case, however, the cosmological evolution of the brane configuration suggests that they are not produced by the creation of daughter branes. Cosmological formation of domain walls and monopoles, which is induced by the creation of daughter branes, will be discussed in ref.\\cite{matsuda_future}. In this case, the production of the branes extended between branes is crucial.} Domain walls are corresponding to the spatial deformations of the vacuum branes, which can be formed by the thermal effect\\cite{thermal-brane} or brane oscillation after inflation. Thermal effect can induce attractive forces between branes. Then the observer in the four-dimensional spacetime sees the restoration of the corresponding symmetry. The spontaneous breaking of the symmetry triggers the formation of the cosmological defects, which is parameterized by the relative position between branes.  In Section 2, we begin with a short review of a model\\cite{HHK-braneinflation} for brane inflation due to the Hanany-Witten\\cite{Hanany-Witten} type brane dynamics. We show how the extended branes are produced by the brane dynamics after brane inflation. Although our result seems to be contradicting to the previous arguments about daughter brane production, we stress that we are not considering a counter example of the mechanism of tachyon condensation. We think it is not difficult to understand how one can avoid the serious criteria given in ref.\\cite{Brane-defects, D-brane-strings}. The extended branes are formed after brane inflation. The correspondence between the extended brane and the D-term string in the effective action is examined. Then, the formation of axionic strings and domain walls is discussed in Section 3. Unlike the D-term strings, these defects are formed by the spatial deformations of the vacuum branes. The domain wall that we are considering in this paper is different from the usual BPS domain walls\\cite{Witten_Wall} in SQCD and MQCD. In the effective action, we add a soft mass for the adjoint scalar field in $N=2$ SYM, which introduces a shallow potential on the coulomb branch. In the brane counterpart, we are considering $D4$-branes separated by a weak repulsive force between them. Thus our defect configurations are not stable in the supersymmetric limit. Our discussions in this paper compensate the analysis in ref.\\cite{Alice-string} and \\cite{incidental}, in which defects were constructed in the classical brane configurations. It should be noted that the conventional BPS domain walls in MQCD are not suitable for our argument, because they cannot exist in the classical limit. We also consider brane Q-ball\\cite{BraneQball}, which is the configuration of branes in motion. Our conclusions and discussions are given in Section 4. ",
        "conclusions": "In this paper, we have considered the formation of D-term strings, axionic strings, domain walls and Q-balls in the Hanany-Witten type brane dynamics. Here we summarize our conclusions. \\begin{itemize} \\item For the D-term string, we have considered D2-brane that is stretched between the splitting D4-branes. Contrary to the previous arguments, the production of the extended D2-branes is not suppressed. Our arguments are general, because the brane collision with a huge kinetic energy will inevitably induce chaotic process of the production/annihilation and the recombination of the branes, which makes it possible to produce many kinds of extended daughter branes. Further discussions of this topic is given in ref.\\cite{matsuda_angleddefect, matsuda_future}. \\item We have considered the production of axionic strings and domain walls. In our case, the defects are not produced by the tachyon condensation but are formed by the spatial deformations of the mother branes. The parameter of the deformation is the position of the branes in the compactified space. These defects are first constructed in the classical brane configuration, and then lifted to MQCD. We have shown that quantum effect is crucial for axionic strings and domain walls. On the other hand, since the defects are  the non-trivial excitations of the system, they are also important in examining the consistency between SQCD and MQCD\\cite{Witten_Wall}. \\item We have discussed brane Q-balls in the Hanany-Witten brane dynamics. We have found that there is a distinguishable difference between brane Q-balls and conventional Q-balls. \\end{itemize}  In our future work\\cite{matsuda_future}, it is shown that monopoles and domain walls can be produced by daughter brane creations, which are extended between mother branes. For the cosmic strings, in our next paper\\cite{matsuda_angleddefect}, the consideration of the extended daughter brane is used to solve the long-standing problem of the $\\theta$-dependence. {\\bf In either case, the production of the extended brane is crucial for the cosmological defect formation}. We show in ref.\\cite{matsuda_future} that the cosmological process that is required for the creation of the monopoles and the domain walls is the same as the one required for the formation of incidental brane defects. {\\bf Then, the actual cosmological relics after brane inflation are the mixture of the two kinds.} The cosmological evolution of such brane defects is quite interesting and deserves further discussions."
    },
    "0402/astro-ph0402330_arXiv.txt": {
        "abstract": "{We compare performances of ground-based single-mode and multimode (speckle) interferometers in the presence of partial Adaptive Optics correction of atmospheric turbulence. It is first shown that for compact sources (i.e. sources smaller than the Airy disk of a single telescope) not entirely resolved by the interferometer, the remarkable property of spatial filtering of single-mode waveguides coupled with AO correction significantly reduces the speckle noise which arises from residual wavefront corrugations. Focusing on those sources, and in the light of the AMBER experiment (the near infrared instrument of the VLTI), we show that single-mode interferometry produces a better Signal-to-Noise Ratio on the visibility than speckle interferometry. This is true for bright sources ($K < 5$), and in any case as soon as Strehl ratio of $0.2$ is achieved. Finally, the fiber estimator is much more robust -- by two orders of magnitude -- than the speckle estimator with respect to  Strehl ratio variations during the calibration procedure. The present analysis theoretically explains why interferometry with fibers can produce visibility measurements with a very high precision, $1\\%$ or less. ",
        "introduction": "The great interest of using spatial filtering properties of optical waveguides in astronomical interferometers has been proven in the past years (\\cite{coude_du_foresto_etal_1}, \\cite{berger_1}). As a consequence, integrated optics and fibers are more and more introduced in the design of present and future interferometers to carry the signal to the detector. Furthermore, practical and theoretical studies (\\cite{haguenauer_1}, \\cite{guyon_1},  \\cite{mege_3}) have been undertaken to investigate the physical and optical properties in waveguided interferometers.  The present work aims at comparing the sensitivity and robustness of single-mode and multimode (speckle) interferometry. In Section \\ref{sec_modvis}, we recall the basic concepts of the fibered interferometric equation and the modal visibility. We derive, in Section \\ref{sec_error_vis}, the formal expression of the Signal to Noise Ratio (SNR) of the modal visibility which takes into account photon, detector and atmospheric noise. In Section \\ref{sec_pseudo_speck}, we propose an analytical approach to estimate the profile of the visibility SNR as a function of the magnitude, from partially Adaptive Optics (AO) corrected interferograms. We also derive the performance of single-mode interferometry applied to the AMBER experiment (the near infrared instrument of the VLTI), in the case of single Gaussian sources. Finally, in Section \\ref{sec_multi}, we compute the performances of the multispeckle method (\\cite{roddier_lena_1}, \\cite{mourard_etal_1}) currently used to estimate visibility from non-fibered interferometers and we compare the performances and the robustness of single-mode and speckle interferometry. ",
        "conclusions": ""
    },
    "0402/astro-ph0402106_arXiv.txt": {
        "abstract": "In the present paper some consequences of the assumption that in the center of the Galaxy there is a supermassive compact object without the events horizon are considered. The possibility of existence of such object has been argued earlier. It is shown, that accretion of a surrounding gas onto the object can cause nuclear burning in a superficial layer which owing to comptonization in a hotter layer, laying above, can manifest itself in observable IR and X spectra. The contribution of an intrinsic magnetic moment of the object in the observable synchrotron radiation is considered, using transfer equations, taking into account influence of gravitation on the energy and movement of  photons. ",
        "introduction": "An analysis of stars motion in the dynamic center of the Galaxy give strong evidence for the existence of a compact object with mass about $3\\cdot 10^{6}M_{\\odot}$ or more that is associated with  Sgr A* \\cite{gensel1},\\cite{gensel2},\\cite{ghes1}, \\cite{ghes2}. There are three kinds of an explanation of observable peculiarities of the object radiation:  1 - The gas accretion onto the central object -- a supermassive black hole (BH) \\cite{melia},\\cite{narayan},  2 - An ejection of the magnetized plasma from the vicinity of the Schwarzschild radius of the BH \\cite{falcke}, \\cite{melia1}  3 - Explanations based on hypotheses about another nature of the central object (a cluster of dark objects \\cite{maoz}  ,  a fermion ball hypothesis \\cite{viollier}, boson stars  \\cite{torres}.  In the present paper we consider some consequence of the assumption that  radiation of Sgr A* is caused by existence of a supermassive compact object without events horizon in the Galaxy Center. Such steady configurations of the degenerated Fermi-gas with masses $10^{2}\\div10^{10}$ $M_{\\odot}$ and with the radiuses $R$ less than the Schwarzschild radius $r_{g}$ are one of the consequence of the metric-field equations of gravitation \\cite{verozub95}, \\cite{verkoch01}, \\cite{verozub01}. In the theory  gravitational field of an attractive mass manifests itself as a field in Minkowski space-time for a remote observer in an inertial frame of reference, and as space-time curvature for the observer in a comoving (with the free falling particles) frame of reference. Physical consequences from the gravitation equations under consideration are very close to the ones in general relativity at the distances from the central mass much more than $r_{g}$ . However they are completely different at the distances nearby $r_{g}$ or less than that. The spherically-symmetric solution of the gravitation equations have no the event horizon and physical singularity in the center \\cite{verozub91}.  Since the gravitational equations was tested by the binary pulsar PSR 1913+16 \\cite{verkoch00} and stability of the supermassive configurations was studied sufficiently rigorously \\cite{verkoch01}, it is meaningful to investigate the possibility of the existence of such objects at the Galaxy Center as an alternative to the supermassive black hole hypothesis.  The gravitational force of a point mass $M$ affecting a free falling particle  of mass $m$ is given by \\cite{verozub91}. \\begin{equation} F=-m\\left[  c^{2}B_{1}+(B_{2}-2B_{3})\\overset{\\cdot}{r}^{2}\\right]  , \\label{gravaccel1}% \\end{equation} where \\begin{equation} B_{1}=C^{\\prime}/2A,\\ B_{2}=A^{\\prime}/2A,\\ B_{3}=C^{\\prime}/2C \\end{equation} and% \\begin{equation} A=f^{\\prime2}/C,\\ C=1-r_{g}/f,\\ f=(r_{g}^{3}+r^{3})^{1/3}. \\label{ABC} \\end{equation} In this equation $r$ is the radial distance from the center, $r_{g}% =2GM/c^{2},$ $M$ is the mass of the object, $G$ is the gravitational constant, $c$ is speed of light, the prime denotes the derivative with respect to $r$. The force affecting the test particle in rest is%  \\begin{equation} F=-\\frac{GmM}{r^{2}}\\left[  1-\\frac{r_{g}}{(r^{3}+r_{g}^{3})^{1/3}}\\right] \\label{ForceStat}% \\end{equation}   Fig. \\ref{GRForce} shows the force $F$ (in arbitrary units ) affecting the test particle at rest (curve 1) and the free falling particle (curve 2) as the function of the distance $\\overline{r}=r/r_{g}$ from the center.  \\begin{figure} \\resizebox{\\hsize}{!} {\\includegraphics[]{Force.eps}} \\caption{The gravitational force (arbitrary units) affecting a test particle at rest (curve 1) and free falling particle (curve 2) near the point attractive mass $M$.} \\label{GRForce} \\end{figure}    It follows from the above plot that the gravitational force affecting free falling particles decreases when $r$ approach to $r_{g}$ and changes its sign at $\\thicksim1.5$ $r_{g}$ . Although we never observed the motion of the particle at distances close to $r_{g}$ , we can test this conclusion for very distant objects in our Universe because for its observed mass $M_{u}$ the value of $2GM_{u}/c^{2}$ is close to the observed radius $r_{u}$ of the Universe. At such distances the gravitational force affecting the particles change the sign. And the repulsion force in a simple model of the expanded selfgraviting dust ball gives a simple and clear explanation of the acceleration of the Universe expansion \\cite{verozubA02}, \\cite{verozubB02}, %  It is seems in the first sight that the accretion onto the object give rise to a too large energy release at the surface that contradicts the low bolometric luminosity ( $\\thicksim10^{36}$ $erg\\ s^{-1}$ ) of Sgr A*. However, it must be taken into account that in the gravitation theory under consideration the velocity of free falling test particles decrease inside the Schwarzschild radius \\cite{verozub91}. If we assume that the the radius $R$ of the object with mass of $2.6\\cdot M_{\\bigodot}$ in the Galactic Center is equal to $0.04$ $r_{g}$ ( which follows from the solution of the equation of the hydrostatic  equilibrium \\cite{verozub95},\\cite{verkoch01}, \\cite{verozub01} ), then the value of the velocity $v$ of  free falling particles at the surface is $4\\cdot10^{8}$ $cm\\ s^{-2}$ . Therefore, even at the accretion rate $\\overset{\\cdot}{M}=10^{-6}$ $M_{\\odot}\\ yr^{-1}$ the amount of the released energy $\\overset{\\cdot}{M}v^{2}/2$\\  is only $\\thicksim10^{36}$ $erg/s$ . ",
        "conclusions": "The assignment of  nature of compact objects in the galactic  centers is one of basic problems of fundamental physics and astrophysics. The results received above, certainly, yet do not allow to draw the certain conclusions. However they show that the opportunity investigated here does not contradicts the observant data, and, therefore, demands the further study."
    },
    "0402/astro-ph0402276_arXiv.txt": {
        "abstract": "We use the deep wide-field optical imaging data of the Subaru/XMM-Newton Deep Survey (SXDS) to discuss the luminosity (mass) dependent galaxy colours down to $z'$=25.0 (5$\\times$10$^9$$h_{70}^{-2}$M$_{\\odot}$) for $z\\sim1$ galaxies in colour-selected high density regions. We find an apparent absence of galaxies on the red colour--magnitude sequence below $z'\\sim24.2$, corresponding to $\\sim$$M^*$+2 ($\\sim$10$^{10}$M$_{\\odot}$) with respect to passively evolving galaxies at $z\\sim1$.  Galaxies brighter than $M^*$$-$0.5 (8$\\times$10$^{10}$M$_{\\odot}$), however, are predominantly red passively evolving systems, with few blue star forming galaxies at these magnitudes.  This apparent age gradient, where massive galaxies are dominated by old stellar populations while less massive galaxies have more extended star formation histories, supports the `down-sizing' idea where the mass of galaxies hosting star formation decreases as the Universe ages.  Combined with the lack of evolution in the shape of the stellar mass function for massive galaxies since at least $z\\sim1$, it appears that galaxy formation processes (both star formation and mass assembly) should have occurred in an accelerated way in massive systems in high density regions, while these processes should have been slower in smaller systems.  This result provides an interesting challenge for modern CDM-based galaxy formation theories which predict later formation epochs of massive systems, commonly referred to as ``bottom-up''. ",
        "introduction": "Galaxy properties depend strongly on the mass of the system.  Based on the 122,808 galaxies drawn from the {\\it Sloan Digital Sky Survey (SDSS)}, Kauffmann et al. (2003) have recently shown an interesting bimodality of local galaxy properties separated at a stellar mass of $\\sim$3$\\times$10$^{10}$\\msun.  In particular, in contrast to massive galaxies which are dominated by old stellar populations showing little sign of recent star formation, less massive galaxies have a much larger contribution from young stars and a significant fraction of these low mass galaxies are likely to have experienced recent starbursts (see also Gavazzi \\& Scodeggio 1996; Baldry et al.\\ 2004). The morphological signatures also depend on the mass or luminosity, with massive galaxies showing centrally-concentrated light profiles (early-type/bulge morphologies), and less massive galaxies showing less concentrated profiles (late-type/disk morphologies; e.g., Kauffmann et al.\\ 2003; Treu et al.\\ 2003).  Based on the $z\\gsim1$ galaxies in the {\\it Hawaii Deep Fields}, Cowie et al.\\ (1996) have suggested that the most massive galaxies form earliest in the Universe, and star formation activity is progressively shifted to smaller systems, although their data are limited to galaxies brighter than $M^*$.  They have termed this phenomenon `down-sizing' in star forming galaxies. This apparent age gradient as a function of mass, where the massive galaxies are uniformly old while the less massive galaxies tend to be younger, appears to be at odds with cold dark matter (CDM) models of the Universe. In these models, galaxies form in a `bottom-up' or hierarchical manner, with small systems collapsing first and massive galaxies forming later via the assembly of these small systems.  Motivated by this interesting and fundamental puzzle in galaxy formation, we have chosen to investigate the mass dependence of galaxy properties at high redshift ($z\\sim1$) in more detail by studying much lower mass systems than previous work. To do this, we have used a statistically large sample of galaxies drawn from the Subaru/XMM-Newton Deep Survey (Sekiguchi et al.\\ 2004), which has the unique advantage of depth ($z'$=25.2 at 5--8$\\sigma$) and width (1.2 deg$^2$) achieved via the wide-field (30') camera Suprime-Cam on the 8.2-m Subaru Telescope.  Bell et al.\\ (2004) have recently analysed the redshift-dependent colour distributions of $\\sim$25000 galaxies over 0.78 deg$^2$ based on the COMBO-17 optical survey (Wolf et al.\\ 2003). This survey, however, is $\\sim$3 magnitudes shallower than the SXDS, reaching only to $\\sim$$M^*$ at $z\\sim1$. Our data therefore provides the first opportunity to study the properties of faint galaxies at high redshift.  We adopt the cosmological parameters ($H_0$, $\\Omega_m$, $\\Omega_{\\Lambda}$)=(70, 0.3, 0.7) throughout this paper, and define $h_{70}$ as $H_0$/(70~km s$^{-1}$Mpc$^{-1}$).  With this parameter set, 1 arcmin corresponds to 0.48~Mpc at $z\\sim1$.  All the magnitudes in this paper will be given in the AB-magnitude system.  We structure the paper as follows.  After an introduction in \\S1, we briefly describe in \\S2 the Subaru imaging data upon which the following analyses are based. In \\S3, we identify the high density regions at $z\\sim1$ by using the red sequence colour slice technique and determine the galaxy demographics in these regions by subtracting off foreground and background contaminations in a statistical manner. We investigate the photometric properties of this statistical sample of $z\\sim1$ galaxies in \\S4, with particular emphasis on the faint end. A discussion of our results and conclusions are given in \\S5 and \\S6, respectively. ",
        "conclusions": "In this paper, we have presented a photometric analysis of galaxies in colour-selected high density regions at $z\\sim1$, constructed from the unique deep ($z'$=25, 6--10$\\sigma$) and wide (1.2 deg$^2$) optical multi-colour imaging data taken as a part of the Subaru/XMM-Newton Deep Survey (SXDS) project. Our analysis has been based primarily on the field-corrected colour--magnitude diagram and the colour-dependent stellar mass functions.  Benefitted by the depth of the survey, we have found a deficit of faint red galaxies along the colour--magnitude sequence below $M^*$+2 with respect to the passive evolution, or $\\sim$10$^{10}$M$_{\\odot}$.  Almost all galaxies below this luminosity/mass at $z\\sim1$ are still undergoing significant star formation.  The luminous/massive end ($<$$M^*$$-$0.5), however, is dominated by old red systems with almost no blue galaxies.  The clear distinction of the `red + bright' and `blue + faint' populations at $z\\sim1$ on the colour--magnitude diagram suggests the existence of `down-sizing' in galaxy formation, where star formation is switched off from the massive systems towards the less massive ones as the Universe ages.  We find that the mass assembly process is also largely complete by $z\\sim1$ in the high density regions.  How to accommodate this down-sizing phenomena in galaxy formation within the context of the bottom-up scenario of the CDM Universe is an interesting puzzle."
    },
    "0402/astro-ph0402040_arXiv.txt": {
        "abstract": "{Cal 87 was observed with with {\\sl XMM-Newton} in April of 2003. The source shows a rich emission spectrum, where lines can be identified if they are red-shifted by 700-1200 km s$^{-1}$. These lines seem to have been emitted in a wind from the system. The eclipse is observed to be shifted in phase  by 0.03 $\\phi_{\\rm orb}$, where  $\\phi_{\\rm orb}$ is the phase of the optical light curve.}     \\addkeyword{Stars: binaries, white dwarfs} \\addkeyword{X-rays: stars}   \\begin{document} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0402/gr-qc0402007_arXiv.txt": {
        "abstract": "s{This talk reviews the constraints imposed by binary-pulsar data on gravity theories, focusing on ``tensor-scalar'' ones which are the best motivated alternatives to general relativity. We recall that binary-pulsar tests are qualitatively different from solar-system experiments, because of nonperturbative strong-field effects which can occur in compact objects like neutron stars, and because one can observe the effect of gravitational radiation damping. Some theories which are strictly indistinguishable from general relativity in the solar system are ruled out by binary-pulsar observations. During the last months, several impressive new experimental data have been published. Today, the most constraining binary pulsar is no longer the celebrated (Hulse-Taylor) PSR B1913$+$16, but the neutron star-white dwarf system PSR J1141$-$6545. In particular, in a region of the ``theory space'', solar-system tests were known to give the tightest constraints; PSR J1141$-$6545 is now almost as powerful. We also comment on the possible scalar-field effects for the detection of gravitational waves with future interferometers. The presence of a scalar partner to the graviton might be detectable with the LISA space experiment, but we already know that it would have a negligible effect for LIGO and VIRGO, so that the general relativistic wave templates can be used securely for these ground interferometers.} ",
        "introduction": "The most efficient way to test a theory is to contrast its predictions with alternative models. Instead of just confirming or ruling out a particular theory, this method allows us to understand what features have actually been tested, and what kind of observations could be performed to test the remaining features.  The best known example of such an embedding of general relativity in a space of alternatives is the so-called ``parametrized post-Newtonian'' (PPN) formalism,\\cite{w93} which describes all possible metric theories of gravity in weak-field conditions, at order $1/c^2$ with respect to the Newtonian interaction. The basic idea was formulated by Eddington,\\cite{edd} who wrote the usual Schwarzschild metric in isotropic coordinates, but introduced some phenomenological parameters $\\beta^\\text{PPN}$ and $\\gamma^\\text{PPN}$ in front of the different powers of the dimensionless ratio $Gm/rc^2$: \\begin{subequations} \\label{1} \\begin{eqnarray} -g_{00}&=& 1 - 2\\frac{Gm}{rc^2} + 2 \\beta^\\text{PPN} \\left(\\frac{Gm}{rc^2}\\right)^2 + \\mathcal{O}\\left(\\frac{1}{c^6}\\right)\\,, \\label{1a}\\\\ g_{ij}&=&\\delta_{ij}\\left(1+2\\gamma^\\text{PPN}\\frac{Gm}{rc^2}\\right) + \\mathcal{O}\\left(\\frac{1}{c^4}\\right)\\,. \\label{1b} \\end{eqnarray} \\end{subequations} General relativity corresponds to $\\beta^\\text{PPN} = \\gamma^\\text{PPN} = 1$, and is in perfect agreement with solar system experiments. At the time of the 10th Marcel Grossmann Meeting (MGX), the tightest published bounds on these parameters were\\cite{w01} \\begin{equation} |\\gamma^\\text{PPN}-1| < 2\\times 10^{-3}\\,, \\qquad |\\beta^\\text{PPN}-1| < 6\\times 10^{-4}\\,. \\label{2} \\end{equation} An unpublished\\cite{eubanks} stronger constraint was also known, $|\\gamma^\\text{PPN}-1| < 4\\times 10^{-4}$, but an impressive new result\\cite{cassini} has been released two months after MGX: \\begin{equation} \\gamma^\\text{PPN} - 1 = (2.1\\pm2.3)\\times 10^{-5}\\,. \\label{3} \\end{equation} Such bounds tell us that general relativity is basically the only theory consistent with experiment at the first post-Newtonian order. However, they do not constrain higher order terms in metric~(\\ref{1}), and the correct theory of gravity might differ significantly from general relativity in strong field conditions. Indeed, if $R$ denotes the radius of a body, the ratio $Gm/Rc^2 \\approx 10^{-9}$ for the Earth and $\\approx 10^{-6}$ for the Sun, but it reaches $\\approx 0.2$ for neutron stars, not far from the theoretical maximum of $\\frac{1}{2}$ for black holes. Pulsar observations can thus be used to test the strong-field regime of gravity. ",
        "conclusions": "Binary pulsars are ideal tools for testing relativistic gravity in the strong field regime. In the most natural class of alternatives to general relativity, tensor-scalar theories, their probing power has been shown to be qualitatively different from weak-field experiments: They have the capability of testing theories which are strictly equivalent to general relativity in the solar system.  Two fantastic binary pulsars have been timed recently. The double pulsar J0737$-$3039 promises to be the best laboratory for testing general relativity itself, and for studying the physics of pulsars. On the other hand, the neutron star-white dwarf system PSR J1141$-$6545 is by far the most constraining binary pulsar known at present, because its asymmetry implies that it generically emits strong dipolar gravitational waves in scalar-tensor theories. It is already almost as constraining as solar-system tests even in the region of positive $\\beta_0$'s, where binary pulsars never competed with weak-field experiments up to now. It should probe values of the Eddington parameter $|\\gamma^\\text{PPN}-1|\\sim 10^{-6}$ by the end of the decade, \\textit{i.e.}, more than one order of magnitude better than present solar-system limits.  Binary pulsars are so precise that they already exclude the models which would have predicted some significant scalar-field contributions to the gravitational wave templates necessary for LIGO and VIRGO. Therefore, one may use securely the general relativistic templates for these interferometers. It is still possible that such scalar-field effects be detectable with the LISA space interferometer, but binary pulsars will probably give us tighter bounds before it is launched."
    },
    "0402/astro-ph0402510_arXiv.txt": {
        "abstract": "We investigate a class of Cardassian scenarios of the universe evolution in notions of the qualitative theory of dynamical systems. This theory allows us to analyze all solutions for all possible initial conditions on the phase plane. In the Cardassian models we find that big-rip singularities are present as a typical behavior in the future if $n<0$. Some exact solutions for the flat Cardassian models as well as a duality relation were found. In turn from the statistical analysis of Knop's SNIa data, without any priors on matter content in the model, we obtain that at the $99\\%$ confidence level this big-rip scenario will reach. The potential function for the Hamiltonian description of dynamics is reconstructed from the SNIa data (inverse dynamical problem). We also pointed out the statistical analysis results depend oversensitively on the choice of the model parameter $\\mathcal{M}$. ",
        "introduction": "In this note we apply the qualitative analysis of differential equations to the Cardassian models which have become popular during last two years due to some of their interesting features. This class of models was proposed as an alternative to the cosmological constant model \\cite{freese02,wang03} to explain the current acceleration of the Universe \\cite{perlmutter99,riess98}. Freese and Lewis \\cite{freese02} claimed that the Cardassian model, in which the Friedmann-Robertson-Walker (FRW) equation is modified, explains the expansion of the universe without any dark energy component. A certain additional term in the FRW equation which may arise from the fundamental physics (brane epoch) drives the present acceleration of the Universe.  The Cardassian model survives several observational tests like the magnitude-redshift for the present type Ia supernovae data \\cite{zhu03,zhu03b,avelino03,wang03,szydlowski03c,godlowski03}, $\\theta$--$z$ test of the angular size of high-$z$ compact radio sources \\cite{zhu02} or $d_{A}$--$z$ test for the SZ/X-ray clusters proposed by Zhu and Fujimoto \\cite{zhu03a}.  The main aim of this paper is to demonstrate the theoretical power in explanation of the cosmological problems by the Cardassian model. Moreover, the dynamical system methods allow to reveal some unexpected properties of this model. We construct the phase spaces for these models and discuss how their structure differs from the canonical model with the cosmological constant (Do new types of solutions appear? What is their physical meaning? Does any change of parameter values lead to some change of dynamical behavior?).  From the theoretical point of view it is interesting to analyze all evolutional paths of the Cardassian models for all initial conditions. On the other (empirical) hand it is important to check which Cardassian evolutional paths fit to astronomical observations, namely type Ia supernovae data. The both issues will be examined in this paper.  There is a widespread opinion that physically realistic models of the universe should possess some kind of structural stability---the existence of too many dramatically different mathematical models which agree with observations would be fatal for the empirical method of science \\cite{thom77}. A dynamical system is said to be structurally stable if other dynamical systems which are close to it (in a metric sense) are topologically equivalent (i.e. modulo homeomorphism). Although the question how to ensure such stability is an open problem for the dynamical systems of higher dimension than two \\cite{smale80} in the case of two-dimensional dynamical systems following the Peixoto theorem that states the structurally stable systems form open and dense subsets in the space of all dynamical systems on the plane \\cite{peixoto62}. Moreover, in the case of two-dimensional systems there is a simple test of structural stability. Namely, if the right-hand sides of the dynamical systems are in the polynomial form, the global phase portraits are structurally stable if the number of critical points and limit cycles is finite, each point is hyperbolic and there are no trajectories connecting saddle points.  In the case considered the dynamics is reduced to the form of a two dimensional dynamical system with right-hand sides in the polynomial form \\begin{align} \\dot{x} \\equiv \\frac{dx}{dt} = P(x,y),\\\\ \\dot{y} \\equiv \\frac{dy}{dt} = Q(x,y), \\label{eq:1} \\end{align} where $P, Q \\in C^{\\infty}$ class of functions; $(x,y)$ is a differential space $\\mathcal{M}$, which is useful for visualization of the dynamics, called the phase space (or state space). The right-hand sides define a vector field $\\mathcal{V}=(P,Q)$ belonging to the tangent bundle $\\mathcal{M}$. Now we can define a phase curve as a integral curve of the vector field. All phase curve with the critical points $(P(x_{0},y_{0})=0$, $Q(x_{0},y_{0})=0)$ constitute the phase portrait of the system. Two phase portraits are equivalent if there exists an orientation preserving a homeomorphism which maps integral curves of both systems into each other.  From the physical point of view a critical point represents asymptotic states or equilibria. The main aim of qualitative analysis of differential equations is constructing the phase portrait of the system. Following the Hartman-Grobman theorem: the nonlinear dynamics near the hyperbolic critical points ($\\forall i$ ${\\rm Re} \\lambda_{i} \\neq 0$, where $\\lambda_{i}$ is an eigenvalue of a linearization matrix) is qualitatively equivalent to its linear part \\begin{align} \\label{eq:2} \\dot{x}&=\\frac{\\partial P}{\\partial x}(x_{0},y_{0})(x-x_{0}) +\\frac{\\partial P}{\\partial y}(x_{0},y_{0})(y-y_{0}) \\\\ \\label{eq:3} \\dot{y}&=\\frac{\\partial Q}{\\partial x}(x_{0},y_{0})(x-x_{0}) +\\frac{\\partial Q}{\\partial y}(x_{0},y_{0})(y-y_{0}). \\end{align}  Full knowledge of the dynamical system comprises also its behavior at infinity. To achieve this one usually transforms the phase plane into a Poincar\\'e sphere. Then infinitely distant points of the plane are mapped into the equator of the sphere. Of course, the character of critical points (which depends on the solutions of characteristic equation $\\lambda^{2}-\\lambda{\\rm Tr}\\mathcal{A}+\\det\\mathcal{A}=0$ for linearization matrix $\\mathcal{A}$) is conserved but new critical points can appear at the equator. Now, the orthogonal projection of any hemisphere onto the tangent plane gives the compactified portrait on compact projective plane.  There are two main aims of the paper. First, theoretical investigations of the dynamics and second, its reconstruction from Knop's SNIa sample (inverse dynamical problem). We found that the Cardassian model very well fits the SNIa data and that its unexpected future (the big-rip singularity) is consistent with SNIa data on the $99\\%$ confidence level without any priors on matter content. We have demonstrated how some information about dark energy can be deduced from the potential of the Hamiltonian system describing the evolution of the universe. Using the potential function instead of the equation of state parameter $w=p/\\rho$ seems to be attractive in the context of searching the adequate description of the present stage of evolution of the Universe. ",
        "conclusions": "In this paper we apply the qualitative cosmology analysis to the general class of Cardassian models which have become rather an alternative to the dark energy models. We show that while the application of qualitative methods allows to reveal some unexpected properties like the big-rip singularities, the SNIa data uniquely reconstruct the phase plane in the model independent way. The reconstruction of this phase plane should be treated as a necessary condition for any model which wants to explain the present acceleration of the universe.  Theory of dynamical systems which offer a possibility of investigating the space of all solutions for all admissible initial conditions is used in analysis of Cardassian models. We demonstrate a simple method of reducing of the dynamics to certain two-dimensional dynamical system. One of the features of this reduction is the possibility of representing the model as a Hamiltonian system in which the properties of the potential function can serve as a tool for qualitative classification of all possible evolution scenarios. It is shown that some important features like resolution of the flatness problem or horizon problem and its degree of generality can be visualized as domains on the phase plane. Then one is able to see how large is the class of solutions (labelled by the initial conditions) leading to the desired property, i.e., this class is generic or non-typical.  Applying the techniques of the particle-like description, developed by us earlier, to a new data set from the Knop's sample we show that the Cardassian model with $n<0$ is favored by the data on the $99\\%$ confidence level. On the other hand it is just the case of appearance of unexpected big-rip singularities in the future evolution of the model.  Although the developed formalism is mainly adopted to the investigation of the Cardassian models it can be useful to include additionally dissipative effects of bulk viscosity to the model. Moreover for these class of models de Sitter solution is admissible as a global attractor in the future.  The main results are the following:  \\begin{enumerate} \\item We show the effectiveness of using of dynamical system methods in analysis of the wide class of Cardassian models and their generalizations. \\item The unwanted big-rip singularities at a finite time were detected as a generic for the case $n<0$ (a global attractor in the future). \\item The exact solutions and some duality relations were found. \\item Additionally we showed that the existence of this kind of singularities is acceptable on the $99\\%$ confidence level from Knop's sample of SNIa data. \\item The particle-like approach was adopted to describe the dynamics as well as the potential function was reconstructed from the observations. In the fitting procedure the power law function are used because it gives the lowest value of $\\chi^{2}$. \\item Because of interpretation of the Cardassian term as an effect of additional fluid one can claim that if FRW equation is ``correct'' then phantom fields are required at the $99\\%$ confidence level. \\end{enumerate}  Finally, our study showed that the Cardassian model with $n<0$ is strongly favored by the Knop's SNIa data (without any priors on matter content). But, this model predicts rather unexpected future of the universe -- a big-rip singularity. However, the estimation of characteristic time of the big-rip singularity will be a subject of our next paper."
    },
    "0402/astro-ph0402660_arXiv.txt": {
        "abstract": "We report on a long ($100$~ks) \\xmm\\ observation of the bright Seyfert 1 galaxy Arakelian 120.  The source previously showed no signs of intrinsic reddening in its infrared--ultraviolet continuum and previous observations had shown no evidence for ionized absorption in either the ultraviolet or X-ray bands.  The new \\xmm\\ RGS data place tight limits on the presence of an ionized X-ray absorber and confirm that the X-ray spectrum of \\ark\\ is essentially unmodified by intervening matter. Thus \\ark\\ can be considered a `bare' Seyfert 1 nucleus. This observation therefore offers a  clean view of the X-ray spectrum of a `normal' Seyfert galaxy free from absorption effects. The spectrum shows a Doppler broadened iron emission  line ($FWHM\\sim 3\\times 10^4$~km s$^{-1}$) and a smooth, continuous soft excess which appears to peak at an energy $\\approx 0.5$~keV.  This adds weight to the claim that genuine soft excesses  (i.e. those due to a real steepening of the underlying continuum below $\\sim 2$~keV) are ubiquitous in Seyfert 1 spectra.  However, the detailed shape of the excess could not be reproduced by any of the simple models tested (power-laws, blackbodies, Comptonised blackbodies, accretion disc reflection).  This observation therefore demonstrates both the need to understand the soft excess (as a significant contributor to the luminosity of most Seyfert 1s) and the inability of the existing, simple models to explain it. ",
        "introduction": "\\label{sect:intro}  Arakelian 120 (aka Mrk 1095) is a luminous Seyfert 1 galaxy at a redshift $z=0.0323$. It was the subject of an early attempt at reverberation mapping, and was important in demonstrating the compact size of the optical broad line region (Peterson \\et 1985; Peterson \\& Gaskell 1991).  More recent optical monitoring campaigns have yielded an estimate of the mass of the central black hole of $\\sim 2 \\times 10^{8}$~\\Msun (Wandel, Peterson \\& Malkan 1999).  The bolometric luminosity for the nucleus is $L_{\\rm bol} \\gs 10^{45}$ erg s$^{-1}$ (Edelson \\& Malkan 1996 estimated the total $0.1-100~\\mu{\\rm m}$ luminosity to be $\\approx 8\\times 10^{44}$  erg s$^{-1}$).  This would suggest it is radiating at $L/L_{\\rm Edd} \\gs 0.05$, where $L_{\\rm Edd}$ is the Eddington luminosity for a  $2 \\times10^{8}$~\\Msun\\ black hole. The nucleus is radio-quiet  but shows a slight extension in its radio image (Condon \\et 1998; Ho 2002). Ward \\et (1987) used broad-band photometry to identify \\ark\\ as a `bare' Seyfert nucleus, i.e. one free from significant reddening or contamination from the host galaxy.  The host is a low-inclination spiral galaxy (Hubble type S0/a, inclination $i \\approx 26\\deg$; Nordgren \\et 1995) .  \\ark\\ has been observed with most of the major X-ray observatories. An \\exosat\\ observation showed \\ark\\ to have a steep soft X-ray spectrum (Turner \\& Pounds 1989), as did a subsequent \\rosat\\ observation (Brandt \\et 1993). Furthermore, these X-ray observations showed no indication of any `warm absorption' features -- i.e. discrete absorption features often found in the soft X-ray band and caused by absorption in photoionized gas along the line-of-sight to the nucleus. Warm absorption systems are common in Seyfert 1s (Reynolds 1997; Crenshaw, Kraemer \\& George 2003).  Observations in the ultraviolet (Crenshaw \\et 1999; Crenshaw \\& Kraemer 2001) showed \\ark\\ to be one of the Seyfert 1 galaxies that showed no intrinsic ultraviolet absorption. Thus \\ark\\ is a rare example of a bright Seyfert 1 galaxy that is not significantly affected by any kind of complex absorption -- its emission spectrum is that of a `bare' Seyfert 1 nucleus.  The paper presents the results of a long \\xmm\\ observation designed to characterise the intrinsic X-ray emission spectrum of a Seyfert 1 galaxy. The rest of this paper is organised as follows. Section~\\ref{sect:obs} discusses the observation details and data reduction. Section~\\ref{sect:timing} gives details of the X-ray variability observed in \\ark\\ and  section~\\ref{sect:fluxed} describes a `first look' at the X-ray spectrum. The spectrum is then examined in detail using the Reflection Grating Spectrometer (RGS) in section~\\ref{sect:rgs}. This is followed by an analysis of the European Photon Imaging Camera (EPIC) data first over the $3-10$~keV band (section~\\ref{sect:epic}) and then the over the full band-pass (section~\\ref{sect:broad}). Section~\\ref{sect:xte}  briefly discusses an analysis of archival \\xte\\ observations of \\ark.  Finally, the results are discussed in section~\\ref{sect:disco}. ",
        "conclusions": "\\label{sect:disco}  \\subsection{The X-ray properties of \\ark}  This paper describes the results of a  long \\xmm\\ observation of the luminous Seyfert 1 galaxy Arakelian 120. The X-ray emission from the source was only weakly variable on the short timescales probed by \\xmm\\ (section~\\ref{sect:timing}) but archival \\xte\\ data revealed `normal' Seyfert 1 variability traits on longer timescales (section~\\ref{sect:xte}; cf. Edelson \\& Nandra 1999; Uttley \\et 2002; Markowitz \\et 2003b).  The X-ray spectrum showed a notable absence of warm absorption features (section~\\ref{sect:rgs}). The lack of absorption means that these data represent the `bare' X-ray emission spectrum of a fairly typical Seyfert 1.  The spectrum above $\\sim 3$~keV can be explained using a fairly conventional spectral model (sections~\\ref{sect:iron} and \\ref{sect:refln}) comprising a power-law continuum ($\\Gamma \\approx 2$) plus Doppler broadened emission from the surface of a weakly ionized reflector (with relative reflection strength $R \\approx 0.5$). The \\xte\\ data support this model (section~\\ref{sect:xte}).  The spectrum below $\\sim 3$~keV becomes dominated by a steep, smooth, broad soft excess component which appears to peak at $\\sim 0.5$~keV.  \\subsection{The iron line of \\ark}  The results of the spectral fitting indicated an iron emission line that is the composite of a weak, narrow line originating in distant material and a stronger line showing significant Doppler broadening.  A narrow, neutral iron line appears to be ubiquitous in Seyfert 1 galaxies (e.g. Yaqoob, George \\& Turner 2002; Page \\et 2004b; Yaqoob \\& Padmanabhan 2004) and may have an original in the optical broad-line region or the putative molecular torus (see discussion in Yaqoob \\et 2001; Reeves \\et 2004). The lack of neutral absorption intrinsic to \\ark\\ (section~\\ref{sect:rgs}) confined this material to lie out of the line of sight, which in turn implies the covering fraction of the line emitting material must be below unity. The relative strength of the narrow line ($EW \\sim 40$~eV) further suggests the sky covering fraction of the material is small ($f_{\\rm C} \\ls 0.5$; Nandra \\& George 1994) or the optical depth is small ($\\tau \\ls 0.1$; Leahy \\& Creighton 1993) or both.  The broad component to the line has a velocity width  $FWHM \\sim 3\\times 10^4$ km s$^{-1}$, far broader than the broad optical lines (e.g. $FWHM({\\rm H}\\beta) \\approx 5800$~km s$^{-1}$; Wandel \\et 1999). However, there is no requirement for the line emitting region to extend into the region of strong gravity about the black hole ($\\ls 20$~\\rg) which would produce an asymmetric redward tail  on the line profile (section~\\ref{sect:iron}). The best-fitting parameters of the disc line model are slightly unusual; the inclination angle is high ($i > 73 \\deg$), the inner radius is rather large ($r_{\\rm in} \\sim 140 r_{\\rm g}$) and the rest-frame energy is unusual ($E \\approx 6.56$~keV). One plausible origin for this line is a weakly ionized accretion disc. The fact that the inner radius is greater than $6r_{\\rm g}$ could mean the disc is truncated. However, the unusual energy of  the line (corresponding to mildly ionized iron) might allow for a simpler alternative. If the disc survives down to the ISCO but rapidly becomes ionized with decreasing radius then, depending on the detailed ionization structure, the innermost regions may produce little observable line emission (see e.g. Ross, Fabian \\& Young 1999).  However, as demonstrated by Fig.~\\ref{fig:contour} there is sufficient leverage in the fit to allow for a smaller inner radius ($r_{\\rm in} \\ls 100 r_{\\rm g}$), and also lower inclination angle, provided the emissivity law is quite flat ($q \\approx -2$). As mentioned in section~\\ref{sect:iron} a good fit can be obtained with $r_{\\rm in}=6\\rg$ and $i \\approx 30 \\deg $ provided that there are two co-existing Fe K$\\alpha$ emission lines arising in the disc (at energies of $6.4$ and $6.8$~keV) and no intrinsically narrow component. This may be feasible if the disc is clumpy or inhomogeneous (Ballantyne, Turner \\& Blaes 2004). The required flat emissivity can be produced if the disc is illuminated from a great height. Figure~\\ref{fig:emissivity} shows the effect of altering $r/h$ in a `lamppost' geometry [i.e. the disc was assumed to be flat and illuminated by a central point source raised by a height $h$, in which case the emissivity law is given by $J(r) \\propto h/(r^2+h^2)^{3/2}$]. This demonstrates that if the X-illumination is provided from a height $h \\sim 100 r_{\\rm g}$ the emissivity will be quite flat out to similarly large radii. Illumination from a large height would require a non-standard geometry for the X-ray emitting region, one possibility is the `aborted jet' model of Ghisellini, Harrdt \\& Matt (2004).  \\subsection{\\ark\\ as a `bare' Seyfert 1}  These new X-ray observations showed \\ark\\ to posses no evidence of an X-ray warm absorber and placed upper limits on the ionic column densities that are substantially  lower than those of more typical, absorbed Seyfert 1s. For example, NGC 3783 shows O~\\textsc{vii} and O~\\textsc{viii} absorption with column densities two orders of magnitude higher (Kaspi \\et 2002). Kaastra \\et (2000) observed NGC 5548 and measured C~\\textsc{vi}, N~\\textsc{vi}, O~\\textsc{vii} and O~\\textsc{viii} absorption lines with corresponding ionic column densities an order of magnitude higher than the limits obtained for \\ark.  A last counter example is IRAS 13349+2438 (Sako \\et 2001) which shows absorption lines  of C~\\textsc{vi}, N~\\textsc{vi}, O~\\textsc{vii}, O~\\textsc{viii}, Ne~\\textsc{ix} and Ne~\\textsc{x}, with column densities an order of magnitude higher than \\ark.  These X-ray observations therefore confirm that \\ark\\ is indeed a `bare' Seyfert 1 nucleus, as suspected based on its broad band spectral energy distribution (SED; Edelson \\& Malkan 1986; Ward \\et 1987) and the lack of ultraviolet absorption seen in \\hst/FOS spectrum (Crenshaw \\et 1999). Other ultraviolet observations have been made with the  Goddard High-Resolution Spectrograph (Penton, Stocke \\& Shull 2000) and with \\fuse\\ (Wakker \\et 2003). Using these data Wakker \\et (2003) showed that the line-of-sight through the Galaxy towards \\ark\\ possesses a very low column of O~\\textsc{vi}.  A preliminary examination of these data suggest there may be weak traces of ionized absorption intrinsic to \\ark.  \\begin{figure} \\centering \\includegraphics[width=6.0 cm, angle=270]{emissivity.ps} \\caption{ The local power-law slope of the emissivity function  ($J(r) \\sim r^q$) for a lamppost geometry. } \\label{fig:emissivity} \\end{figure}  \\subsection{Other bare Seyferts}  Shortly after the original discovery of the X-ray warm absorber (Halpern 1984), this phenomenon was found to be common in Seyfert 1 galaxies. Low resolution spectra from \\exosat\\ and \\ginga\\ suggested $\\sim 50$ per cent of Seyfert 1s showed evidence for warm absorption (Turner \\& Pounds 1989; Nandra \\& Pounds 1994).  Subsequent observations with the better spectral resolution afforded by \\asca\\ suggested the incidence of warm absorbers was possibly higher ($\\sim 50-70$ per cent; Reynolds 1997; George \\et 1998).  Similar results were found for the ultraviolet absorber based on  \\hst\\ FOS spectra (Crenshaw \\et 1999).  The increased sensitivity and resolution offered by \\xmm\\ and \\chandra\\ have allowed for even more sensitive searches for absorption. These have confirmed that there remains a significant population of Seyfert 1s that lack a strong X-ray warm absorber.  In addition to \\ark,  other Seyfert 1s that show a distinct lack of warm absorption  in both their \\hst\\ FOS ultraviolet and \\xmm/\\chandra\\ X-ray spectra include Mrk 478 (Marshall \\et  2003a),  Mrk 335 (Gondoin \\et 2002), Fairall 9 (Gondoin \\et 2001) and Mrk 205 (Reeves \\et 2001). All four of these were in the sample of Crenshaw \\et (1999)  and classified as having no intrinsic ultraviolet warm absorber.   Other objects known to lack X-ray absorbers include  PKS 0558--504 (O'Brien \\et 2001a), MCG--2-58-22 (Weaver \\et 1995) and Ton S180\\footnote{ There is some debate about Ton S180.  R\\'{o}\\.{z}a\\'{n}ska \\et (2004) re-examined the \\chandra\\ LETGS observation and claim to have found several resonance absorption lines.  The lines are very weak and in most cases of borderline significance (many have equivalent widths consistent with zero).  Furthermore, the spectral analysis was complicated by the presence of contaminant on the \\chandra\\ ACIS which affects the LETGS/ACIS calibration (see Marshall \\et 2003b).  Thus the existence of an X-ray warm absorber in Ton S180 is still not well established and in any case must be extremely weak.  } (Turner \\et 2001b; Vaughan \\et 2002).  These objects are probably not completely without ionized absorption systems. Kriss (2002) reported weak absorption by the O~\\textsc{vi} $\\lambda \\lambda 1032, 1038$~\\A\\ resonance doublet in  both Ton S180 and Mrk 478 based on high resolution \\fuse\\ data (see also Turner \\et 2002 for the \\fuse\\ observation of Ton S180). Such absorption is too weak to have been detected in the Crenshaw \\et (1999) ultraviolet survey. The corresponding X-ray warm absorbers in these objects may be present but so weak as to have a negligible effect on the available data.  The reason for the lack of an X-ray warm absorber is unclear.  The most obvious explanation is that overall column density  is lower in these objects (by at least an order of magnitude) compared to  more typical, absorbed Seyferts. However, it is also plausible that a similar column of ionized gas exists but is either too highly ionized to show significant spectral features or lies out of the line of sight (requiring a covering fraction less than unity).  A detailed survey comparing the X-ray and ultraviolet emission/absorption line spectra of a large sample of bright Seyferts may be able to answer this question (see discussions in Crenshaw \\et 2003). A more speculative solution, recently proposed by Gierli\\'{n}ski \\& Done (2004), is that objects that lack narrow absorption lines may be dominated by a deep, broad absorption trough produced by an absorption system with such high dispersion velocity that no individual lines can be resolved. This model would require the underlying power-law continuum to be rather steep, which then makes explaining the upturn above $\\sim 10$~keV more difficult to explain.  \\subsection{How common are soft X-ray excesses?}  These observations have clearly revealed a luminous soft excess in \\ark, in the sense that an extrapolation of the hard ($\\gs 3$~keV) continuum into the soft X-ray band revealed a very significant upturn. Steep soft X-ray spectra were first seen in Seyferts using \\heao\\ data (Pravdo \\et 1981)   and subsequently seen as excesses over the hard power-law  by \\exosat\\ (Arnaud \\et 1985) and \\einstein\\ (Bechtold \\et 1987).  Early soft X-ray surveys (e.g. the \\exosat\\ survey of Turner \\& Pounds 1989) suggested that $\\gs 50$ per cent of all Seyferts possessed a soft excess component.  The \\rosat\\ survey of Walter \\& Fink (1993) also suggested a high incidence of soft excesses.  These numbers were highly uncertain, however,  because many of the sample members were absorbed.  Complex soft X-ray absorption can mask, or in some cases even mimic, the steeper soft X-ray spectrum indicative of a soft excess.  The recent \\xmm\\ and \\chandra\\ observations would seem to indicate that all Seyfert 1s without strong X-ray warm absorption do show a strong soft excess.  However, many of the well-known examples (e.g. Ton S180, PKS 0558-504, Mrk 478, NGC 4051 and Mrk 359) are narrow-line Seyfert 1s (NLS1s; Osterbrock \\& Pogge 1985). NLS1s are a subclass of Seyfert 1s defined by their narrow permitted optical lines ($FWHM({\\rm H}\\beta) \\ls 2000$~km s$^{-1}$) but noted for their often exceptionally steep soft X-ray spectra (Boller \\et 1996; Laor \\et 1997; Vaughan \\et 1999; Leighly 1999). Until recently it was possible that the soft excess appeared ubiquitous in unabsorbed Seyferts only because many of the well-studied examples were NLS1s (i.e. the sample of unabsorbed Seyferts was biased towards the soft NLS1s).  \\ark\\ is an interesting counter example, being both unabsorbed and a `normal,' broad-line Seyfert 1 (BLS1) with  $FWHM({\\rm H}\\beta) \\approx 5800$~km s$^{-1}$ (Wandel \\et 1999). Yet it too shows a strong soft excess. Other notable BLS1s that also lack absorption include Mrk 205 (Reeves \\et 2001), Mrk 335  (Gondoin \\et 2002) and Fairall 9 (Gondoin \\et 2001),  all of which showed soft excesses in their \\xmm\\ observations.  Although a complete and unbiased survey of the soft X-ray spectra of Seyfert 1s has yet to be conducted, it does seem highly plausible that all unabsorbed Seyfert 1s (whether NLS1 or BLS1) possess a soft excess. Assuming that there is no fundamental difference in the underlying X-ray continuum spectra of absorbed and unabsorbed Seyfert 1s then implies that soft X-ray excesses are ubiquitous to Seyfert 1s. The mini-survey of \\xmm\\ spectra by Pounds \\& Reeves (2002) showed that in all six Seyfert 1s they studied, the soft X-ray spectrum ($\\ls 0.5$~keV) was always  $\\gs 50$ per cent higher than an extrapolation of the hard X-ray power-law would predict.  Furthermore, detailed studies of Seyfert 1s with complex, strong X-ray warm absorbers often conclude that an additional soft X-ray emission component is required behind the warm absorber (e.g. Collinge \\et 2001; Netzer \\et 2003;  Blustin \\et 2003).  This therefore underlines the need to understand the soft excess as a common (perhaps ubiquitous) and significant contributor to the luminosity of Seyferts.  In the case of \\ark\\ the best-fitting models are a doubly-broken power-law with an anomalously flat slope below $0.5$~keV, or multiple, soft blackbodies. The blackbody origin is difficult to explain as the temperatures are far too high to correspond to any standard accretion disc. The expected temperature for the inner region ($\\approx 6r_{\\rm g}$) of a geometrically thin, optically thick disc about a $2 \\times10^{8}$~\\Msun\\ black hole is $kT \\sim 11$~eV if radiating at $L/L_{\\rm Edd}=0.1$ and $kT \\sim 20$~eV if radiating at $L/L_{\\rm Edd}=1$. This emission should therefore not contribute to the observed X-ray spectrum. The temperatures of the best-fitting blackbody components were at at least an order of magnitude higher than this. In addition, the size of the emission region implied by these high temperatures is far too small ($ r_{\\rm BB} \\ll r_{\\rm g}$).  The alternative models tested (reflection, disc blackbodies and bremsstrahlung) all produced the wrong spectral shape.  At present there is no single model that can account for the known properties of the soft excess (see also  the discussion in Vaughan \\et 2002)."
    },
    "0402/astro-ph0402383_arXiv.txt": {
        "abstract": "The gas at the surfaces of molecular clouds in galaxies is heated and dissociated by photons from young stars both near and far.  \\HI\\ resulting from the dissociation of molecular hydrogen \\Htwo\\ emits hyperfine line emission at $21$ cm, and warmed CO emits dipole rotational lines such as the $2.6$ mm line of CO(1--0).  We use previously developed models for photodissociation regions (PDRs) to compute the intensities of these \\HI\\ and CO(1--0) lines as a function of the total volume density $n$ in the cloud and the far ultraviolet flux \\Gzero\\ incident upon it and present the results in units familiar to observers.  The intensities of these two lines behave differently with changing physical conditions in the PDR, and, taken together, the two lines can provide a ground--based radio astronomy diagnostic for determining $n$ and \\Gzero\\ separately in distant molecular clouds.  This diagnostic is particularly useful in the range \\Gzero\\ $\\lesssim 100$, 10 \\pcmcub $\\lesssim n \\lesssim 10^5$ \\pcmcub, which applies to a large fraction of the volume of the interstellar medium in galaxies.  If the molecular cloud is located near discrete sources of far--UV (FUV) emission, the PDR--generated \\HI\\ and CO(1--0) emission on the cloud surface can be more easily identified, appearing as layered ``blankets'' or ``blisters'' on the side of the cloud nearest to the FUV source.  As an illustration, we consider the Galactic object G216 -2.5, i.e.\\ ``Maddalena's Cloud'', which has been previously identified as a large PDR in the Galaxy.  We determine that this cloud has $n\\approx 200$ \\pcmcub\\ and \\Gzero\\ $\\approx 0.8$, consistent with other data. ",
        "introduction": "\\label{sec:intro}  The interstellar medium (ISM) in galaxies is excited, dissociated, and ionized by far-ultraviolet (FUV) photons produced by young O and B stars.  Atomic gas in the ISM recombines into molecular form mainly through the catalytic action of dust grain surfaces.  In particular, hydrogen nuclei in the ISM cycle repeatedly from the molecular (\\Htwo) to the atomic (\\HI) phase and back again, at rates depending on the incident FUV flux \\Gzero, the total volume density $n$ of the gas, and the dust-to-gas ratio $\\delta$.  Regions in the ISM where the physics is dominated by FUV photons are called photodissociation regions (PDRs).  The surfaces of giant molecular clouds (GMCs) are important (and ubiquitous) examples of PDRs in galaxies.  The physics of PDRs has been explored in detail over the past $\\sim\\,30$ years by many workers including D.\\ J.\\ Hollenbach, B.\\ T.\\ Draine, A.\\ Dalgarno, J.\\ H.\\ Black, A.\\ G.\\ G.\\ M.\\ Tielens, E.\\ van Dishoeck, J.\\ Le Bourlot, and their students and collaborators.  One major focus of this work has been to explain the $\\approx 1$\\% line-to-continuum ratios of the far-infrared lines of \\CII\\ and \\OI\\ in Galactic sources observed from high-altitude aircraft and balloon platforms.  With the advent of more sensitive space-based observations, the models were extended to include the weak rotational-vibrational spectrum of excited \\Htwo\\ observed on the active PDR surfaces of GMCs that are exposed to relatively intense FUV fluxes from nearby young stars.  An excellent review of the observational and theoretical state of the field is given by \\citet{hol99}.  PDR model computations have become very detailed and comprehensive, raising the possibility that observations of multiple spectral lines including those from trace molecules such as CO can be used to \"invert\" the models in order to determine physical conditions in the ISM.  For example, a set of such models has been computed by \\cite{kau99} for use in the interpretation of far-infrared and submillimeter spectra from Galactic and extragalactic sources.  However, we wish to point out that, although there is qualitative agreement on most of the physics, there are differences in the details of the various model codes currently in use around the world as well as differences in the numerical values of the parameters adopted.  These differences can lead to disagreements in the values derived for physical conditions even for the same observational data:  an example of such differences is illustrated in \\cite{li02}.  Efforts are underway to intercompare models from different groups in a more systematic manner (E.\\ van Dishoeck 2003, private communication).  The results obtained from any one model computation (such as the one we use here) can therefore be revealing in general, but the specific numerical values obtained should be taken with some caution.  The purpose of the present paper is to draw attention to the fact that the \\HI\\ produced from photodissociated \\Htwo\\ can also be a useful diagnostic for determining physical conditions in PDRs.  In particular, the combination of the 21~cm \\HI\\ and 2.6~mm CO(1--0) lines provides a means of independently estimating $n$ and \\Gzero\\ in distant PDRs using ground-based radio astronomy data.  These two radio lines form a diagnostic that is particularly useful in the area of parameter space where \\Gzero\\ $\\lesssim 100$, 10 \\pcmcub $\\lesssim n \\lesssim 10^5$ \\pcmcub; this range of FUV flux and total density is representative of most of the volume of the ISM in galaxies. ",
        "conclusions": "\\subsection{Complementarity}  Figure \\ref{fig:modelboth} clearly shows the complementarity of 21~cm \\HI\\ and 2.6~mm CO(1--0) radio line emission as diagnostics of PDRs.  Over the range of validity of our calculations, the CO(1--0) line brightness $I_{\\rm CO}$ depends mostly on the volume density of the gas for $n \\lesssim 10^3$ \\pcmcub, no matter how intense the FUV flux.\\footnote{The insensitivity of $I_{\\rm CO}$ to the FUV flux comes about because increasing \\Gzero\\ merely leads to more photodissociation of the surface CO, and the CO-emitting layer retreats deeper into the static cloud where the temperatures are about the same as they were for the case of lower \\Gzero.}  On the other hand, for constant $n$ the \\HI\\ brightness increases nearly linearly with increasing FUV fluxes up to \\Gzero\\ $\\approx 10-100$, becoming logarithmic for higher FUV fluxes.  Accordingly, the combination of measurements of $N(\\HI)$ and $I_{\\rm CO}$ from the same PDR can provide unique values for $n$ and \\Gzero\\ over a substantial range in the diagram.  \\subsection{Identifying the Relevant \\HI\\ and CO Emission}  How can extraneous emission from gas that lies far outside the cloud be identified?  The spatially layered structure of PDRs provides a direct clue, as has been convincingly shown e.g.\\ by the observations of the Orion Bar region (see Fig.\\ 2 in \\citet{hol99}) on the $\\sim 1$ pc scale.  A more or less edge-on viewing angle is required for a certain identification, and one ought to remove any extended emission that is associated with other clouds along the line of sight or produced by other excitation mechanisms.  Galaxies that are viewed at low to intermediate inclination angles are favorably oriented for this identification to work, and \\cite{all97} and \\cite{smi00} have shown that the signature morphology can be found at 100 pc scales in the highest resolution \\HI\\ images of nearby galaxies M81 and M101.  As to the CO, it is generally assumed that this arises in PDRs, although the morphological signature is not always clear.  \\subsection{G216~-2.5: An Example}  Figure \\ref{fig:modelboth} provides a method for determining $n$ and \\Gzero\\ for a specific PDR if $N(\\HI)$ and $I_{\\rm CO}$ can be determined observationally.  As an example, we consider the Galactic object G216~-2.5 (also called ``Maddalena's Cloud'') which has been proposed to be a large PDR in the Galaxy by \\citet{wil96}.  \\subsubsection{\\HI\\ Column Density}  Determining the column density of that part of the \\HI\\ associated with G216~-2.5 is made quite difficult by confusion from unassociated \\HI\\ superposed along the line of sight through the Galaxy.  Several estimates can be made from the single-dish observations reported by \\citet{wil96}.  First, and perhaps simplest, we estimate that fraction of the total averaged \\HI\\ profile shown in their Figure 1 that is in the velocity range of the CO(1--0) emission (also shown on the same figure); the \\HI\\ profile shows three overlapping peaks and integrates to $\\approx 48$ K $\\times \\; 58$ \\kmps $\\approx 2800$ K \\kmps = \\pwr{5.1}{21} \\pcmsq.  The averaged $I_{\\rm CO}$ profile corresponds to the central \\HI\\ peak and is $8.5$ \\kmps\\ wide \\citep{mad85}, so the average amount of \\HI\\ column associated with the CO emission is, by this estimate, probably not more than $\\approx (8.5/58) \\times$ \\pwr{5.1}{21} = \\pwr{7.5}{20} \\pcmsq.  In their \\S 3.2.2, \\citet{wil96} carry out a more involved analysis with area integrations of the \\HI\\ in the velocity range 16--38 \\kmps\\ and conclude that the ``excess'' \\HI\\ associated with the PDR is $\\approx$ \\pwr{2}{20} \\pcmsq; the uncertainty in this method is approximately a factor of 2 (R.\\ Maddalena 2002, private communication).  A value in the range \\pwr{1-8}{20} \\pcmsq\\ may therefore be considered typical for this cloud, with some preference for the lower end of that range.  \\subsubsection{CO(1-0) Intensity}  From Figure 1 of \\citet{wil96}, the total $I_{\\rm CO}$ profile over the whole cloud integrates to $\\approx 7.1$ K \\kmps.  The CO(1--0) map for the cloud integrated over the same velocity range of 16--38 \\kmps\\ as the \\HI\\ is shown in Figures 2 and 4 of \\citet{wil96}, based on earlier data of \\citet{mad85}.\\footnote{Note that the $l - b$ maps in \\citet{wil96} are similar in shape but brighter by about a factor of 2 when compared to the earlier map of $I_{\\rm CO}$ in Fig.\\ 2 of \\citet{mad85}.  This is not likely to be an effect of differences in the velocity range over which the data have been integrated, since that range is smaller than the range of 15--40 \\kmps used by Maddalena \\& Thaddeus.  We use the more recent results of Williams \\& Maddalena.} From Figure 2 of \\citet{wil96} the CO(1--0) brightness ranges from 3--7 contours with a few positions reaching 10 contours.  A range of 3--8 contours or 6--16 K \\kmps\\ therefore seems representative of the CO data, again with some preference for the lower end of that range.  \\subsubsection{$n$ and \\Gzero\\ for G216~-2.5}  The range of values for $I_{\\rm CO}$ and $N(\\HI)$ obtained above is plotted as a black box on our Figure \\ref{fig:modelboth}.  In the context of our model, these values describe gas with a density range of \\pwr{0.7-5}{2} \\pcmcub\\ and an incident FUV flux of 0.1--7 of the standard value \\Gzero\\ near the Sun, with likely values being approximately \\pwr{2}{2} \\pcmcub\\ and 0.8, respectively.  Although no independent measurements of the total volume density are available for G216~-2.5, the value we have obtained is in the range generally indicated for GMCs in the Galaxy ($n = n(\\HI) + 2n(\\Htwo) \\approx 2n(\\Htwo) \\sim 2 \\times 50$ \\pcmcub, e.g., \\citet{bli93}).  As to the FUV flux, \\citet{wil96} estimated \\Gzero\\ $\\approx 1$ from what is known about the two nearby young stars identified in their study as likely to be responsible for the photodissociation.  Our model thus agrees well with the observations.  \\subsection{Approximate Detection Limits for \\HI\\ and CO(1--0)} \\label{subsec:DetLims}  The practical detection limits for the \\HI\\ 21 cm line are $\\approx 5$ and $\\approx 100$ K \\kmps\\ for single-dish survey and interferometer array imaging observations, respectively, corresponding to $N(\\HI) \\approx$ \\pwr{1}{19} and \\pwr{2}{20} \\pcmsq.  These are shown as the dashed (single-dish survey) and dotted (interferometer array) lines for $N(\\HI)$ on Figure \\ref{fig:modelboth}.  The corresponding values for the 2.6~mm CO line are $\\approx 1$ K \\kmps\\ (dashed line:  single dish survey) and $\\approx 10$ K \\kmps\\ (dotted line:  interferometer array).  In all cases, this assumes that the gas clouds are resolved; otherwise these limits need to be increased further by the ratio of the beam area to the cloud area.  Locating the detection limits on our Figure \\ref{fig:modelboth}, we can conclude that the range in parameter space over which the combination of \\HI\\ and CO(1--0) can provide useful constraints on $n$ and \\Gzero\\ is given approximately by \\Gzero\\ $\\lesssim 100$, 10 \\pcmcub\\ $\\lesssim n \\lesssim 10^5$ \\pcmcub.  These boundaries are set by the decreasing sensitivity of the \\HI\\ emission at high FUV flux levels ($N(\\HI)$ increases only logarithmically for high \\Gzero\\ at constant $n$) and by observational detection limits ($I_{\\rm CO}$ disappears at low $n$ owing to insufficient excitation of the transition; \\HI\\ disappears at high $n$ as the gas stays mostly molecular).  \\subsubsection{Working with Synthesis Imaging Data}  Many observers who report synthesis imaging (interferometer array) data in the literature use units of Jy beam$^{-1}$ instead of K.  The required conversion factors are (unfortunately) telescope dependent; Appendix \\ref{app:conversions} provides an approximate recipe for these conversions.  \\subsection{Conclusions}  We have shown that a combination of observations in the 21~cm line of \\HI\\ and the 2.6~mm line of CO can provide a useful diagnostic for physical conditions in distant PDRs.  Both of these lines are readily observed in the Galaxy and in nearby galaxies using ground-based single-dish and interferometer-array radio telescopes.  The useful range in parameter space for this combined radio line diagnostic is approximately bounded by \\Gzero\\ $\\lesssim 100$, 10 \\pcmcub $\\lesssim n \\lesssim 10^5$ \\pcmcub, a range that covers a large fraction of the volume of the ISM in galaxies.  The unique layered morphology of the emission in PDRs provides a means of identifying that part of the \\HI\\ and CO emission that is related when that morphology is observable, but confusion along the line of sight makes this separation difficult in the Galaxy.  This method will be especially useful for the interpretation of high-resolution \\HI\\ and CO synthesis imaging data in nearby galaxies where the PDR morphology can be more easily identified.\\footnote{However, in this ``extragalactic'' case, the effects of beam-smearing are severe and must be carefully considered in relating the observational parameters of the models to the measured brightnesses and profile velocity widths.}  Improvements in the model parameters and in the details of the computations are anticipated in the near future that will permit even more precise values of \\Gzero\\ and $n$ to be obtained."
    },
    "0402/hep-th0402075_arXiv.txt": {
        "abstract": "We investigate a string-inspired scenario associated with a rolling massive scalar field on D-branes and  discuss its cosmological implications. In particular, we discuss cosmological evolution of the massive scalar field on the ant-D3 brane of KKLT vacua. Unlike the case of tachyon field, because of the warp factor of the anti-D3 brane, it is possible to obtain the required level of the amplitude of density perturbations. We study the spectra of scalar and tensor perturbations generated during the rolling scalar inflation and show that our scenario satisfies the observational constraint  coming from the Cosmic Microwave Background anisotropies and other observational data. We also implement the negative cosmological constant arising from the stabilization of the modulus fields in the KKLT vacua  and find that this leads to a successful reheating in which the energy density of the scalar field effectively scales as a pressureless dust. The present dark energy can be also explained in our scenario provided that the potential energy of the massive rolling scalar does not exactly cancel with the amplitude of the negative cosmological constant at the potential minimum. ",
        "introduction": "Cosmological inflation has become an integral part of the standard model of the universe \\cite{review}. Apart from being capable of removing the shortcomings of the standard big-bang cosmology, this paradigm has gained a good amount of support from the accumulated observational data. The recent measurement of the Wilkinson Microwave Anisotropy Probe (WMAP) \\cite{WMAP1,WMAP2} in the Cosmic Microwave Background (CMB) made it clear that (i) the current state of the universe is very close to a critical density and that (ii) primordial density perturbations that seeded large-scale structure in the universe are nearly scale-invariant and Gaussian, which are consistent with the inflationary paradigm.  Inflation is often implemented with a single or multiple scalar-field models \\cite{LR}. In most of these models, at least one of the scalar fields undergoes a slow-roll period allowing an accelerated expansion of the universe. It then enters the regime of quasi-periodic oscillations, quickly oscillates and decays into particles leading to reheating. The late time acceleration of universe is supported by observations of high redshift supernovae and indirectly, by observations of the cosmic microwave background and galaxy clustering. The cosmic acceleration can be sourced by an exotic form of matter (dark energy) with a large negative pressure \\cite{phiindustry}. Therefore, the standard model in order to comply with the logical consistency and observation, should be sandwiched between inflation at early epoch and quintessence at late times. It is natural to ask whether one can construct a natural cosmological model using scalar fields to join the two ends without disturbing the thermal history of the universe. Attempts have been made to unify both these concepts using models with a single scalar field \\cite{unifiedmodels}. \\par Inspite of all the attractive features of cosmological inflation, its mechanism of realization still remains ad hoc. As inflation operates around the Planck's scale, the needle of hope points towards the string theory. It is, therefore, not surprising that M/String theory inspired models are under active consideration in cosmology at present. It was recently been suggested that a rolling tachyon condensate, in a class of string theories, might have interesting cosmological consequences. Using the boundary conformal field theory (BCFT) technique, Sen \\cite{s1} has shown that the decay of D-branes produces a pressure-less gas with a finite energy density that resembles a classical dust. He also shown that the same results can be extracted from the tachyon DBI effective action \\cite{as1}. Rolling tachyon matter associated with unstable D-branes has an interesting equation of state $w$ which smoothly interpolates between $-1$ and 0. As the tachyon field rolls down the hill, the universe undergoes an accelerated expansion and at a particular epoch, the scale factor passes through the point of inflection marking the end of inflation. At late times the energy  density of tachyon matter scales as $a^{-3}$, where $a$ is a scale factor. The tachyonic matter was, therefore, thought to  provide an explanation for inflation at the early epochs and could contribute to some new form of cosmological dark matter at late times \\cite{tachyonindustry}. Unfortunately, the effective potentials for rolling tachyon do not contain free parameters that could be tuned to make the roll sufficiently slow to obtain enough inflation and required level of density perturbations \\cite{KL}. The situation could be remedied by invoking the large number of D-branes separated by distance much larger than $l_s$ (string scale). However, the number of branes turns out to be typically of the order of $10^{10}$. This scenario also faces difficulties associated with reheating and the formation of acoustics/kinks \\cite{Frolov}.\\par  In this paper we consider a DBI type effective field theory of rolling massive scalar boson on the D-brane or anti-D brane obtained from string theory. We then consider this effective action for the massive excitation of the anti-D3 brane of the KKLT vacua, and study the cosmological evolution of the scalar rolling from some initial value.  The warp factor, $\\beta$, of the anti-D3 brane provides us an interesting possibility to resolve the problem of the large amplitude of density perturbations. We also take into account the contribution of the negative cosmological constant arising from the stabilization of the modulus fields in the KKLT vacua \\cite{KKLT}. This is important to avoid that the energy density of the rolling massive scalar over dominates the universe after inflation. The present critical density can be explained by considering both the minimum potential of the rolling scalar and the negative cosmological constant. We also evaluate the inflationary observables such as the spectral index of scalar perturbations and the tensor-to-scalar ratio, and examine the validity of this scenario by using a complication of latest observational data. ",
        "conclusions": "In this paper we have presented a scenario based upon a massive scalar field $\\phi$ rolling on the D-brane which was shown in Ref.~\\cite{s2} as a possible solution in string theory. Using a DBI type effective action presented in Ref.~\\cite{mrg1}, we have discussed the cosmological dynamics of the rolling scalar of the anti-D3 brane of KKLT vacua, and have demonstrated that it can lead to inflationary solutions at early epochs. In our scenario it is possible to obtain a sufficient amount of inflation without tuning a fundamental string scale, unlike the case of a rolling tachyon \\cite{KL}. The warp metric in the KKLT vacua provides us a free parameter, the warp factor $\\beta$. The presence of this parameter allows us to obtain the COBE normalized value of density perturbations without changing the brane tension and the masses of massive string states.    We further investigated scalar and tensor perturbations generated during the rolling scalar inflation and showed that our scenario is compatible with recent observational data. It is interesting to note that the steep inflation driven by an exponential potential in the Randall-Sundrum braneworld II scenario is out of the observational contour bounds \\cite{cmbbrane}. This is certainly related to the fact that the dynamics of the Born-Infeld scalar is very different from that of an ordinary scalar field as well as that of a scalar field on the brane.  We have also implemented the contribution of the negative cosmological constant which arises from the stabilization of the modulus fields \\cite{KKLT}. Although this effect is not important during inflation, the dynamics of reheating drastically changes by taking into account the negative cosmological constant that nearly cancels the potential energy of the massive rolling scalar at the potential minimum. One of the problems of the tachyon inflation is that the energy density of the tachyon scales slower than that of the pressure less matter and the radiation \\cite{KL,Frolov}, which means that the tachyon over-dominates the universe after inflation. In our scenario this problem is solved by the negative cosmological constant. As shown in Sec.\\,VII, the average equation of state of the field $\\phi$ approaches that of the pressure-less dust during reheating. We also found that the negative instabity of the field fluctuations in the tachyon case \\cite{Frolov} is not present for the  rolling massive scalar field. This suggests that the reheating proceeds in a similar way to the standard one driven by a massive inflaton field.  The massive rolling scalar can be used to explain the origin of dark energy provided that the potential energy at the minimum ($\\beta^2 T_3$) is very close to the amplitude of the negative cosmological constant. The presence of the negative cosmological constant is crucially important to lead to a successful reheating and to explain the origin of dark energy. Note that in the absence of the $-\\Lambda$ term it is difficult to obtain the amplitude of the critical density unless we choose very small values of $\\beta$.  While we performed a detailed analysis about the spectra of scalar and tensor perturbations generated in the inflationary stage, we did not precisely study the dynamics of reheating including the decay of the field $\\phi$. One interesting aspect is the generation of gauge fields coupled to the rolling massive scalar as was done in Ref.~\\cite{jmc} in the tachyon case. We leave the future work about the precise analysis of the reheating dynamics in our scenario."
    },
    "0402/astro-ph0402442_arXiv.txt": {
        "abstract": "\\baselineskip 11pt The $B$ modes generated by the lensing of CMB polarization are a primary target for the upcoming generation of experiments and can potentially constrain quantities such as the neutrino mass and  dark energy equation of state.   The net sample variance on the small scale $B$ modes out to $l=2000$ exceeds Gaussian expectations by a factor of 10 reflecting the variance of the larger scale lenses that generate them.   It manifests itself as highly correlated band powers with correlation coefficients approaching 70\\% for wide bands of $\\Delta l/l \\sim 0.25$. It will double the total variance for experiments that achieve a sensitivity of approximately 4 $\\mu $K-arcmin and a beam of several arcminutes or better This non-Gaussianity must be taken into account in the analysis of experiments that go beyond first detection. ",
        "introduction": "As a step on the road toward the ultimate goal of detecting primordial gravitational waves, upcoming cosmic microwave background (CMB) polarization experiments will target the distortion to the acoustic polarization induced by gravitational lensing. As with the polarization induced by gravitational waves, the gravitationally lensed polarization contains a component with handedness, the so-called  $B$ mode component \\cite{ZalSel98}. Unlike gravitational waves, gravitational lensing provides a guaranteed signal.  In the standard cosmological model, the predicted amplitude of the $B$ modes can only vary at the tens of percent level within current constraints \\cite{Speetal03}. Moreover these fine variations provide an opportunity to measure the dark side of the universe, namely the dark energy and neutrino dependent growth of structure, as well as another handle on the reionization optical depth \\cite{Hu01c,KapKnoSon03}.  Although both the intrinsic distribution and the density perturbations that lens the CMB are expected to be Gaussian, the lensed distribution is non-Gaussian at second order in the perturbations.    The non-Gaussianity is therefore relatively small in the temperature distribution \\cite{Hu01}.  However because the $B$ modes are generated by the lensing itself, its non-Gaussianity is a first order effect but fortunately one that is precisely calculable. Gravitational lensing therefore also provides a unique testing ground for experimentally extracting a non-Gaussian signal in the presence of foregrounds and systematic errors. Ultimately, the non-Gaussianity of the lensed polarization also provides the key to mapping the dark matter \\cite{HuOka01,HirSel03} and hence the separation of the lensing and gravitational wave $B$ mode components \\cite{KnoSon02,KesCooKam02}.   For the upcoming generation of experiments, the non-Gaussianity will provide an important source of uncertainty for power spectrum measurements.  Fisher information studies have shown that the information contained on cosmological parameters in the $B$ mode power spectrum under the Gaussian approximation unphysically exceeds that contained in the two underlying Gaussian fields \\cite{Hu01c}.   In this paper, we study the origin and quantify the impact of the non-Gaussian sample variance  on $B$ mode power spectra measurements.   The basic reason for the large sample variance is that the fluctuations that lens the CMB are mainly on degree scales.  All of the arcminute scale $B$ modes fluctuate jointly with the lens and so precision measurements will require many degree scale patches not simply many arcminute scale patches.   We begin in \\S \\ref{sec:bmodes} by briefly reviewing the generation of $B$ modes through gravitational lensing.  We calculate the non-Gaussian sample covariance in \\S \\ref{sec:covariance} and explore its impact on measurements in \\S \\ref{sec:impact}. We conclude in \\S \\ref{sec:discussion}. For illustrative purposes we employ throughout a fiducial cosmology that is consistent with WMAP determinations: an initial scale invariant spectrum of curvature fluctuations with amplitude $\\delta_\\zeta =5.07\\times 10^{-5}$  ($\\sigma_8=0.91$, $\\tau=0.17$), a baryon density $\\Omega_bh^2 =0.024$ and a matter density $\\Omega_m h^2 =0.14$ in a flat $\\Omega_\\Lambda=0.73$ cosmology. ",
        "conclusions": "\\label{sec:discussion}  The non-Gaussianity of the $B$ modes in the lensed CMB polarization substantially degrades the amount of information contained in the $B$ mode power spectrum. It both increases the variance of band powers and makes them strongly covary across a wide range in $l$ surrounding the peak power.   Ultimately it will increase the variance of the amplitude of the power spectrum by an order of magnitude.  As experiments move from the upper limit and first detection stage to using the $B$ mode power spectrum to constrain the properties of dark components such as the neutrinos and dark energy, this non-Gaussianity will have to be included in the analysis.  By quantifying the sample covariance, we have provided the analytic and numerical tools that will be the basis for such an analysis.  The advantage of the Monte Carlo approach is that it can be straightforwardly applied to any estimator of $B$ power.  In principle, one can include the sample covariance in an effective $\\chi^2$ as is done for power spectrum errors of the temperature field (e.g. \\cite{Veretal03}).  However since the computation of the covariance is much more costly than the computation of the $B$ power spectrum, minimization in a large-dimensional cosmological parameter space is impractical even with Monte Carlo Markov Chain techniques. Since most of the parameters affecting the high redshift universe will be fixed prior to these measurements from the $T$ and $E$ mode spectra, as a first order correction one can follow the Fisher matrix approach and calculate the effect in a fiducial model.  More specifically, given the correlation matrix ($R_{ij}$) and the relative variance degradation ($D_{i}$) in a fiducial model, one can scale the covariance matrix to the model $B$ mode spectrum ($C_l^{BB}$) as calculated from Boltzmann codes.   This approach would capture the main scaling of the covariance through the amplitude of the lensing power spectrum. Implementing such a pipeline though is beyond the scope of this paper.  As experiments move from parameter constraints based on power spectra to mapping the lensing potential \\cite{HuOka01,HirSel03}, the non-Gaussianity of the polarization becomes the signal and not the noise. The extent to which this ultimate goal will be achievable instrumentally and in the presence of foregrounds  \\cite{Bowetal03} awaits the results of the upcoming generation of experiments.  \\smallskip \\noindent{\\it Acknowledgments:}  We thank T. Okamoto and B. Winstein for useful discussions. KMS and WH were supported by NASA NAG5-10840, the DOE and the Packard Foundation. MK was supported by the NSF and NASA NAG-11098. It was carried out  in part at  the CfCP under NSF PHY-0114422."
    },
    "0402/astro-ph0402168_arXiv.txt": {
        "abstract": "We present an overview of our ongoing systematic search for wide (sub)stellar companions around the stars known to host rad-vel planets. By using a relatively large field of view and going very deep, our survey can find all directly detectable stellar and massive brown dwarf companions (m$>$40\\,$M_{Jup}$) within a 1000\\,AU orbit. ",
        "introduction": "Circumstellar disks are discovered with sizes up to 1000\\,AU and also binary stars with the same comparable separations are known. Because the formation of stars and brown dwarfs seems to follow a similar scheme (fragmentation of large gas clouds) substellar objects may indeed reside in that distance around stars hosting rad-vel planets. Adaptive optics search programs to find very close companions around those stars already exist, but they leave out an interesting regime of objects, namely the wide companions because of a too small field of view (e.g. Patience et al. 2002).  As of October 2003, more than one hundred extra-solar planets were discovered. Many of those have extremely close orbits which could be explained by a migration process in the early history of the system. During this migration angular momentum is transferred from the inner part of the accretion disk to its outer border. A wide companion can cut off the disk and be a sink for the lost angular momentum. Furthermore theories predict that wide companions can induce rapid instability in disks which otherwise would be stable, hence they could have a strong influence on the planet formation and on the longtime evolution of planetary orbits.  Actually, some extrasolar planets were found to reside in binary stellar systems. Those few cases are intriguing, and might exhibit some statistically different features than the planets around single stars (Zucker \\& Mazeh 2002). This could be a first hint about an interaction between the (sub)stellar wide companions and the extrasolar planets. Nevertheless, the whole sample of extrasolar planetary systems has not been surveyed completely for wide companions with sensitive IR cameras that are able to find faint low-mass companions. For this reason we have started in 2001 a systematic deep imaging of all the stars known to harbor planets, in order to look for faint companions in wide orbits. The companionship of those faint objects can be established only by follow up observations which will detect common proper motion with the nearby star that host the planet.  To find all the companion-candidates around the planet hosting stars we secure deep IR images, obtained with the IR cameras SOFI at the 3.58\\,m NTT and UFTI at the 3.8\\,m UKIRT, with detection limit of H\\,$\\sim$\\,19.5\\,mag. To detect common proper motion we obtain two images about one year apart. We also make use of the 2 micron all sky survey (2MASS) images, which were taken several years before our exposures. However, the limit of 2MASS is H\\,$\\sim$\\,15\\,mag (Cutri et al. 2003), and the proper motion of fainter objects needs to be measured only by our two images. ",
        "conclusions": ""
    },
    "0402/astro-ph0402397_arXiv.txt": {
        "abstract": "{ We present the first results of a 2-year high-resolution spectroscopy campaign of 59 candidate $\\gamma$\\,Doradus stars which were mainly discovered from the HIPPARCOS astrometric mission. More than 60\\,\\% of the stars present line profile variations which can be interpreted as due to pulsation related to $\\gamma$\\,Doradus stars. For all stars we also derived the projected rotation velocity (up to more than 200\\,\\kms). The amplitude ratios $2K/\\Delta m$ for the main HIPPARCOS frequency are in the range 35 - 96\\,\\kms\\,mag$^{-1}$. Less than 50\\,\\% of the candidates are possible members of binary systems, with 20 stars being confirmed $\\gamma$\\,Doradus. At least 6 stars present composite spectra, and in all but one case (for which only one spectrum could be obtained), the narrow component shows line profile variations, pointing towards an uncomfortable situation if this narrow component originates from a shell surrounding the star. This paper is the first of a series concerning mode identification using both photometric and spectroscopic methods for the confirmed $\\gamma$\\,Doradus stars of the present sample. ",
        "introduction": "In the coming decade, thanks to dedicated satellites (COROT, EDDINGTON), the detailed knowledge of the internal structure of stars should be achieved through the technique of asteroseismology. The goal of this, relatively new, research domain is to derive the internal processes in stars with an unprecedented precision through a detailed study of their oscillations.  This paper deals with a class of non-radial pulsators along the main sequence, namely the $\\gamma$\\,Doradus stars (see e.g. Kaye et al. (\\cite{kh99a}) for the main observational characteristics of this class of variables). These stars are multiperiodic high-order gravity-mode oscillators with spectral types around F0. The origin of the mode destabilization is not clearly known yet, and driving mechanisms have been proposed by Guzik et al. (\\cite{gk00}), Wu (\\cite{w02}) and L\\\"offler (\\cite{l02}).  Much effort is currently made to find new members of this group, to constrain their pulsation characteristics and their position in the HR Diagram, especially the $\\gamma$\\,Doradus star's red border in relation with the solar-like star's blue border. Indeed, they show quite a large variety in their observational behaviour, and the number of confirmed members is still low. This observational campaign should contribute to the necessary comparison between the observational HR diagram and the theoretical one recently defined by Warner et al. (\\cite{wk03}). Because of their relatively low amplitude (few tens of mmag in photometry, of the order of 1\\,\\kms\\ in radial velocity), and due to the long time scales of the variation (between 0.3 and 3\\,d), the detection of such variables is still difficult. Up to now the best tool has been the HIPPARCOS satellite. The HIPPARCOS sampling does not suffer from the aliasing problems of a single Earth site which is of particular annoyance for $\\gamma$\\,Dor studies. However, there are two major drawbacks: the precision of photometric individual measurements degrades quite rapidly for fainter stars and the non-continuous sampling makes the detection/interpretation of multiperiodic phenomena difficult.  Several studies selected $\\gamma$\\,Dor candidates from HIPPARCOS: Eyer (\\cite{e98}) proposed a list of such candidates extracted from the periodic variable stars in the HIPPARCOS variability annex which have well defined absolute magnitude and colour.  Aerts, Eyer and Kestens (1998) used stars from the same catalogue which have furthermore Geneva photometry. It permitted to use a multivariate discriminant analysis which proved to be very efficient for detecting new slowly pulsating B stars (Waelkens et al.\\ \\cite{wa98}), which are also main-sequence gravity-mode oscillators. Handler (\\cite{h99}) broadened the search for $\\gamma$\\,Dor stars to the unsolved variable stars of the HIPPARCOS variability annex and relaxed selection criteria, focusing more on the nature of the power spectra. These studies proposed about 60 bona fide and prime candidates stars. One star in our sample, \\object{HD\\,173977}, which was in Handler's list (\\cite{h99}), has been discarded since it is now classified as a $\\delta$\\,Scuti variable (Chapellier et al. \\cite{cm03}).  However, the spectroscopic studies of most of the candidates having well-known photometric properties are much less detailed. In 2001, we undertook a spectroscopic campaign whose objective was twofold: \\begin{itemize} \\item to derive basic spectroscopic parameters (rotation velocity, line profile variations, duplicity, etc.) for better identification of the pulsation modes. This is particularly important for stars with similar rotation and pulsation frequencies (Dintrans \\& Rieutord \\cite{dr00}). \\item to prepare the COROT and EDDINGTON space missions by including at least one $\\gamma$\\,Doradus star in the core program (Mathias et al. \\cite{mc03a}, \\cite{mc03b}). \\end{itemize}  This paper presents the first descriptive part of the campaign concerning 59 $\\gamma$\\,Doradus candidates. Only the very homogeneous OHP spectroscopy is discussed. The observations are described in Sect.\\,2. In Sect.\\,3, results are given for individual stars, depending on the detection of line profile variations, in Sect.\\,3. Sect.\\,4 and 5 present the pulsation and stellar environments for some candidates. Concluding remarks are given in Sect.\\,6. ",
        "conclusions": "We have presented spectroscopic observations of 59 candidate $\\gamma$\\,Doradus stars detected mainly from the HIPPARCOS space mission. The main goal was to confirm these stars as real members of the group through the presence of line profile variations typical of $g$-mode pulsations. The $\\gamma$\\,Doradus stars that are confirmed by the present work, in addition to the ``bona fide'' candidates given in Table\\,\\ref{table1} are \\object{HD\\,48271}, \\object{HD\\,70645}, \\object{HD\\,80731}, \\object{HD\\,100215}, \\object{HD\\,113867}, \\object{HD\\,175337}, \\object{HD\\,195068}. We were unable to detect LPV in less than 40\\,\\% of the candidates, but most stars being (spectroscopically) faint, the signal to noise ratios were not always sufficient to detect very weak variations. Moreover, for most stars we have a very limited number of spectra, so LPV cannot be ruled out for these candidates.  In only a very few cases were we able to impose the main HIPPARCOS frequency on the radial velocity curves deduced from the LPV. The deduced $2K$ amplitudes are generally low (between 0.6 and 4.2\\,\\kms), pointing towards a mean amplitude ratio of about 60\\,\\kms\\,mag$^{-1}$. The pulsation behaviour for the most interesting stars (observations are on-going) will be described in subsequent papers.  Fekel et al. (\\cite{fw03}) suggest a percentage of $\\gamma$\\,Doradus members of multiple systems as high as 74\\,\\%. Our larger sample, containing however a larger proportion of stars which are not confirmed $\\gamma$\\,Doradus stars, shows that this percentage seems to be smaller, i.e. 50\\,\\%. This value is still larger than the one measured for such stars (30\\,\\%) in a previous radial velocity study (Nordstr\\\"om et al. \\cite{ns97}). Similar to that occuring in a number of $\\delta$\\,Scuti-type pulsators (Lampens \\& Boffin \\cite{lb00}), we also found several $\\gamma$\\,Doradus variables in binary systems with eccentric orbits.  Our sample contains 6 stars that show composite spectra. This behaviour can be due either to binarity or to the presence of a shell surrounding the star. Our data easily show that the narrow component presents LPV in 5 out of the 6 candidates. If a shell is really present, one has to find the mechanism that induces LPV in this shell. A first step would be to detect the period, if existing, of the variations of this narrow component."
    },
    "0402/astro-ph0402674_arXiv.txt": {
        "abstract": "We interpret the rapid correlated  UV/optical/ X-ray variability of XTE~J1118+480 as a signature of the coupling between the X-ray corona and a jet emitting synchrotron radiation in the optical band. We propose a scenario in which the jet and the X-ray corona are fed by the same energy reservoir where large amounts of accretion power are stored before being channelled  into either the jet or the high energy radiation. This time dependent model reproduces the main features of the rapid multi-wavelength variability of XTE~J1118+480. Assuming that the energy is stored in the form of magnetic field, we find that the required values of the  model parameters are compatible with both a patchy corona atop a cold accretion disc and a hot thick inner disc geometry. The range of variability timescales for the X-ray emitting plasma are consistent with the dynamical times of an accretion flow between 10 and 100 Schwarzschild radii. On the other hand, the derived range of timescales associated with the dissipation in the jet extends to timescales more than 10 times larger, confirming the suggestion that the generation of a powerful outflow requires large scale coherent poloidal field structures. A strong requirement of the model is that the total jet power should be at least a few times larger than the observed X-ray luminosity, implying a radiative efficiency for the jet $\\epsilon_{\\rm j} \\la 3 \\times 10^{-3}$. This would be consistent with the overall low radiative efficiency of the source. We present independent arguments showing that the jet probably dominates the energetic output of all accreting black holes in the low-hard state. ",
        "introduction": "The high energy spectrum ($>$ 1 keV) of accreting stellar mass black holes in the low/hard state can be roughly described by a power-law with photon index $\\Gamma \\sim 1.4-2$, and a nearly exponential cut-off at a characteristic energy $E_{\\rm c}$  of a few hundred keV (see e.g. Tanaka \\& Lewin 1995; Gierlinski et al. 1997; McClintock \\& Remillard 2004). Such a spectrum is generally interpreted as due to thermal Comptonisation in a plasma with electron temperature $kT_{\\rm e}\\sim$ 100 keV and Thomson optical depth $\\tau\\sim$ 1 (see e.g. Poutanen 1998). There are two possible explanations for the presence of this very hot plasma (often called {\\it corona}). It could be either a geometrically thick, optically thin innermost part of the accretion flow (Shapiro, Lightman \\& Eardley 1976; Narayan \\& Yi 1994) or a collection of small scale active regions located atop a cold, geometrically thin and optically thick accretion disc (Haardt, Maraschi \\& Ghisellini 1994), possibly powered by magnetic reconnection.  When the X-ray luminosity increases above a few percent of the Eddington luminosity ($L_{\\rm Edd}$), accreting black holes are observed to switch from the low/hard state to the so-called high/soft-state (see e.g. McClintock \\& Remillard 2004). Then the X-ray luminosity is dominated by a strong thermal component originating from a geometrically thin, optically thick disc. Current observations show that the power-law component is steeper ($\\Gamma \\sim 2.5-3$) and much less luminous than in the hard state, suggesting the disappearance of the Comptonising plasma in this state.  Recent multi-wavelength observations of accreting  black holes in the hard state have shown the presence of an ubiquitous flat-spectrum radio emission (see e.g Fender 2004), that may extend up to infrared and optical wavelengths. The properties the radio emission indicate it is likely  produced by synchrotron emission from relativistic electrons in compact, self-absorbed jets (\\citeauthor{bk79}, 1979; \\citeauthor{hj88}, 1988). This idea was confirmed by the discovery of a continuous and steady milliarcsecond compact jet around Cygnus X-1  (\\citeauthor{sti01}  2001). Moreover, in hard state sources a tight correlation has been found between the hard X-ray and radio luminosities, holding over more than three decades in luminosity (Corbel et al. 2003; Gallo, Fender \\& Pooley 2003). In contrast, during high-soft state episodes the sources appear to be radio weak (Tananbaum et al. 1972; Fender et al. 1999; \\citeauthor{cor00} 2000), suggesting that the Comptonising medium of the low/hard state is closely linked to the continuous ejection of matter in the form of a small scale jet.  Merloni \\& Fabian (2001) have pointed out that the energy content of the electrons in the Comptonising medium is too low  to account for the observed X-ray luminosities. The Comptonising electrons have to be tightly connected to an energy reservoir where large amount of accretion power is stored before being transfered to them, and ultimately radiated. Independently,  the rapid X-ray variability in Cyg X-1 also suggests the presence of of energy reservoirs (Negoro et al. 1995; Maccarone \\& Coppi 2002). Because angular momentum transport in accretion flow is most likely due to magneto-rotational instability (MRI, see Balbus \\& Hawley, 1998), a natural candidate for the energy repository is the magnetic field amplified in the disc by the MRI-turbulent flow. However, if the dissipation of the tangled field (via magnetic reconnection) and non-adiabatic turbulent heating preferentially energize the protons, rather than the electrons, the hot Comptonising plasma could become two-temperature, and the protons themselves act as the main energy reservoir (\\citeauthor{dbf97}, 1997).  Models and simulations of jet production (Blandford \\& Znajek 1977; Blandford \\& Payne 1982; Meier 2001)  indicate that jets are driven by the poloidal component of the magnetic field. Therefore  storage of energy into  magnetic structures driving the jet and powering the Comptonising electrons is the most straightforward explanation for the observed corona-jet association \\footnote{See however Markoff, Falcke \\& Fender (2001) and Georganopoulos, Aharonian \\& Kirk (2002) for alternative interpretations involving dominant X-ray emission from the jet.}. In this context, the corona would constitute the location where the jet is launched. These idea were developed by several authors in the context of geometrically thin and/or thick accretion flows (Meier 2001; Livio, Pringle \\& King 2003) as well as accretion disc coronae (Merloni \\& Fabian 2002).  Besides the radio/X-ray correlation  observed on long ($>$ 1 day) timescales, there are  indications that the corona-jet coupling operates on timescales as short as a few seconds or less. The best example is provided by the X-ray nova XTE J1118+480 (Remillard et al. 2000; McClintock et al. 2001a). During its outburst in 2000, this black hole showed all the X-ray properties of hard state sources. Fast optical and UV photometry  has shown rapid optical/UV flickering presenting complex correlations with the X-ray variability (Kanbach et al. 2001; Hynes et al. 2003, hereafter K01 and H03 respectively). This correlated variability cannot be caused by reprocessing of the X-rays in the external parts of the disc. Indeed, the optical flickering occurs on average on shorter time-scales than the X-ray one (K01), and reprocessing models fail to fit  the complicated shape of the X-ray/optical cross correlation function (H03). Spectrally, the jet emission  seems to extend at least up to the optical band (McClintock et al. 2001b; Chaty et al. 2003, hereafter C03), although the external parts of the disc may provide an important  contribution to the observed flux at such wavelengths. The jet  activity is thus the most likely explanation for the rapid observed optical flickering. For this reason, the properties of the optical/X-ray correlation in XTE J1118+480 might be of primary importance for the understanding of the jet-corona coupling and the ejection process.  The simultaneous optical/X-ray observations are described at length in a number of papers (K01; Spruit \\& Kanbach 2001; H03; Malzac et al. 2003, hereafter M03). As discussed in these works, the observations are very challenging for any accretion model. The most puzzling pieces of evidence are the following: (a) The optical/X-ray Cross-Correlation Function (CCF) shows the optical band lagging the X-ray by $~$0.5 s, but with a dip 2-5 seconds in advance of the X-rays (K01); (b) The correlation between X-ray and optical light curves appears to have timescale-invariant properties: the X-ray/optical CCF maintains a similar, but rescaled, shape on timescales ranging at least from 0.1 s to few 10 s (M03); (c) The correlation does not appear to be  triggered by a single type of event (dip or flare) in the light curves; instead, as was shown by M03, optical and X-ray  fluctuations of very different shapes, amplitudes and timescales are correlated in a similar way, such that the optical light curve is related to the time derivative of the X-ray one. Indeed, in the range of timescales where the coherence is maximum, the optical/X-ray phase lag are close to $\\pi/2$, indicating that the two lightcurves are related trough a differential relation. Namely, if the optical variability is representative of fluctuations in the jet power output  $P_{\\rm j}$, the data suggest that the jet power scales roughly like $P_{\\rm j} \\propto -\\frac{dP_{\\rm x}}{dt}$, where $P_{\\rm x}$ is the X-ray power.  Here we will show that, if indeed there is a common energy reservoir feeding both the jet and the corona, this differential relation is naturally satisfied, provided that the jet power dominates over the X-ray luminosity.  We will first present the energy reservoir model and suggest a simple physical scenario for the energy reservoir and jet disc/coupling (Section \\ref{sec:reserv}). We  will then present a time dependent model that captures the main features of the multi-wavelength variability observed in XTE J1118+480, and discuss our main results obtained by comparing the properties of the simulated lightcurves with the observations (Section \\ref{sec:results}). Section \\ref{sec:flow} will be devoted to a discussion of the constraints on the nature of the accretion flow in XTE J1118+480 derived from both spectral and temporal analysis, while in Section \\ref{sec:jet}, we make an attempt to generalize these results to other accreting black hole sources. Finally, we summarize our conclusions in section \\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  We have shown that the puzzling optical/X-ray correlations of XTE~J1118+480, can be understood in terms of a common energy reservoir for both the jet and the Comptonizing electrons. For illustration purpose, we have presented a specific shot noise variability model that reproduces {\\it all} the main observed features of the multi-wavelength correlated variability.  Our time dependent model is fairly general, and our main conclusions hold regardless of the specific geometrical and dynamical properties of the system. The main results can be summarized as follows: \\begin{itemize} \\item{Any energy reservoir model for XTE J1118+480 requires that the total jet power dominates over the X-ray luminosity. In particular, assuming that the compact jet synchrotron emission extends up to the optical band, this implies a radiative efficiency for the jet $\\epsilon_{\\rm j}\\la 3\\times 10^{-3}$. Following the same line of arguments as FGJ03, we showed that this situation is likely and probably represents a common feature of all black holes in the low-hard state} \\item{The range of typical variability timescales of dissipation rate of the X-ray  emitting plasma is consistent with the dynamical times of an accretion disc between $\\sim$10 and 100 Schwarzschild radii (for a 10 $M_{\\odot}$ black hole). As expected, the X-ray variability can be associated with either a hot, thick inner accretion flow or with a patchy corona.} \\item{As jet launching  requires large-scale coherent magnetic structures, the energy dissipation in the jet should vary on longer time scales. Indeed, the derived range of timescales associated with the dissipation in the jet extends to timescales more than 10 times larger than the X-ray emitting one.} \\item{By combining the information obtained from our time variability model with (time averaged) spectral analysis we conclude that, whatever the accretion geometry, the whole disc-corona-jet system must be radiatively inefficient. It is therefore possible, in principle, that the system is energetically dominated by the jet, as suggested by FGJ03, although whether the bulk of the accretion power is advected into the black hole or into the jet remains at present an open question. } \\item{In terms of specific dynamical models, we conclude that the observed properties of XTE J1118+480 during its low/hard state outburst are consistent with either an inner hot, quasi spherical, radiatively inefficient flow, from which the jet originates (Meier 2001), surrounded by a geometrically-thin, optically-thick cold disc, or with a powerful patchy, outflowing corona on top of an extremely cold standard thin disc. In the first case, multicolour disc fits to the UV/EUV spectrum indicate a very large inner disc radius, implying a large total accretion rate ($\\dot m\\simeq 0.2$), which might be in conflict with the hypothesis of standard advection dominated flow theory. In the second case, in order to reproduce the very low inner disc temperature, an (uncomfortably) extreme coronal power is needed, together with substantial relativistic bulk motion of the coronal plasma, both possibly associated with very high magnetic viscosity.} \\end{itemize}"
    },
    "0402/astro-ph0402504_arXiv.txt": {
        "abstract": "The specific angular momentum of Cold Dark Matter (CDM) halos in a $\\Lambda$CDM universe is investigated. Their dimensionless specific angular momentum $\\lambda'=\\frac{j}{\\sqrt{2}V_{vir} R{vir}}$ with $V_{vir}$ and $R_{vir}$ the virial velocity and virial radius, respectively depends strongly on their merging histories. We investigate a set of $\\Lambda$CDM simulations and explore the specific angular momentum content of halos formed through various merging histories. Halos with a quiet merging history, dominated by minor mergers and accretion until the present epoch, acquire by tidal torques on average only 2\\% to 3\\% of the angular momentum required for their rotational support ($\\lambda'=0.02$). This is in conflict with observational data for a sample of late-type bulgeless galaxies which indicates that those galaxies reside in dark halos with exceptionally high values of $\\lambda' \\approx 0.06-0.07$. Minor mergers and accretion preserve or slowly increase the specific angular momentum of dark halos with time. This mechanism is however not efficient enough in order to explain the observed spin values for late-type dwarf galaxies. Energetic feedback processes have been invoked to solve the problem that gas loses a large fraction of its specific angular momentum during infall. Under the  assumption that  dark halos hosting bulgeless galaxies acquire their  mass  via  quiescent accretion, our results indicate yet another serious problem: the specific angular momentum gained during the formation of these objects is not large enough to explain their observed rotational properties, even if no angular momentum would be lost during gas infall. ",
        "introduction": "The dynamical structure of disc galaxies is dominated by angular momentum.  Therefore, understanding the origin of angular momentum in these systems is crucial in any theory of galaxy formation. In the current paradigm for structure formation, dark matter is assumed to be cold and collisionless and luminous galaxies form by gas infall into  dark matter halos,  which grow by gravitational accretion and merging in a hierarchical fashion (White $\\&$ Reese 1978). Fall $\\&$ Efstathiou (1980, hereafter FE80) proposed that the sizes of galactic disks are linked to the angular momentum of their parent dark matter halos. This theory is able to produce disks with sizes that are in agreement with observations, if the gas initially had the same specific angular momentum as dark matter halos show today and if the gas preserved its specific angular momentum during the protogalactic collapse phase. The angular momentum of galaxies results from torques due to tidal interactions with neighbouring structures, acquired early, before the halo decoupled from the Hubble expansion.  Many models for the formation of galactic disks have been proposed, based on the picture of FE80. Most of them incorporate the mass accretion history of halos and are able to reproduce many properties of observed disk galaxies (e.g. Mo, Mao $\\&$ White 1998; Firmani $\\&$ Avila-Reese 2000; van den Bosch 2000). However, in  these models, the angular momentum of the dark matter halos is assigned without accounting for their merging history. Recent results from numerical N-body simulations have pointed out that the effect of major mergers is to increase the mean angular momentum content of the halos (Gardner 2001, G01 hereafter; Vitvitska et al. 2001). This is explained by the orbital angular momentum of the merging halos which dominates the final net angular momentum of the remnant (G01). However, numerical simulations that incorporate gas dynamics have difficulties to make realistic disk galaxies in the current cosmological paradigm. In most simulations, the disks are smaller, denser and have much lower angular momenta than observed disk systems.  Simulations show that galaxies are built up by merging of baryonic subclumps, rather than smooth accretion of gas. Most of the gas cools at the centre of subhalos and spirals toward the center of the parent halo, transferring orbital angular momentum to the surrounding dark matter (e.g.  Navarro $\\&$ Steinmetz 2000). More acceptable fits to real disk systems can be found if heuristic prescriptions of stellar feedback are included in the simulations (Sommer-Larsen et al. 2003; Abadi et al. 2003, Robertson et al. 2004). However, even in these simulations the disk systems typically contain denser and more massive bulges than observed in late-type galaxies.  Most of the previous work has focused on angular momentum properties of halos that had at least one dominant major merger during their evolution. In this Letter we explore the angular momentum properties of halos that did not experience any major merger since redshift 3 and that are in principle good candidates to host bulgeless galaxies. We demonstrate that even these galaxies have an angular momentum problem that is directly related with the spin of their dark halos and cannot be solved easily by energetic feedback processes. ",
        "conclusions": "If bulgeless galaxies have not experienced any major mergers during their evolution our analysis shows that their dark halos  are characterized by systematically lower angular momentum than observed. Halos without major mergers acquire their specific angular momentum through tidal torques in the early epochs of evolution (Barnes $\\&$ Efsthatiou 1987), when the density contrasts were small, in accordance with the prediction of the linear theory. In Fig.3, top panel, the spin parameter evolution of the major progenitors of two representative halos is shown. Due to  the lack of any major merger event, there is no sharp increase in $\\lambda'$ which is characteristic for major mergers. Instead, $\\lambda'$ is almost constant and slightly increasing with time, proving that minor mergers and accretion does not affect or significantly  decrease $\\lambda'$, in contradiction with claims of Vitvitska et al. 2002. The bottom panel shows that the halos aggregate mass gradually by  accretion of small subhalos. Robertson et al. (2004) have recently    presented  the  results  of  a   cosmological hydrodynamic simulation where they form an extended, bulge-less disk galaxy. In our sample of \"quiescent\" halos, there is one case with high spin value ($\\lambda'=0.05$ at z=0), increasing more than a factor of 2 from z=0 and z=3 and no major merger during that time. In its accretion history, mergers with mass ratio 1:5, 1:6 seem to be enough to spin up the halo, although they are unlikely in CDM scenario. Thus, of course this object could be a good candidate to be a bulgeless galaxies and we are resimulating it with higher resolution (D'Onghia et al. in preparation), but cannot be considered as a representative ``quiescient'' halo. Robertson et al. 2004 could have simulated an object with similar properties, although the authors don't report the spin parameter value or the accretion history  of their object.  The net result seems to be that  tidal torques generate halos with typical spin parameters of $\\lambda' = 0.02$ whereas data for bulgeless galaxies indicate halos with values of $\\lambda' = 0.06 - 0.07$. One might worry that our result is affected by the adopted small cosmological volume of 15 $h^{-1}$ Mpc box size, that might suppress torques by by larger-scale structures. However tidal forces scale as $F \\propto 1/r^3$. Compared to the forces that are generated by structures at a Mpc distance, structures located at distances of 10 Mpc need to have the mass $10^3$ larger than at 1 Mpc to have the same effect on protogalaxies. At 100 Mpc, the mass of the structure would have to be even$10^6$ times larger to torque protogalaxies efficiently, which is unlikely. In addition, our results are in agreement with the models of G01 for a larger volume of 100 Mpc box size.  At the moment it is not at all clear how one can reconcile the observations and theory. It is known that the spin parameter distribution for the collapsed objetcs is insensitive to the shape of the initial power spectrum of density fluctuations, to the environment and the adopted cosmological model (Lemson $\\&$ Kauffmann 1999). A modification of the nature of dark matter does not seem to solve the problem either. Recent works show that warm dark matter  halos have systematically smaller spins than their counterparts in $\\Lambda$CDM (Knebe et al. 2002; Bullock, Kravtsov $\\&$ Colin 2002), although in this scenario the presence of pancakes could provide more efficient torques on the protohalos. Feedback was invoked as a mechanism to prevent the process of drastic angular momentum loss of infalling gas. (BBS; Maller, Dekel $\\&$ Somerville 2002). However, we have shown here that the dark halos that experienced no major mergers have already too low an angular momentum to produce the observed disks and no feedback process is know that would increase the the specific angular momentum of the gas. The origin of extended bulgeless disk galaxies remains a puzzle."
    },
    "0402/astro-ph0402218_arXiv.txt": {
        "abstract": "We investigate the process of galaxy formation as can be observed in the only currently forming galaxies -- the so-called Tidal Dwarf Galaxies, hereafter TDGs -- through observations of the molecular gas detected via its CO (Carbon Monoxide) emission. Molecular gas is a key element in the galaxy formation process, providing the link between a cloud of gas and a {\\it bona fide} galaxy. We have now detected CO in 9 TDGs with an overall detection rate of 80\\%, showing that molecular gas is abundant in TDGs, up to a few $10^8 M_\\odot$.  The CO emission coincides both spatially and kinematically with the HI emission, indicating that the molecular gas forms from the atomic hydrogen where the HI column density is high. A possible trend of more evolved TDGs having greater molecular gas masses is observed, in accord with the transformation of HI into H$_2$. Although uncertainties are still large for individual objects as the geometry is unknown, we find that the ``dynamical\" masses of TDGs, estimated from the CO line widths, do not seem to be greater than the ``visible\" masses (HI + H$_2$ + a stellar component), i.e., TDGs require no dark matter. We provide  evidence that TDGs are self-gravitating entities, implying that we are witnessing the ensemble of processes in galaxy formation: concentration of large amounts of gas in a bound object, condensation of the gas, which is atomic at this point, to form molecular gas and the subsequent star formation from the dense molecular component. ",
        "introduction": "Tidal Dwarf Galaxies (TDGs) are small galaxies which are currently forming from material ejected from the disks of spiral galaxies through collisions. They allow us to observe processes -- galaxy formation and evolution -- similar to what occurred in the very early universe but in very local objects.  As a consequence, they can be studied with a sensitivity and a resolution inconceivable for high-redshift sources.  Because galactic collisions can be well reproduced through numerical simulations, it is possible to obtain good age estimates for the individual systems (e.g. Duc et al. 2000). The formation of TDGs is not exactly the same as what happened during the major episode of galaxy formation in that the material which TDGs are made from is ``recycled\", as it was already part of a galaxy.  In particular, the presence of metals, both in gas and as dust, facilitates the cooling of the gas and the formation of H$_2$ molecules. Nevertheless, both for TDGs and in the early universe, the galaxy formation process involves clouds of atomic hydrogen (HI) gas gradually condensing through their own gravity, becoming progressively denser, fragmenting, forming molecular gas from the atomic material, and then forming stars.  How this occurs in detail at high redshift is unknown and is one reason for studying TDGs.  Perhaps the least well known of these processes is the transformation of atomic into molecular gas because of the difficulty of observing molecular gas in very low-metallicity environments (e.g. Taylor et al. 1998). Because TDGs condense from matter taken from the outer disks of spiral galaxies, the metallicity of the gas they contain is typically only slightly subsolar as opposed to highly subsolar for small dwarf galaxies (Duc et al. 2000). The metallicity dependent CO lines can thus be used as a probe of the molecular gas content as in spiral galaxies. The study of TDGs influences three areas of astronomy: star formation, dark matter, and galaxy formation.  Much of the material described here has been published in Braine et al. (2000, 2001; hereafter Papers I and II.). Readers are referred to these articles for the lengthy details of the sample as well as many figures showing the CO spectra.  Here we focus on the results and the most recent detection.  \\begin{figure}[h] \\psfig{figure=braine1.fig1.ps,width=120mm,angle=270.0} \\caption{R band image of NGC 4694 and VCC 2062 from Koopman et al. (2001) with the CO(1--0) and HI spectra of VCC 2062 inset.  The angular resolutions are 21\\arcsec\\ and 38\\arcsec\\ respectively.} \\vspace*{-0.5cm} \\end{figure} ",
        "conclusions": ""
    },
    "0402/astro-ph0402054_arXiv.txt": {
        "abstract": "{ We present here the results of a search for new microquasars at low galactic latitudes, based on a cross-identification between the {\\it ROSAT} all sky Bright Source Catalog (RBSC) and the NRAO VLA Sky Survey (NVSS) and follow-up observations. The results obtained up to now suggest that persistent/silent microquasars such as LS~5039 are rare objects in our Galaxy, and indicate that future deeper surveys, and harder than the RBSC in X-rays, will play a fundamental role in order to discover them. }    \\addkeyword{Radio continuum: stars} \\addkeyword{X-rays: binaries}  \\suppressfulladdresses  \\begin{document} ",
        "introduction": "\\label{sec:intro}  Microquasars are X-ray binary systems with the ability to generate relativistic jets (see Mirabel \\& Rodr\\'{\\i}guez 1999 and Fender 2004 for detailed reviews on the topic). These sources mimic, on smaller scales, many of the phenomena seen in AGNs and quasars, but with timescales several orders of magnitude shorter (since $t$ scales with $M$). This property allows us to study, in a few minutes, the accretion/ejection processes that take place near galactic compact objects. Unfortunately, the population of known microquasars is still very small with merely around 15 cases. Therefore, it is worth to search for new microquasars in order to increase the known population to allow meaningful statistical studies. ",
        "conclusions": "\\label{sec:summary}  After a detailed analysis of all the follow-up observations, we show in Table~\\ref{table:summary} a summary of the obtained results. The first two sources, 1RXS J001442.2+580201 and 1RXS J013106.4+612035, show featureless spectra indicative of an extragalactic nature. Nevertheless, the possibility of having highly reddened stars cannot be completely excluded (see Mart\\'{\\i} et~al.\\@ 2004 for details). 1RXS J042201.0+485610 is a Seyfert~1 galaxy with broad Hydrogen emission lines and $z=0.114$. The source 1RXS J062148.1+174736 is probably an extragalactic object due to the featureless spectrum and extended nature of the optical counterpart. 1RXS J072259.5$-$073131 shows properties common to BL Lac objects, while 1RXS J072418.3$-$071508 is an already identified Flat Spectrum Radio Quasar (FSRQ) with $z=0.270$.   Assuming that none of our candidates is galactic, it appears that the population of new and persistent microquasars is not very numerous in the Galaxy. The corresponding density of new (bright) Cygnus~X-3 and (faint) LS~3039-like system is constrained to be $\\la 1.1 \\times 10^{-12}$~pc$^{-3}$ and $\\la 5.6 \\times 10^{-11}$~pc$^{-3}$, respectively. Although we plan to expand our cross-identification studies to $5^{\\circ} \\leq |b| \\leq 10^{\\circ}$, the basic limitation of the RBSC low-energy range, where X-ray photons are highly absorbed, will persist. Therefore, sensitive surveys in hard X-rays and $\\gamma$-rays, such as the current {\\it INTEGRAL} Galactic Plane Survey or the planned {\\it EXIST} mission, will play a fundamental role in order to reveal the {\\it real} population of persistent microquasars in the Galaxy."
    },
    "0402/astro-ph0402438_arXiv.txt": {
        "abstract": "s{ A derivation of the $\\gamma\\gamma\\rightarrow$ e$^+$ e$^-$ optical depth for $\\gamma$ rays produced in a comoving spherical emitting region is presented.  Employing a simplified expression for the $\\gamma\\gamma$ absorption cross section, analytic expressions for the minimum Doppler factor implied by the requirement of $\\gamma$-ray transparency are derived for a broken power-law spectrum of target photons which are isotropically distributed in the comoving frame. Application to specific systems is illustrated.  } ",
        "introduction": "One particularly powerful probe of relativistic motions of AGN and GRB jets, as revealed by the EGRET instrument on the {\\it Compton Observatory}, is the use of $\\gamma$-ray observations to infer minimum Doppler factors of the radiating plasma. The basic idea is simple: the measured FWHM variability time scale $t_{var}$ of a blazar flare or GRB pulse implies a maximum comoving radius of the emitting region from causality considerations. The measured flux and redshift implies the corresponding density of photons which provide targets for $\\gamma\\gamma\\rightarrow$ e$^+$ e$^-$ pair production attenuation of $\\gamma$ rays.  The requirement that the emitting region have small optical depth for observed $\\gamma$ rays places a lower limit on the Doppler factor $\\delta$ of the emitting region. In many cases, such arguments indicate that the radiating plasma in blazar and GRB jets must be relativistic\\cite{dg95,ls01}.  The dramatic improvements in sensitivity of the upcoming {\\it GLAST} mission and the ground-based imaging air Cherenkov telescopes {\\it VERITAS} and {\\it HESS} over previous instruments offer the opportunity to place better limits on $\\delta$, to monitor changes of $\\delta$ in a given source, and to compare $\\delta$ between members of different source classes, e.g., BL Lac objects and flat spectrum radio quasars (FSRQs), and X-ray rich, short duration, and classical GRBs. Because {\\it GLAST} is a scanning mission, blazar flares can be correlated with radio outflows to infer the locations of the sites of $\\gamma$-ray emission.  The $\\gamma$-ray observations, coupled with correlated multifrequency data, will provide knowledge of the jet-disk connection, jet dynamics, and radiation fields in the vicinity of the jet. For these reasons, it seems appropriate to revisit the problem of $\\gamma$-ray photoabsorption in relativistic jets. ",
        "conclusions": "The results here are in essential agreement with previous treatments, taking into account the different approximations made for the cross sections\\cite{dg95,ls01}. Note an implicit co-spatial assumption in the derivation, namely that the $\\gamma$-rays are formed in the same region as the lower energy target photons. Without this assumption, only much smaller values of minimum Doppler factors can be confidently asserted, which depend more on observations at MeV energies rather than at GeV or TeV energies\\cite{dg95a}. To demonstrate the reliability of the co-spatial assumption, and to measure the spectrum of photons that attenuate the $\\gamma$ rays, {\\it GLAST} and ground-based $\\gamma$-ray observations of blazars and GRBs should be correlated with observations made by X-ray/soft $\\gamma$-ray detectors such as {\\it INTEGRAL}, {\\it Chandra}, {\\it XMM Newton} and {\\it Swift}."
    },
    "0402/astro-ph0402112_arXiv.txt": {
        "abstract": " ",
        "introduction": "In Big Bang Nucleosynthesis free neutrons and protons fuse to form gradually heavier nuclei via nuclear and weak reactions.  In the Standard BBN model the universe is homogeneous and isotropic.  The baryon-to-photon ratio $\\eta$ is a variable in the SBBN model.  Figure (1) shows the mass fraction of $ ^{4}$He and the log of the abundance ratios of deuterium to hydrogen and $ ^{7}$Li to hydrogen for the SBBN case, wherein they are functions of $\\eta$.  These results can be compared with observations of primordial abundances.  The $ ^{4}$He mass fraction has been measured as low as 0.228 \\cite{oss97} and as high as 0.248 \\cite{it03}. Kirkman et al \\cite{ktsol03} have measured [d/H] = $2.78_{-0.38}^{+0.44} \\times 10^{-5}$ while Ryan et al \\cite{rbofn00} have measured [$ ^{7}$Li/H] = $1.23_{-0.32}^{+0.68} \\times 10^{-10}$.  [d/H] corresponds to $\\eta = ( 5.6 - 6.6 ) \\times 10^{-10}$ while [$ ^{7}$Li/H] to a different range, $\\eta = ( 1.6 - 4.1 ) \\times 10^{-10}$.  \\begin{figure}[t] \\begin{center} \\includegraphics[height=22pc]{Fig01.ps} \\end{center} \\caption{SBBN:  $ ^{4}$He, [d/H], and [$ ^{7}$Li/H] as functions of $\\eta$. Constraints are in color.} \\end{figure} ",
        "conclusions": "The existence of a region of concordance between [d/H] and $ ^{4}$He constraints in IBBN models is independent of the geometry of the model.  In IBBN models the $ ^{7}$Li constraints still need a depletion factor in order to be in concordance with the other constraints.  A factor of 3.4 can bring all the constraints in concordance for both SBBN and small size IBBN models. Larger size IBBN models alone can satisfy the constraints with a factor of 6.1.  A future article will discuss the influence of neutron diffusion on final abundances in much greater detail.  The cases where $ ^{7}$Be production is greatly increased and when neutron diffusion and BBN coincide will be discussed as those cases determine the sizes of the concordance regions.  The influence of other details such as proton and isotope diffusion and neutrino degeneracy \\cite{kajino02} also needs to determined.  Models for $ ^{7}$Li depletion will be kept in mind to compare with the depletion factors determined in this article."
    },
    "0402/astro-ph0402324_arXiv.txt": {
        "abstract": "Questions such as whether we live in a spatially finite universe, and what its shape and size may be, are among the fundamental open problems that high precision modern cosmology needs to resolve. These questions go beyond the scope of general relativity (GR), since as a (local) metrical theory GR leaves the global topology of the universe undetermined.  Despite our present-day inability to \\emph{predict\\/} the topology of the universe, given the wealth of increasingly accurate astro-cosmological observations it is expected that we should be able to \\emph{detect} it. An overview of basic features of cosmic topology, the main methods for its detection, and observational constraints on detectability are briefly presented. Recent theoretical and observational results related to cosmic topology are also discussed. ",
        "introduction": "\\label{sec:intro}  Is the space where we live finite or infinite? The popular ancient Greek finite-world response, widely accepted in medieval Europe, is at a first sight open to a devastating objection: in being finite the world must have a limiting boundary. But this is impossible, because a boundary can only separate one part of the space from another: why not redefine the universe to include that other part? In this way a common-sense response to the above old cosmological question is that the universe has to be infinite otherwise something else would have to exist beyond its limits.  This answer seems to be obvious and needing no further proof or explanation. However, in mathematics it is known that there are compact spaces (finite) with no boundary. They are called closed spaces. Therefore, our universe can well be spatially closed (topologically) with nothing else beyond its 'spatial limits'. This may be difficult to visualize because we are used to viewing from 'outside' objects which are embedded in our regular $3$--dimensional space. But there is no need to exist any region beyond the spatial extent of the universe.  Of course, one might still ask what is outside such a closed universe. But the underlying assumption behind this question is that the ultimate physical reality is an infinite Euclidean space of some dimension, and nature needs not to adhere to this theoretical embedding framework. It is perfectly acceptable for our $3$--space not to be embedded in any higher-dimensional space with no physical grounds.  Whether the universe is spatially finite and what is its size and shape are among the fundamental open problems that high precision modern cosmology seeks to resolve. These questions of topological nature have become particularly topical, given the wealth of increasingly accurate astro-cosmological observations, especially the recent observations of the cosmic microwave background radiation (CMBR)~\\cite{WMAP}. An important point in the search for answers to these questions is that as a (local) metrical theory general relativity (GR) leaves the global topology of the universe undetermined. Despite this inability to \\emph{predict\\/} the topology of the universe we should be able to devise strategies and methods to \\emph{detect} it by using data from astro-cosmological observations.  The aim of the article is to give a brief review of the main points on cosmic topology addressed in the talk delivered by one of us (MJR) in the XXIV Brazilian National Meeting on Particles and Fields, and discuss some recent results in the field. % The outline of our paper is as follows. In section~\\ref{sec:origin} we discuss how the cosmic topology issue arises in the context of the standard cosmology, and what are the main observational consequences of a nontrivial topology for the spatial section of the universe. In section~\\ref{sec:statmethods} we review the two main statistical methods to detect cosmic topology from the distribution of discrete cosmic sources. In section~\\ref{sec:CinSky} we describe the search for circles in the sky, an important method which has been devised for the detection of cosmic topology from CMBR. In section~\\ref{sec:detect} we discuss the detectability of cosmic topology and present examples on how one can decide whether a given topology is detectable or not according to recent observations. Finally, in section~\\ref{sec:news+remarks} we briefly discuss recent results on cosmic topology, and present some concluding remarks. ",
        "conclusions": "\\label{sec:news+remarks}  In this section we shall briefly discuss some recent results and advances in the search for the shape of the universe, which have not been treated in the previous sections. We also point out some problems, which we understand as important to be satisfactorily dealt with in order to make further progress in cosmic topology.  One of the intriguing results from the analysis of WMAP data is the considerably low value of the CMBR quadrupole and octopole moments, compared with that predicted by the infinite flat $\\Lambda$CDM model. Another noteworthy feature is that, according to WMAP data analysis by Tegmark \\emph{et al.\\/}, both the quadrupole and the octopole moments have a common preferred spatial axis along which the power is suppressed \\footnote{Incidentally, it was the fitting to the observed low values of the quadrupole and the octopole moments of the CMB temperature fluctuations that motivated Poincar\\'e dodecahedron space topology% ~\\cite{Poincare}, which according to 'cirlces in the sky' plus WMAP analysis is excluded~\\cite{CSSK2003}. Nevertheless, the Poincar\\'e dodecahedron space proposal was an important step in cosmic topology to the extent that for the first time a possible nontrivial cosmic topology was tested against accurate CMBR data.}.  This alignment of the low multipole moments has been suggested as an indication of a direction along which a possible shortest closed geodesics (characteristic of multiply connected spaces) of the universe may be~\\cite{OCTZH2003}. Motivated by this as well as the above anomalies, test using $S$-statistics~\\cite{OCSS1996} and matched circles furnished no evidence of a nontrivial topology with diametrically opposed pairs of correlated circles% ~\\cite{OCTZH2003}. It should be noticed, however, that these results do no rule out most multiply connected universe models because $S$-statistics is a method sensitive only to Euclidean translations, while the search for circles in the sky, which is, in principle, appropriate to detect any topology, was performed in a limited three-parameter version, which again is only suitable to detect translations.  At a theoretical level, although strongly motivated by high precision data from WMAP, it has been shown that if a \\emph{very} nearly flat universe has a detectable nontrivial topology, then it will exhibit the generic local shape of (topologically) $\\mathbb{R}^2 \\times \\mathbb{S}^1$% \\footnote{Or more rarely $\\mathbb{R} \\times \\mathbb{T}^2\\,$.}, irrespective of its global shape~\\cite{MGRT2003a}. In this case, {}from WMAP and SDSS the data analysis, which indicates that $\\Omega_0 \\approx 1$~\\cite{Tegmark-et-al2003}, one has that if the universe has a detectable topology, it is very likely that it has a preferred direction, which in turn is in agreement with the observed alignement of the quadrupole and octopole moments of the CMBR anisotropies. In this context, it is relevant to check whether a similar alignment of higher order multipole ($\\ell>3$) takes place in order to reinforce a possible nontrivial local shape of our $3$--space. In this connection it is worth mentioning that Hajian and Souradeep~\\cite{HajianSouradeep2003a,HajianSouradeep2003b} have recently suggested a set of indicators $\\kappa_{\\ell}$ ($\\ell=1,2,3,...$) which for non-zero values indicate and quantify statistical anisotropy in a CMBR map. Although $\\kappa_{\\ell}$ can be potentially used to discriminate between different cosmic topology candidates, they give no information about the directions along which the isotropy may be violated, and therefore other indicators should be devised to extract anisotropy directions from CBMR maps.  In ref.~\\cite{MGRT2003a} it has also been shown that in a very nearly flat universe with detectable nontrivial topology, the observable (detectable) isometries will behave nearly like translations. Perhaps if one use Euclidean space to locally approximate a nearly flat universe with detectable topology, the detectable isometries can be approximated by Euclidean isometries, and since these isometries are not translations, they have to be screw motions. As a consequence, an approximate local shape of a nearly flat universe with detectable topology would look like a \\emph{twisted} cylinder, i.e. a flat manifold whose covering group is generated by a screw motion. Work toward a proof of this conjecture is being carried out by our research group.  Before closing this overview we mention that the study of the topological signature (possibly) encoded in CMBR maps as well as to what extent the cosmic topology CMBR detection methods are robust against distinct observational effects such as, e.g., Suchs-Wolfe and the thickness of the LSS effects, will benefit greatly from accurate simulations of these maps in the context of the FLRW models with multiply connected spatial sections. A first step in this direction has been achieved by Riazuelo \\emph{et al.\\/}~\\cite{RULW2002}, with special emphasis on the effect of the topology in the suppression of the low multipole moments. Along this line it is worth studying through computer-aided simulations the effect of a nontrivial cosmic topology on the nearly alignments of the quadrupole and the octopole moments (spatial axis along which the power is suppressed).  To conclude, cosmic topology is at present a very active research area with a number of important problems, ranging from how the characterization of the local shape of the universe may observationally be encoded in CMBR maps, to the development of more efficient computationally searches for matching circles, taking into account possible restrictions on the detectable isometries, and thereby confining the parameter space which realistic `circles in the sky' searches need to concentrate on. It is also of considerable interest the search for the statistical anisotropy one can expect from a universe with non-trivial space topology. Finally, it is important not to forget that there are almost flat (spherical and hyperbolic) universes, whose spatial topologies are undetectable in the light of current observations with the available methods, and our universe can well have one of such topologies. In this case we have to devise new methods and strategies to detect the topology of the universe.   \\bigskip \\noindent{\\bf Acknowledgments} \\medskip  We thank CNPq and FAPESP (contract 02/12328-6) for the grants under which this work was carried out. We also thank A.A.F. Teixeira and B. Mota for the reading of the manuscript and indication of relevant misprints and omissions."
    },
    "0402/hep-ph0402168_arXiv.txt": {
        "abstract": " ",
        "introduction": "A few years ago, in 1998, a novel idea was proposed according to which the so-called {\\it hierarchy problem} -- in other words, our difficulty in answering the question of why the characteristic scale of gravity, $M_P\\sim 10^{19}$~GeV, is 16 orders of magnitude larger than the Electro-Weak scale, $M_{EW} \\sim 1$~TeV -- could be solved by assuming the existence of extra dimensions in the Universe \\cite{ADD,AADD}. The novelty in this idea was that the traditional picture of Planck-length-sized additional spacelike dimensions ($\\ell_P \\simeq 10^{-33}$ cm) was abandoned, and the extra dimensions could have a size as large as 1 mm. The upper bound on the size of the proposed {\\it Large Extra Dimensions} actually matched the smallest length scale down to which the force of the gravitational interactions, and thus their $1/r^2$ dependence, had been measured. If extra dimensions of that size did exist, gravitational interactions would have a completely different dependence on $r$ in scales smaller than 1 mm, however no gravitational experiment at that time could rule this out.  On the other hand, electromagnetic, weak and strong forces are indeed sensitive to the existence of extra dimensions. If, for example, gauge bosons were allowed to propagate in the extra-dimensional spacetime, their interactions would be modified beyond any acceptable phenomenological limits unless the size of the extra dimensions was smaller than $10^{-16}$\\,cm. This problem was resolved under the assumption that all particles experiencing this type of interactions, in other words, all ordinary matter, is restricted to live on a (3+1)-dimensional hypersurface, a {\\it 3-brane}, that has a width along the extra dimensions of, at most, the above order. The 3-brane, playing the role of our four-dimensional world, is then embedded in the higher-dimensional spacetime, usually called the {\\it bulk}, in which only gravity can propagate.  In this model, the large volume of the extra dimensions can help us solve, or at least recast, the hierarchy problem: the traditional Planck scale, $M_{P}$, is only an effective energy scale derived from the fundamental higher-dimensional one, $M_*$, through the following relation \\cite{ADD} \\begin{equation} M_P^2 \\sim M_*^{2+n}\\,R^n\\,. \\end{equation} In the above, it has been assumed that each one of the extra spacelike, compact dimensions has the same size $R$. From the above, it becomes clear that, if the volume of the compact space, $V \\sim R^n$, is large, i.e if $R \\gg \\ell_P$, then the  $(4+n)$-dimensional Planck mass, $M_*$, will be much lower than the 4-dimensional one, $M_P$. If one chooses $M_*=M_{EW}$, then the above expression provides a relation between the scale of gravity and the scale of particle interactions.  The above idea makes use of geometrical features of the extra, compact space in order to connect two completely different energy scales. The same goal was achieved by an alternative, but similar idea, that was proposed a year later \\cite{RS} (for some early works on brane-world models and their implications, see Refs. \\refcite{Akama}-\\refcite{Lykken}). According to this alternative proposal, the magnitude of the effective energy scale on the 3-brane, taken to be of the order of the Electro-Weak scale, follows from the fundamental higher-dimensional scale $M_P$ after being suppressed by an exponential factor involving the distance of our observable brane from a hidden brane. In both models, however, the complete resolution of the hierarchy problem would demand also an explanation of why the volume of the compact space, or the inter-brane distance, has the value that leads to the observed ratio of $M_P/M_{EW}$.  In this review, we will concentrate on the scenario with Large Extra Dimensions. A number of experiments and theoretical analyses have tried over the years to put bounds on the size of extra, compact dimensions, and the produced bounds are constantly updated. In the regime $r \\ll R$, the extra dimensions `open up' and Newton's law for the gravitational interactions is modified assuming a $1/r^{n+2}$ dependence on the radial separation between two massive particles.  Torsion-balance experiments which measure the gravitational inverse-square law at short scales can provide limits on the size of the extra dimensions, or equivalently on the value of the fundamental scale $M_*$. On the other hand, since gravitons can propagate both in the bulk and on the brane, massive Kaluza-Klein (KK) graviton states can modify both the cross sections of Standard Model particle interactions and astrophysical or cosmological processes. Assuming modifications that are below the current observable limits puts bounds on the mass of the KK gravitons, and consequently, on the size of extra dimensions. In Table 1, we summarize some of those limits. The constraints from collider experiments, although more accurate, are particularly mild, while the cosmological and astrophysical ones are much more stringent; however, they contain large systematic errors. If we ignore these errors, the latter type of constraints exclude by far even the $n=3$ case, while low-gravity models with $M_* \\sim 1$~TeV are still allowed for $n \\geq 4$.  \\begin{table}[t] \\tbl{Current limits on the size of extra, compact dimensions} {\\begin{tabular}{@{}crr@{}} \\toprule Type of Experiment/Analysis &  $M_* \\ge$ \\hspace*{0.5cm} & $M_* \\ge$ \\hspace*{0.5cm}\\\\ \\colrule \\begin{tabular}{l}Collider limits on the production \\\\ of real or virtual KK gravitons \\cite{Delphi,Opal,CDF}\\end{tabular} & 1.45 TeV ($n=2$) & 0.6 TeV ($n=6$)\\\\[4mm] Torsion-balance Experiments\\cite{Hoyle}  & 3.5 TeV ($n=2$) &  \\\\[2mm] Overclosure of the Universe \\cite{Hall} & 8 TeV ($n=2$) & \\\\[2mm] Supernovae cooling rate \\cite{Cullen,Barger,Reddy,Reddy2} & 30 TeV ($n=2$) &  2.5 TeV ($n=3$)\\\\[2mm] Non-thermal production of KK modes \\cite{Pospelov} & 35 TeV ($n=2$) & 3 TeV ($n=6$)\\\\[2mm] Diffuse gamma-ray background \\cite{Hall,HR1,Hann} & 110 TeV ($n=2$) & 5 TeV ($n=3$)\\\\[2mm] Thermal production of KK modes~\\cite{Hann} & 167 TeV ($n=2$) &  1.5 TeV ($n=5$)\\\\[2mm] Neutron star core halo \\cite{HR2} & 500 TeV ($n=2$) & 30 TeV ($n=3$) \\\\[2mm] Neutron star surface temperature~\\cite{HR2}& 1700 TeV ($n=2$) & 60 TeV ($n=3$)\\\\[2mm] BH absence in neutrino cosmic rays \\cite{Feng} & & 1-1.4 TeV ($n \\geq 5$)\\\\ \\botrule \\end{tabular}} \\end{table}   If indeed present, the extra dimensions will inevitably change our notion for the universe. The introduction of extra dimensions affects both gravitational interactions and particle physics phenomenology, and leads to modifications in standard cosmology. Already existing theories would need to be extended or modified in order to accommodate the effects resulting from the presence of extra dimensions (for an incomplete list of works on the cosmological and phenomenological implications in theories with large extra dimensions, see Refs. \\refcite{DDG}-\\refcite{Paul}). Similarly, the properties and physics of black holes are also bound to change in the context of a higher-dimensional theory.  As in the four-dimensional case, it seems natural to assume that, when matter trapped on the brane undergoes gravitational collapse, a black hole is formed which is centered on the brane and extends along the extra dimensions. If the horizon of the formed black hole is much larger than the size of the extra dimensions, $r_H \\gg R$, the produced black hole is effectively a four-dimensional object. If, however, $r_H \\ll R$, then this small black hole is virtually a higher-dimensional object that is completely submerged into the extra-dimensional spacetime. As we will see, these small black holes have significantly modified properties compared to a four-dimensional black hole with exactly the same mass $M_{BH}$: for example, they are larger, colder and thus live longer compared to their four-dimensional analogues. Another striking consequence of the introduction of extra dimensions is that, by lowering the Planck scale $M_*$ closer to the Electro-Weak scale, the idea of the production of miniature black holes during high-energy scattering processes, with trans-Planckian center-of-mass energy $\\sqrt{s} \\gg M_*$, now becomes more realistic. Theoretical arguments have shown that the presence of the extra dimensions facilitates further the production of black holes during such collisions by increasing the production cross-section, thus leading to striking consequences for the high-energy interactions of elementary particles either at colliders or at cosmic rays.  The produced black holes are characterized by a non-vanishing temperature $T_H$, whose value is inversely proportional to the horizon radius. They decay by the emission of Hawking radiation, i.e. emission of elementary particles with rest mass smaller than $T_H$. This is expected to be their most prominent observable signature with a characteristic thermal radiation spectrum and an almost blackbody profile. The non-trivial metric in the region exterior to the horizon of the black hole creates an effective potential barrier which backscatters a part of the outgoing radiation back into the black hole. The amount of radiation that finally reaches the observer at infinity depends on the {\\it energy} of the emitted particle, its {\\it spin} and the {\\it dimensionality} of spacetime. The dependence on all the aforementioned parameters is encoded into the expression of a filtering function, the `greybody factor' $\\sigma_{n}^{(s)}(\\omega)$, which is present in the radiation spectrum. The greybody factors can be important experimentally since they modify the spectrum in the low- and intermediate-energy regime, where most particles are produced, thus altering the characteristic spectrum by which we hope to identify a `BH event'. In addition, as we will explain later in this review, by studying the Hawking radiation emitted by this type of small black holes, one would be able to `read' the total number of dimensions that exist in nature.  A higher-dimensional black hole emits radiation both in the bulk and on the brane. According to the assumptions of the theory with Large Extra Dimensions, only gravitons, and possibly scalar fields, can propagate in the bulk and thus, these are the only types of fields allowed to be emitted in the bulk during the Hawking evaporation phase. On the other hand, the emission on the brane can take the form of scalar Higgs particles, fermions and gauge bosons. From the perspective of the brane observer, the radiation emitted in the bulk will be a missing energy signal, while radiation on the brane may lead to experimental detection of Hawking radiation and thus of the production of small black holes. Nevertheless, in order to have a clear picture of the characteristics of the radiation spectrum on the brane, it is important to know exactly how much energy is lost in the bulk.  As we mentioned above, the greybody factor depends on the dimensionality of spacetime; it also depends on whether the emitted particle is brane-localized or free to propagate in the bulk. The greybody factor is actually the outgoing transmission cross-section associated with propagation in the aforementioned gravitational background, however, due to the thermal character of the radiation spectrum, it is equal to the incoming absorption cross-section \\cite{Birrell}. Therefore, all we need to do is to solve the equation of motion of a particle incident on the background metric that describes the black hole, either higher-dimensional or four-dimensional. After the absorption coefficient is computed, the corresponding cross-section, and thus the greybody factor, can easily follow.  We need to stress here that the above semiclassical calculation of Hawking emission is only reliable when the energy of the emitted particle is small compared to the black hole mass, $\\omega \\ll M_{BH}$, since only in this case is it correct to neglect the back reaction to the metric during the emission process. This in turn requires that the Hawking temperature obeys the relation $T_H \\ll M_{BH}$, which is equivalent to demanding that the black hole mass $M_{BH} \\gg M_*$. As the decay proceeds and the mass of the black hole decreases, inevitably this condition will break down during the final stages of the evaporation process. Nevertheless, for black holes of initial mass much larger than $M_*$ most of the evaporation process is within the semi-classical regime.  We will start this article by reviewing, in Section 2, the existing literature on the creation of black holes during high-energy particle collisions, both in four-dimensional and higher-dimensional spacetimes. In Section 3, we will discuss the properties of the produced higher-dimensional black holes, namely the horizon radius, the temperature and their life-time, and point out the differences between them and their four-dimensional analogues. The physics that governs the evaporation of these higher-dimensional objects, through the emission of Hawking radiation, is covered in Section 4. Section 5 focuses on the emission of Hawking radiation directly on the brane: we start with the derivation of the {\\it master} equation for the propagation of fields with arbitrary spin in the induced-on-the-brane black hole background, and then we present all existing results in the literature for the emission of scalars, fermions and gauge bosons during the {\\it spin-down} and {\\it Schwarzschild} phases of the life of the black hole; both analytical and numerical results for the greybody factors and radiation spectra are presented as well as exact results for the number and type of fields emitted on the brane as a function of the dimensionality of spacetime. Section 6 deals with the emission of Hawking radiation in the bulk and the question of the amount of the missing energy: analytical and numerical results on the emission of bulk scalar fields, including greybody factors and radiation spectra, are reviewed, and the ratio of the missing energy over the `visible' one emitted on the brane is presented for the Schwarzschild phase and for different values of the number of extra dimensions. Our conclusions are summarized in Section~7. ",
        "conclusions": "The proposal of the existence of extra dimensions in the Universe has opened many different pathways that, in principle, lead to important modifications in the four-dimensional cosmology, particle physics phenomenology and black hole physics. The recent revival of the idea of the existence of extra space-like dimensions, that can have an almost macroscopic size or be even non-compact, introduced a novel feature in the theory: the scale at which the gravity becomes strong may be much lower than the traditional four-dimensional Planck mass $M_P$. This has led to the exciting prospect that high-energy collisions between elementary particles, that will take place at next-generation ground-based accelerators or have been already taken place at the atmosphere of the Earth, can probe the energy regime of quantum gravity. Some of the most striking consequences would be the possible creation of mini black holes, or even strings and D-branes, as the products of a high-energy collision of particles with center-of-mass energies at just a few times the new scale of gravity $M_*$.  In Section 2, we have reviewed the existing results in the literature regarding the possibility of the creation of mini black holes during high-energy collisions. Well-known, four-dimensional analyses have been generalized to cover the case where the colliding particles propagate in a higher-dimensional spacetime. Some of the new studies have shown that, in the case of head-on collisions, the mass of the produced black hole gets suppressed, as the number of extra dimensions increases, while some recent complimentary results indicate that the emission of gravitational waves during the collision is actually suppressed when $n$ increases: this obviously brings the two types of results in an apparent disagreement unless we accept the possibility that a significant amount of the energy, lost during the head-on collision, is emitted in a form different from that of gravitational radiation. On the other hand, for collisions with a non-vanishing impact parameter, which are the most likely to take place, the black hole production cross-section is in fact enhanced with the number of extra dimensions. Putting the above into the framework of a realistic collision between composite particles have led to large estimates, by new physics standards, for the corresponding black hole production cross section, either at the LHC or at the atmosphere of the Earth.  The effect of extra dimensions is not restricted to the production of mini black holes; the properties of the produced, higher-dimensional black holes are also modified. In Section 3, we have discussed some of those properties, namely, the horizon radius, temperature and lifetime, all of which are crucial parameters for the successful production and detection of these elusive, up to now, objects. As we saw, the horizon radius of a ($4+n$)-dimensional black hole is many orders of magnitude larger than the one of a four-dimensional black hole with the same mass, which simply means that a mass $M$ needs to be compacted less in a higher-dimensional spacetime to create a black hole. The temperature of these small black holes, in turn, comes out to be lower than in four-dimensions, which means that the emission rate of Hawking radiation is smaller and their lifetime longer. What is most favourable for the possibility of detecting these objects is the fact that, for a black hole with a mass $M_{BH}=5$ TeV or so, the Hawking radiation spectrum reaches its peak at energies close to the temperature of the black hole which is of the order of 100-600 GeV (for $n=1,...,7$); this is exactly the energy regime that present and next-generation collider experiments can probe. Such a small black hole will be extremely short-lived, i.e. $\\tau \\simeq 10^{-26}$ sec, nevertheless, the proximity of the evaporating black holes to our detectors increases significantly the possibility of their detection.  These mini black holes, upon implementation of quantum effects, emit Hawking radiation into the higher-dimensional spacetime in the form of both bulk and brane modes. The generalization of the four-dimensional expressions for the emission rates in the case of  a higher-dimensional spacetime is straightforward, nevertheless, the exact expression of the greybody factors for different types of fields propagating in a higher-dimensional background was, until recently, unknown. As we explained in Section 4, the greybody factors encode valuable information for the background around the emitting black hole and depend on the energy of the emitted particle, its spin and the dimensionality of spacetime. This means that the presence of the greybody factor in the radiation spectrum will cause the modification of the low-energy emission rate from the high-energy one, and will lead to different emissivities for particles with different spin. Moreover, both the number and the type of particles emitted will depend on the number of extra dimensions that exist in nature, a feature that may possibly lead to the determination of the dimensionality of spacetime upon detection of Hawking radiation.  Having concluded that the implementation of the greybody factor in the radiation spectrum is imperative in order to derive accurate estimates for the emission spectrum of the black hole, we moved on, in Section 5, to derive a master equation for the propagation of fields with different spin on the black-hole background induced on the four-dimensional brane. By solving this master equation, one can compute the transmission cross-section, in other words the greybody factor, for brane-localized modes emitted by the black hole. This type of emission during the spin-down phase of a black hole has been only partially studied: radiation spectra for fields with different spin $s$ have been computed only for the case of $n=1$ in the limit of low energy and low angular momentum. On the contrary, the spherically-symmetric Schwarzschild phase has been thoroughly investigated. Both analytical and numerical methods were used, with the former one leading to analytical, but low-energy-only-valid, expressions, and the latter one yielding exact numerical results valid at all energy regimes. We were thus able to compare the radiation spectra for different types of fields and different number of extra dimensions. For all types of particles, the total emissivities are greatly enhanced with the number of extra dimensions, with the enhancement reaching orders of magnitude for large values of $n$. As the increase in the emission rate depends strongly on the spin, the relative emissivities for particles with different spin are also strongly $n$-dependent: while scalar fields remain the preferred type of particle emitted by the black hole for low and intermediate values of $n$, they are outnumbered by the gauge bosons for large values of $n$, with the fermions being the least effective channel during the emission.  The emission of brane-localized modes is undoubtly the most phenomenologically interesting effect since it involves Standard Model particles that can be easily detected during experiments. On the other hand, the emission of bulk modes will be only perceived as a missing energy signal by the observer on the brane. Nevertheless, the amount of energy lost in the bulk is crucially important as it determines the remaining available energy for emission on the brane. The details of the emission of bulk scalar modes during the spin-down phase have been studied in an analytic but qualitative way, and no quantitative results are available. For the Schwarzschild phase, as we saw in section 6, the study has been completed. The greybody factor and emission rates have been calculated and the latter were shown to increase again with the dimensionality of spacetime. The extremely important question of the bulk-to-brane energy emissivity has been answered only for the scalar channel and for the Schwarzschild phase: the amount of energy spent by the black hole for the emission of bulk scalar modes is always smaller than the one for brane modes; nevertheless, the amount of energy lost in the bulk must always be taken into account especially for large values of $n$ when the energies spent in the brane and bulk channel become comparable. No results for the bulk-to-brane emissivity for gravitons have been yet derived.  Although the possibility of the production and evaporation of mini black holes at the LHC is an exciting prospect, this will be possible only in the case where the fundamental scale of gravity $M_*$ is indeed very close to 1 TeV. Nevertheless, there is absolutely no guarantee for that, and the only argument in favour of this particular value is the possible resolution of the hierarchy problem. If $M_*$ is larger than 1 TeV, even by one order of magnitude, the probability of the production of black holes at the LHC vanishes (although we might still witness these type of effects in cosmic ray particles). Nevertheless, all the analytical and numerical results presented in this review have $M_*$ as an independent parameter, and are therefore valid for all values of $M_*$. If this scale is pushed upwards by various constraints, the derived results will still be applicable for the production and evaporation of black holes at the new energy regime $\\sqrt{s} \\geq M_*$.  We would also like to stress that the results for the radiation spectra reviewed here refer to individual degrees of freedom and not to elementary particles, like electrons or quarks, which contain more than one polarization. For the number of elementary particles produced, and the energy they carry, one has to use a Black Hole Event Generator \\cite{dl,HRW}. This simulates both the production and decay of small black holes at hadronic colliders and provides estimates for the number and spectra of the different types of elementary particles produced.  In this review, we concentrated on theories postulating the existence of Large Extra Dimensions, and we studied the properties of small black holes that live in a spacetime with $D-1$ flat spacelike dimensions. Our analysis is definitely not valid in highly curved spacetimes, like the ones in five-dimensional warped models~\\cite{RS}. In that case, attempts to  construct a gravitational background that would reduce to a black-hole line-element on the brane while being well defined away from it have failed; numerical methods have instead found five-dimensional localized black holes that do not necessarily have a black hole line-element projection onto the brane (see Refs. \\refcite{CHR}-\\refcite{Kudoh} for some relevant works). Nevertheless, in the case of very small bulk cosmological constant, the `warping' of the five-dimensional spacetime, parametrised by the inverse AdS radius, $\\lambda^{-1}$, is small and the extra spacetime can be considered as an almost flat one. Alternatively, if the horizon radius $r_H$ is much smaller than the AdS radius, no matter what the value of $\\lambda$ is, then the black hole cannot perceive the warping of the fifth dimension. In either case, we may model black holes with $r_H \\ll \\lambda$ arising in a warped spacetime as black holes living in a five-dimensional, flat spacetime. Under this assumption, all the results presented in this review hold equally well also for the `warped' black holes.  As a concluding remark, we would like to stress once again that the detection of signatures of possible black hole production events during high-energy collisions would be a revolutionary development both for particle physics phenomenology and gravitational physics. Any observational signal of this type would automatically prove the existence of extra dimensions and of a fundamental theory of all forces, with a low energy scale, valid in a higher-dimensional spacetime. The detection of Hawking radiation emitted by these small, higher-dimensional black holes is the most direct evidence for the production and evaporation of those objects. The radiation spectrum of a decaying black hole can also reveal the exact dimensionality of spacetime as both the amount and  type of the emitted radiation strongly depend on it. What is also exciting is that a black hole can emit all types of particles that exist in nature, independently of their spin, charge, quantum numbers or interaction properties, as long as their rest mass is smaller than the black hole temperature; therefore, long-sought but yet undiscovered particles, like the Higgs fields or supersymmetric particles, might indeed make their appearance in the decay spectrum of a black hole. The launch of the LHC, or of any other future experiment able to probe even higher energy regimes, deserves to be awaited with great expectations indeed."
    },
    "0402/astro-ph0402031_arXiv.txt": {
        "abstract": "Detailed evidence on the system AX J0049.4-7323 is presented here to show how the passage of the neutron star in the binary system disrupts the circumstellar disk of the mass donor Be star. A similar effect is noted in three other Be/X-ray binary systems. Together the observational data should provide valuable tools for modelling these complex interactions. ",
        "introduction": "The Be/X-ray systems represent the largest sub-class of massive X-ray binaries.  A survey of the literature reveals that of 96 proposed massive X-ray binary pulsar systems, 57\\% of the identified systems fall within this class of binary.  The orbit of the Be or supergiant star and the compact object, presumably a neutron star, is generally wide and eccentric.  X-ray outbursts are normally associated with the passage of the neutron star close to the circumstellar disk (Okazaki \\& Negueruela, 2001). A recent review of these systems may be found in Coe (2000).  The optical light from a Be/X-ray binary is dominated by the mass donor star in the blue end of the spectrum, but at the red end there is normally a significant contribution from the circumstellar disk. Long term optical observations such as those collected by the MACHO experiment (Alcock et al, 1995) provide valuable insights into the behaviour of the circumstellar disk, and hence into some of the details of the binary interactions within the system. ",
        "conclusions": "Detailed evidence on the system AX J0049.4-7323 has been presented here to show how the passage of the neutron star disrupts the circumstellar disk of the Be star. A similar effect is noted in three other Be/X-ray binary systems. Together the observational data should provide valuable tools for modelling these complex interactions."
    },
    "0402/astro-ph0402207_arXiv.txt": {
        "abstract": "We derived the stellar parameters of a sample of Galactic early-O type stars by analysing their UV and Far-UV spectra from \\textit{FUSE} (905-1187\\AA), \\textit{IUE}, \\textit{HST-STIS} and \\textit{ORFEUS} (1200-2000\\AA). The data have been modeled with spherical, hydrodynamic, line-blanketed, non-LTE synthetic spectra computed with the \\textit{WM-basic} code. We obtain effective temperatures ranging from \\Teff = 41,000~K to 39,000~K for the O3-O4 dwarf stars, and \\Teff = 37,500~K for the only supergiant of the sample (O4 If$^+$). Our values are lower than those from previous empirical calibrations for early-O types by up to 20\\%. The derived luminosities of the dwarf stars are also lower by 6 to 12\\%; however, the luminosity of the supergiant is in agreement with previous calibrations within the error bars. Our results extend the trend found for later-O types in a previous work by Bianchi \\& Garcia. ",
        "introduction": "\\label{s_intro}  Hot massive stars have a great impact on the surrounding interstellar medium (\\textit{ISM}) and play a crucial role in the chemical evolution of galaxies. Their strong ultraviolet radiation is responsible for the ionization of nearby $HII$ regions where their stellar winds blow vast bubbles. Their supersonic wind outflows and the supernova explosion at the end of their evolution transfer energy and momentum to the \\textit{ISM} and disperse the material processed in the stellar interiors, thus setting the conditions for the formation of subsequent generations of stars.  The determination of the physical parameters of massive stars is therefore of great interest, yet complicated. High resolution spectroscopy is needed. Modeling the stellar atmosphere requires to account  for the expanding wind, the non local thermodynamic equilibrium (non-LTE) conditions and the so called line-blanketing that modifies the flux distribution, especially at short wavelengths.  Spectroscopy in the ultraviolet and far ultraviolet ranges (hereafter UV and Far-UV) is a powerful tool to study the winds of massive stars since these spectral regions contain the resonance lines of the most abundant ions in the wind. In this work we analyse spectra from the \\textit{Far Ultraviolet Spectroscopic Explorer} (\\textit{FUSE}) \\citep[]{M00}, covering the 905-1187\\AA\\space region, in conjunction with spectra at longer wavelengths (1200-2000\\AA) from the \\textit{International Ultraviolet Explorer} (\\textit{IUE}), the \\textit{Orbiting Retrievable Far and Extreme Ultraviolet Spectrometers} (\\textit{ORFEUS}) and the \\textit{Space Telescope Imaging Spectrograph} (\\textit{STIS}) aboard the \\textit{Hubble Space Telescope} (\\textit{HST}). The \\textit{FUSE} spectra allows us to uniquely constrain the stellar parameters by adding new ionization stages to those accessible to \\textit{IUE}, \\textit{ORFEUS} and \\textit{HST-STIS} (e.g. \\citet[]{LB00}, \\citet[]{BG02}).  This is the second paper of a series devoted to provide accurate and consistent determination  of the stellar parameters of Galactic massive stars with this method. \\citet[hereafter paper~I]{BG02} studied six mid-O type stars and found effective temperatures lower (by 15-20\\%) than previously determined for the sample stars or calibrated for their spectral types. In this work we perform a similar analysis for early-O type stars. The paper is organized as follows. In Section \\ref{s_data} we provide details about the data and the reduction. In Section \\ref{s_stars} we summarize the relevant information from the literature about the program stars. In Section \\ref{s_morpho} we compare the spectral line characteristics. The stellar parameters are derived in Section \\ref{s_models} by modeling the spectra. In Section \\ref{s_conclusion} the results are discussed. ",
        "conclusions": "\\label{s_conclusion}  We have performed a detailed spectroscopic analysis of seven Galactic early-O type stars and derived their stellar parameters. We found effective temperatures ranging from 41,000~K to 39,000~K for the O3-4 dwarf stars and 37,500~K for the O4 supergiant. The X-ray ionization due to shocks in the wind has been constrained by our modeling, primarily from the \\ion{O}{6}~$\\lambda\\lambda$1031.9,1037.6 resonance doublet, accessible with the \\textit{FUSE} telescope. In the range of \\Teff, \\Mdot~ and \\Rstar~ that we derived for the dwarf stars, the spectral morphology hardly changes with  gravity and the derived value of \\logg is less well constrained than other parameters, but consistent with previous estimates.  All the sample stars display unsaturated \\ion{N}{5}~$\\lambda\\lambda$1238.8,1242.8 profiles that we cannot fit consistently with the rest of the spectrum. In one case at least (HD~96715), the unsaturated profile is due to poor background subtraction in the IUE data (as shown by an \\textit{ORFEUS} spectrum), and this may apply to the other stars as well. In the case of HD~190429A, HD~93205 and HD~168076, the unsaturated profile may also be due to flux from the cooler companion star, whose contribution we estimate to be $\\leq$ 15\\% of the observed flux. A comparison of \\textit{WM-basic} and \\textit{CMFGEN} models calculated with parameters suitable for the early-O type stars is under way (\\citet[]{BGH03}; Bianchi, Garcia, \\& Herald, in preparation), to investigate the effect of a different treatement of the shocks on the \\ion{N}{5}~$\\lambda\\lambda$1238.8,1242.8 doublet.  We compile previous determinations of the stellar parameters for our program stars from the literature in Table \\ref{t_prev}. Our effective temperatures are consistently lower than those derived in other works, with the only exception of the Str\\\"omgren and H$\\beta$ photometric studies  by \\citet[]{KG93} and \\citet[]{SJ95}. \\citet[]{KG93} derived temperatures for the Carina stars of our sample lower than ours by $\\sim$4,000~K on average. \\citet[]{SJ95} assigned a temperature of 28,200~K to HD~190429 (without resolving the system), closer to what we would expect for the secondary than to what we found for HD~190429A (37,500~K). \\citet[]{M75} performed Str\\\"omgren and H$\\beta$ photometry of HD~168076 and derived a temperature higher than ours by 16,000~K. The calibration of \\citet[]{C75}, based on the comparison of the equivalent widths of hydrogen and helium lines with plane-parallel hydrostatic non-LTE model predictions, was used by \\citet[]{CW76} and \\citet[]{C77} for several objects of our sample, yielding temperatures higher than our results by $>$ 10,000~K. Temperatures obtained from fitting Balmer lines and optical \\ion{He}{1} and \\ion{He}{2} lines with plane-parallel non-LTE hydrostatic models are $\\sim$  10\\% \\citep[]{Sal83} to $\\sim$20\\% \\citep[]{Kal92,puls96} higher than ours. \\citet[]{WMBAS} compared the \\textit{IUE} spectrum of HD~93250 to \\textit{WM-basic} models of \\Teff=50,000~K (approximately the same as derived by \\citet[]{puls96}), \\logg=4.0 and \\Rstar=12\\Rsun, with different mass-loss rates around \\Mdot=5.6$\\cdot 10^{-6}$\\Myr. % Their synthetic spectra did not include shocks effects and either produce \\ion{O}{5}~$\\lambda$1371.0 in excess or cannot reproduce the \\ion{N}{4}~$\\lambda\\lambda$1718.0,1718.5 line. In our work we used those lines, together with \\ion{C}{4}~$\\lambda$1169+\\ion{C}{3}~$\\lambda$1176 and \\ion{P}{5}~$\\lambda\\lambda$1118.0,1128.0, to set the upper limit of the effective temperature of the dwarf stars to $\\sim$42,000~K (see Section \\ref{s_models}) and we obtained a consistent fit of all spectral features by including the effects of shocks in the calculations.   In Figure \\ref{tefflum} we compare the effective temperatures and luminosities obtained for the total sample of O3-O7 type stars analysed in this work and in paper~I with previous empirical calibrations (\\citet[]{vacca}, \\citet[]{dejag} and \\citet[]{MPRM03}). Our \\Teff~ values for O3-O6 dwarf stars are lower than these calibrations, the differences ranging between 9,000-10,000~K, 6,000-9,000~K and 4,000-7,000~K respectively. The discrepancy is smaller for the O4-O7 supergiants (6,000-10,000~K, 4,000-6,000~K and 2,000-3,500~K, respectively). The derived luminosities for both dwarf stars and supergiants are also lower, by an average of 7\\% and 10\\% respectively, than both the calibrations of \\citet[]{dejag} and \\citet[]{vacca}. For the O4~If$^+$ supergiant the luminosity is also lower, but the discrepancy is within the error bars.  The lower temperatures from our analysis are due to two main improvements. On the one hand, the use of \\textit{FUSE} data enables the assessment of X-rays from shocks and thus a correct (consistent) derivation of the wind ionization. On the other hand, the inclusion of line-blanketing effects in the analysis yields lower effective temperatures than pure hydrogen and helium model analyses. Optical spectra analyses of massive stars with spherical, hydrodynamic, wind-blanketed, non-LTE synthetic spectra also revised the temperature scale downwards \\citep[]{MSH02,HPN02,RPH03}. Our derived \\Teff~ values are still lower. Temperatures of O3~V and O4~V stars, according to the scale of \\citet[]{MSH02} (based on calculations with the CMFGEN code) are $\\simeq$47,500~K and $\\simeq$44,500~K, whereas we have obtained 41,000~K and 40,000-39,000~K. From the work of \\citet[]{HPN02}, who fit Balmer, \\ion{He}{1} and \\ion{He}{2} optical lines with FASTWIND \\citep[]{FASTWIND}, we can interpolate a temperature of 41,250~K for an O4~If$^+$, i.e. $\\sim$4,000~K higher than we derived for HD~190429A. \\citet[]{RPH03}, proceeding similarly, obtained 46,000~K for HD~93250 (in contrast with our result, 40,000~K) and 41,000~K for HDE~303308 (in agreement within the error bars with our value, 40,000~K). The remaining disagreement between these results and ours may originate in the fact that the analyses are based on optical and UV data respectively. We are planning to analyse optical spectra of the sample stars to check the consistency of our results from the UV and Far-UV.  The spectral lines in the FUSE and IUE ranges constrain \\Mdot~ as well. The mass-loss rate we derived for HD~190429A (1.2$\\cdot 10^{-5}$\\Myr) agrees with the values derived from H$\\alpha$ by \\citet[]{C77} ($1.1\\cdot10^{-5}$\\Myr), \\citet[]{L88} (1.26$\\cdot 10^{-5}$\\Myr), \\citet[]{Sal92} (8.6$\\cdot 10^{-6}$\\Myr) and \\citet[]{MPRM03} ($1.42\\cdot10^{-5}$\\Myr). Studies of IR and radio spectra yield lower \\Mdot\\space: $4.6\\cdot10^{-6}$\\Myr~ \\citep[radio]{Sal98}, $<$0.35$\\cdot 10^{-6}$\\Myr \\citep[IR]{PFG83}. For the dwarf stars the mass-loss rates derived in the literature are at least one order of magnitude higher than ours. For HD~93250, the exceptionally high value of \\Mdot=$4.1\\cdot10^{-5}$ \\citep[]{Lal95} was determined from radio fluxes at 8.54 GHz, but the data may be contaminated by non-thermal emission; for this object, there is also one measurement lower than ours by one order of magnitude. In view of the systematic differences for dwarf stars, we used the ``recipe'' of \\citet[]{Vink}, based on the radiation pressure driven wind theory, to predict the mass-loss rates of the sample stars from their derived photospheric parameters. The predictions are presented in Table \\ref{t_models}. The mass-loss rate derived for HD~190429A agrees with the prediction. For dwarf stars our values are lower than predicted by a factor of $\\sim$2 ($\\sim$4 for HD~64568).  In general, the stellar masses we have obtained are lower than values determined in previous works, as shown in Table \\ref{t_prev}. In the case of HD~93205, the mass had been derived from the orbital parameters of the binary system: \\Mstar$\\ge$ 31.5 \\citep[]{Mal01}, \\Mstar$\\ge$ 32.6 \\citep[]{SL93}, \\Mstar$\\ge$ 37 \\citep[]{Cal99}. For this star we derive 30$\\pm10$\\Msun, lower than previous determinations but compatible within the errors. The primary mass yields a mass of the secondary of 13$\\pm5$\\Msun~ for a mass ratio of q=0.423 \\citep[]{Mal01}, a mass significantly lower than derived for other O8~V stars in eclipsing systems like, for example, DN~Cas, for which \\citet[]{dncas} give 24\\Msun. However, it must be emphasized here that the uncertainty in \\logg and \\Rstar~ (see Section \\ref{s_fit_v}) makes the derived value of the stellar masses much more uncertain than other parameters (like \\Teff~ and \\Mdot). Hopefully the mass determination will be refined with a consistent analysis of optical spectra, planned as a future work.  Our results improve the empirical calibration of the Wind momentum-Luminosity Relation (WLR). The WLR relates the so-called modified wind momentum $D_{mom}=$ \\Mdot \\Vinf $(R_{\\ast}/R_{\\odot})^{0.5}$ to the luminosity of the star as $D_{mom} \\propto L^x$ \\citep[]{Kal95}. In the past few years a big effort has been made to calibrate the WLR \\citep[]{puls96,KP00,HPN02,RPH03,MPRM03} for its promising application to distance determinations. % In Figure \\ref{windmom} we plot the WLR for the total sample of stars studied in this paper and in paper~I. Our points are mostly consistent with the theoretical relation of \\citet[]{Vink} within the error bars.  We are planning to extend our modeling of Far-UV and UV spectra to a larger sample of O type stars and to the optical range. The lower \\Teff~ values (compared with previous results/calibrations) have great implications for understanding the ionisation of $HII$ regions, since cooler stars would produce less photons energetic enough to ionize hydrogen ($\\lambda<$ 912\\AA). For instance, the ionizing flux from a star of \\Teff=40,000~K amounts to only 30\\% of the flux from a \\Teff=50,000~K star with the same radius."
    },
    "0402/astro-ph0402177_arXiv.txt": {
        "abstract": "Most of the successful physical theories rely on the constancy of few fundamental quantities (such as the speed of light, $c$, the fine-structure constant, $\\alpha$,  the proton to electron mass ratio, $\\mu$, etc), and constraining the possible time variations of these fundamental quantities is an important step toward a complete physical theory. Time variation of $\\alpha$ can be accurately probed using absorption lines seen in the spectra of distant quasars. Here, we present the results of a detailed many-multiplet analysis performed on a new sample of Mg~{\\sc ii} systems observed in high quality quasar spectra obtained using the Very Large Telescope. The weighted mean value of the variation in ${\\bf \\alpha}$ derived from our analysis over the redshift range ${\\bf 0.4\\le z\\le 2.3}$ is ${\\bf \\Delta\\alpha/\\alpha}$~=~${\\bf (-0.06\\pm0.06)\\times10^{-5}}$. The median redshift of our sample (z$\\simeq$1.55) corresponds to a look-back time of 9.7 Gyr in the most favored cosmological model today. This gives a 3$\\sigma$ limit, ${\\bf -2.5\\times 10^{-16} ~{\\rm yr}^{-1}\\le(\\Delta\\alpha/\\alpha\\Delta t) \\le+1.2\\times 10^{-16}~{\\rm yr}^{-1}}$, for the time variation of $\\alpha$, that forms the strongest constraint obtained based on high redshift quasar absorption line systems. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0402/astro-ph0402388_arXiv.txt": {
        "abstract": "{ We combine samples of spiral galaxies and starburst systems observed with ISOCAM on board ISO to investigate the reliability of mid-infrared dust emission as a quantitative tracer of star formation activity. The total sample covers very diverse galactic environments and probes a much wider dynamic range in star formation rate density than previous similar studies. We find that both the monochromatic 15\\,$\\mu$m continuum and the $5 - 8.5\\,\\mu$m emission constitute excellent indicators of the star formation rate as quantified by the Lyman continuum luminosity $L_\\mathrm{Lyc}$, within specified validity limits which are different for the two tracers. Normalized to projected surface area, the 15\\,$\\mu$m continuum luminosity $\\Sigma_\\mathrm{15\\,\\mu m,ct}$ is directly proportional to $\\Sigma_\\mathrm{Lyc}$ over several orders of magnitude. Two regimes are distinguished from the relative offsets in the observed relationship: the proportionality factor increases by a factor of $\\approx 5$ between quiescent disks in spiral galaxies, and moderate to extreme star-forming environments in circumnuclear regions of spirals and in starburst systems.  The transition occurs near $\\Sigma_\\mathrm{Lyc} \\sim 10^{2}~\\mathrm{L_{\\odot}\\,pc^{-2}}$ and is interpreted as due to very small dust grains starting to dominate the emission at 15\\,$\\mu$m over aromatic species above this threshold. The $5 - 8.5\\,\\mu$m luminosity per unit projected area is also directly proportional to the Lyman continuum luminosity, with a single conversion factor from the most quiescent objects included in the sample up to $\\Sigma_\\mathrm{Lyc} \\sim 10^{4}~\\mathrm{L_{\\odot}\\,pc^{-2}}$, where the relationship then flattens. The turnover is attributed to depletion of aromatic band carriers in the harsher conditions prevailing in extreme starburst environments. The observed relationships provide empirical calibrations useful for estimating star formation rates from mid-infrared observations, much less affected by extinction than optical and near-infrared tracers in deeply embedded \\ion{H}{ii} regions and obscured starbursts, as well as for theoretical predictions from evolutionary synthesis models. ",
        "introduction": "\\label{Sect-intro}  Star formation is a fundamental process of galaxy formation and evolution. Estimates of the star formation rate (SFR) in galaxies at all redshifts are key indicators of the efficiency and mechanical feedback effects of star formation activity, of the chemical evolution of the interstellar and intergalactic medium, and, ultimately, of the cosmic star formation history.  Commonly used probes of the SFR include photospheric emission from hot stars in the ultraviolet, nebular H and He recombination lines as well as fine-structure lines arising in \\ion{H}{ii} regions from optical to radio wavelengths, and the total infrared luminosity ($\\lambda = 8 - 1000\\,\\mu$m), the bulk of which is due to heated dust reprocessing the interstellar radiation field \\citep[see, e.g., the review by][]{Ken98}.  However, ultraviolet, optical, and even near-infrared diagnostics are subject to potentially large uncertainties because of extinction in deeply embedded young star-forming sites and in nuclear regions of galaxies.  Nebular lines may be difficult to measure when intrinsically weak or superposed over a strong continuum. While dust emission suffers very little from extinction effects and is usually strong in star-forming environments, the total infrared luminosity is difficult to evaluate because it is generally derived from observations in a few wavelength intervals which do not constrain the spectral energy distribution accurately.  Moreover, a cirrus-like component unrelated to star-forming regions can contribute substantially to the far-infrared output of galaxies \\citep{Hel86, Sau92}.  Mid-infrared emission (MIR, $\\lambda = 5 - 20\\,\\mu$m) provides an alternative probe of star formation activity.  The ``classical'' spectrum of star-forming sources exhibits broad emission features often referred to as ``unidentified infrared bands'' (UIBs) and of which the most prominent dominate the $6 - 13\\,\\mu$m range, and a continuum rising towards long wavelengths at $\\lambda \\ga 11\\,\\mu$m \\citep[see the reviews by][]{Geb97, Tok97, Ces99, Gen00}. Various carbonaceous materials have been proposed to carry the UIBs, including the popular polycyclic aromatic hydrocarbons (PAHs; e.g. \\citealt{Leg84}) that we adopt hereafter. \\citet{Peeters02} have analysed their shape and relative amplitude variations in different classes of Galactic objects. The continuum emission is generally attributed to very small dust grains (VSGs; e.g. \\citealt{Des90}) about which little is known.  Superposed on these PAH and VSG components, H recombination lines and fine-structure lines of various metals originating in \\ion{H}{ii} and photodissociation regions are observed as well \\citep[e.g.][]{Stu00}.  These lines may however not always be measurable because of their weakness or of insufficient spectral resolution.  Numerous past studies have established that PAH and $\\lambda \\ga 11\\,\\mu$m continuum emission trace well star-forming regions but their usefulness as {\\em quantitative\\/} diagnostics is still debated.  Complications arise from the different nature of the emitting particles and by their being out of thermal equilibrium under most radiation field conditions, undergoing large temperature fluctuations of several 100~K \\citep[e.g.][]{Greenberg74, Draine85, Puget89}. In addition, although both species are predominantly heated by energetic radiation, PAHs can also be excited by softer optical and near-ultraviolet photons as indicated by their detection in the diffuse interstellar medium, in regions of insufficient far-ultraviolet energy density to account for their heating \\citep[e.g.][]{Sel90, Mat96, Uch98, Uch00, Li02}. Furthermore, empirical evidence indicates that the $\\lambda \\ga 11\\,\\mathrm{\\mu m}$ emission is produced by a mixture of dust particles akin to PAHs (or at least whose flux variations follow well those of PAHs) and of VSGs \\citep[e.g.][]{Hony01, Rou01b}. The first component is best seen in quiescent environments such as disks of spiral galaxies while the second becomes prominent in active star formation sites.  It is not yet clear how their combined emission varies over a large dynamic range in star formation intensity.  On the other hand, spatially resolved studies of Galactic and Magellanic Clouds \\ion{H}{ii} regions have revealed that both the PAH features and the VSG continuum are produced in the vicinity of massive stars, the former arising in photodissociation regions (PDRs) at the interface between ionized and molecular gas and the latter peaking closer to the ionizing stars \\citep[e.g.][]{Geb97, Tok97, Ver96, Cre99, Con00}. MIR imaging of external galaxies has shown that bright emission from both components is closely associated with active star-forming sites on large scales as well \\citep[e.g.][]{Mir98, Mat99, Rou01c, FS03}. A strong coupling with the SFR may thus exist and has been demonstrated for disks of spiral galaxies by \\citet{Rou01c}.  Specifically, these authors found that the broadband $5 - 8.5\\,\\mu$m and $12 - 18\\,\\mu$m fluxes vary linearly with the H$\\alpha$ line flux in the disks of 44 spirals. \\citet{FS03} also found a direct proportionality between the monochromatic 15\\,$\\mu$m continuum ($\\Delta\\lambda = 0.4\\,\\mu$m) and the [\\ion{Ar}{ii}] 6.99\\,$\\mu$m line emission in the nearby starbursts M\\,82, NGC\\,253, and NGC\\,1808, down to spatial scales of $\\sim 100$~pc.  In this paper, we pursue the work of \\citet{Rou01c} and \\citet{FS03} by combining samples of spiral and starburst galaxies observed with the ISOCAM instrument \\citep{Ces96} on board the Infrared Satellite Observatory \\citep[ISO;][]{Kes96}. The merged sample covers diverse environments ranging from quiescent galactic disks to infrared-luminous merging systems.  This allows us to extend the investigation to higher activity levels and to test whether previous results restricted to specific environments can be generalized into more universal relationships.  We derive the dependence of the PAH-dominated $5 - 8.5\\,\\mu$m emission and the VSG-probing monochromatic $15\\,\\mu$m continuum on the production rate of Lyman continuum photons $Q_\\mathrm{Lyc}$ quantifying the SFR. The resulting empirical calibrations provide useful tools in MIR studies of star-forming galaxies as well as constraints for models predicting the dust emission of such systems.  The paper is organized as follows.  Section~\\ref{Sect-sample} presents the galaxy sample.  Section~\\ref{Sect-diagn} describes the MIR indicators and the SFR estimates obtained from more classical diagnostics.  Section~\\ref{Sect-res} discusses the derived calibrations and Sect.~\\ref{Sect-conclu} summarizes the results. ",
        "conclusions": ""
    },
    "0402/astro-ph0402427_arXiv.txt": {
        "abstract": " ",
        "introduction": "Recent N--body+SPH cosmological simulations of the formation of disc galaxies reproduce the observed Tully-Fisher (TF) relation (Dale et~al.\\ 1999), provided the mass--to--light (M/L) ratio of the stellar component is rather low, {\\mbox{M/L$_I$=0.7--1}} in the I--band (Sommer-Larsen \\& Dolgov 2001; Sommer-Larsen et~al.\\ 2003; Fig.~\\ref{fig:tully}). The location of the simulated galaxies in the ($M_*$, $V_c$) plane of Fig.~\\ref{fig:tully} is quite independent of the adopted Initial Mass Function (IMF) or feedback efficiency: the baryonic mass that cools out to form a galactic disc and its resulting circular velocity correlate so that data points tend to move along the TF relation, hardly affecting the zero--point (Navarro \\& Steinmetz 2000ab). However, the IMF is crucial for the M/L ratio, to translate the stellar masses $M_*$ to luminosities and compare the simulated TF relation to the empirical one. Although the zero--point of the simulated TF may change with the concentration of the dark matter halos, and hence with the normalization of the power spectrum $\\sigma_8$ (Navarro \\& Steinmetz 2000ab; Eke et~al. 2001), many other arguments support a low stellar M/L ratio in spiral galaxies:  The stellar mass of the Milky Way is {\\mbox{M$_* \\sim 5 \\times 10^{10}~M_{\\odot}$}}; to lie on the observed TF relation as other spirals, its M/L$_I$ must be $\\lsim$1 (Sommer-Larsen \\& Dolgov 2001; Fig.~\\ref{fig:tully}). A low {\\mbox{M/L$_I < 0.8$}} is also derived for the massive Sb galaxy NGC 2841, when compared to the observed TF relation (Portinari et~al.\\ 2004a, hereinafter PST; Fig.~\\ref{fig:tully}).  Based on bar instability arguments, Efstathiou et~al.\\ (1982) suggest an upper limit of {\\mbox{M/L$_B \\leq 1.5~h$}} for discs, i.e.\\ M/L$_B \\lsim 1$ for $h$=0.7. ($h$ indicates the Hubble constant H$_0$ in units of 100~km~sec$^{-1}$~Mpc$^{-1}$).  The stellar M/L ratio is related to the issue as to whether discs are maximal or sub--maximal, i.e.\\ as to whether they dominate or not the dynamics and rotation curves in the inner galactic regions. Even in the case of maximal stellar discs, lower M/L ratios for the stellar component are required, than those predicted by the Salpeter IMF (Bell \\& de Jong 2001). And it is still much debated whether discs are maximal or sub--maximal; for his favoured sub--maximal disc model, Bottema (2002) finds M/L$_I \\sim 0.82$.  Finally, two recent dynamical studies of individual spiral galaxies yield {\\mbox{M/L$\\sim$1}} in the B, V and I band for the Sc galaxy NGC 4414 (Vallejo et~al.\\ 2002) and M/L$_I$=1.1 for the disc of the Sab spiral 2237+0305, Huchra's lens (Trott \\& Webster 2002).  In this paper we discuss if M/L ratios so low are compatible with our understanding of stellar populations and chemical evolution in disc galaxies. We also address the effects of different star formation histories on the TF relation for different Hubble types. ",
        "conclusions": ""
    },
    "0402/astro-ph0402611_arXiv.txt": {
        "abstract": "Images obtained with the CFHTIR camera on the Canada-France-Hawaii Telescope are used to investigate the near-infrared photometric properties of the star-forming M81 group galaxy NGC 3077. The spectral-energy distribution (SED) of the near-infrared light within 10 arcsec of the nucleus (1) is very different from that of `typical' dwarf ellipticals (dEs), blue compact dwarf galaxies, and HII/starburst galaxies, and (2) is consistent with the $2\\mu$m light being dominated by hot young (log(t$_{yr}) < 6.8$) stars reddened by A$_V = 3 - 4$, with A$_V \\geq 8$ mag in some regions, including previously detected areas of CO emission. A population like that near the center of NGC 205 likely contributes only a modest fraction of the light near $2\\mu$m.  A number of candidate star clusters are detected in and around NGC 3077. These objects have near-infrared brightnesses and colors that are consistent with them being classical globular clusters and young star clusters. The specific frequency of globular clusters in NGC 3077 is estimated to be S$_N = 2.5$, which falls within the range of S$_N$ measured in nearby dEs. The candidate young clusters have photometric masses that are similar to those of compact young clusters in other active star-forming systems, and SEDs consistent with ages log(t$_{yr}) \\leq 6.6$. Based on the masses and ages of the young clusters, it is estimated that the star formation rate in NGC 3077 was at least $0.25 - 0.50$M$_{\\odot}$ year$^{-1}$ during the past few million years. ",
        "introduction": "The origin of dwarf spheroidal (dSph) and dwarf elliptical (dE) galaxies, and their relation to dwarf irregular (dIrr) galaxies, has long been a matter of debate (e.g. Tajiri \\& Kamaya 2002; Mayer et al. 2001; Silk, Wyse, \\& Shields 1987; Thuan 1985; Lin \\& Faber 1983). At issue is whether the morphological characteristics of these galaxies are primarily the result of (1) local conditions within the mini-halos from which they first collapsed, or (2) external factors, such as proximity to a much larger companion. Numerical simulations predict that the initial density and dark matter content are critical parameters for defining final morphology (e.g. Dekel \\& Silk 1986; Ferrara \\& Tolstoy 2000; Carraro et al. 2001), and the correlation between the chemical compositions, central surface brightness, and mass-to-light ratios of a wide range of dwarf galaxies is consistent with initial conditions playing a key role in dwarf galaxy evolution (Prada \\& Burkert 2002). However, the occurence of morphological segregation in nearby groups (e.g. Karachentsev et al. 2002) and the relation between gas content and distance from nearby large companions in the Local Group (Blitz \\& Robishaw 2000) indicate that environment also plays a key role in dwarf galaxy evolution. It is also clear that tidal interactions can profoundly affect the structural properties of dwarf galaxies in hierarchical systems, with the Sagitarrius dwarf being a prime example (Ibata, Gilmore, \\& Irwin 1995).  Studies of nearby galaxies and their companions will provide insights into dwarf galaxy evolution. The Milky-Way, M31, and M81 have similar masses (e.g. Kochanek 1996; Schroder et al. 2002; Perrett et al. 2002), morphologies, and environments, and yet have very different satellite systems. The companions of the Milky-Way include the dIrr Magellanic Clouds and a number of dSphs, while the brightest companions of M31 are the dE galaxies NGCs 147, 185, and 205, and the compact elliptical galaxy M32; M31 also has a number of fainter dSph satellites. The companions of M81 show considerable diversity, consisting of dSphs, dEs, and dIrrs, and some of these show morphological peculiarities that are related to tidal interactions. Indeed, the M81 group appears to be undergoing significant evolution at the present day, making this the nearest laboratory for studying the effects of on-going tidal interactions on gas-rich dwarf galaxies.  NGC 3077 is an actively star-forming member of the M81 group that is not easily placed within the Hubble sequence, although Price \\& Gullixson (1989) conclude that there are similarities with the Local Group dE NGC 185 at red wavelengths. There is a  prominent central dust lane, and filamentary H$\\alpha$ emission (Barbieri, Bertola, \\& di Tullio 1974). The optical depth of the central dust lane is $\\sim 0.5$, and this absorption obscures the isophotal center of the galaxy at visible wavelengths (Price \\& Gullixson 1989). The central regions of NGC 3077 are dominated by hot stars (Benacchio \\& Galletta 1981), and the star formation rate in the central 700 parsecs exceeds that in normal discs (e.g. Thronson, Wilton, \\& Ksir 1991; Ott, Martin, \\& Walter 2003). Martin (1998) identified a number of expanding gas shells in the ISM of NGC 3077. These structures have kinematic ages $\\leq 10$ Myr, and spatial scales and expansion velocities that are indicative of energy input from a large number of SNe, thus confirming that there has been considerable recent star-forming activity.  NGC 3077 is not evolving in isolation. Ott et al. (2003) conclude that while the hot gas in NGC 3077 is confined at present, some of it may eventually escape into the M81 group intergalactic medium. An HI bridge connects M81 and NGC 3077 (van der Hulst 1979), while tidal spurs also link NGC 3077 to other members of the M81 group (e.g. Boyce et al. 2001). The interactions with M81 and its companions likely spurred the current star-forming episode in NGC 3077, and may also have triggered the formation of young compact star clusters in the M81 disk (Chandar et al. 2001).  Studies of stars and star clusters in NGC 3077 will provide insight into the past history of this galaxy. The central regions of NGC 3077 have yet to be resolved into stars. However, Sakai \\& Madore (2001) and Karachentsev et al. (2002) resolved stars on the upper RGB in the outer regions of NGC 3077, and the latter conclude that $\\mu_0 = 27.91$, placing the system behind M81. The presence of RGB stars indicates that NGC 3077 is not a recently formed tidal fragment, as has been suggested for some other M81 companions (e.g. Karachentsev, Karachentseva, \\& Boerngen 1985; Yun, Ho, \\& Lo 1994; Boyce et al. 2001).  In the present study, deep $J, H,$ and $K'$ images obtained with the CFHTIR camera are used to probe the near-infrared spectral energy distribution (SED) and morphology of the central regions of NGC 3077, and search for star clusters. Observations of this galaxy in the infrared are of interest because the evolved cool stars that formed during intermediate and early epochs might be expected to dominate the light output at these wavelengths; knowledge of the spatial distribution of such stars may yield insight into the nature of NGC 3077 prior to the most recent tidal interactions, and provide clues about its appearance after the current star-forming episode fades. Light at infrared wavelengths is also less affected by dust absorption than at visible wavelengths, making it easier to probe the heavily obscured central regions of NGC 3077, and detect bright star clusters.  The paper is structured as follows. The observations and the data reduction procedures are described in \\S 2. The infrared SED of integrated light in the central regions of NGC 3077 and the isophotal properties of the galaxy are investigated in \\S 3. Comparisons are also made with the Local Group dE galaxy NGC 205, which shows some similarities with NGC 3077, and may be in a more advanced evolutionary state (Davidge 1992). While the present data do not have the angular resolution needed to detect individual stars, a number of potential star clusters are identified, and the nature of these is investigated in \\S 4. A summary and discussion of the results follows in \\S 5. ",
        "conclusions": "\\subsection{The Near-Infrared SED of NGC 3077}  Sub-arcsec angular resolution $J, H,$ and $K'$ images obtained with the CFHTIR camera have been used to investigate the near-infrared photometric properties of the central regions of the M81 group galaxy NGC 3077. In \\S 3 it was demonstrated that (1) the integrated near-infrared SED of the central regions of NGC 3077 differs from that of `typical' dE's, BCDG's and HII/starburst galaxies, and (2) the light near $2\\mu$m is dominated by very young stars. The latter result is perhaps not surprising, as previous studies have found that the SFR in NGC 3077 is high, although there is significant scatter among the estimates. Thronson et al. (1991) estimate that the SFR is 0.06 M$_\\odot$ year$^{-1}$ based on the integrated H$\\alpha$ flux, and $\\leq 0.25$ M$_\\odot$ year$^{-1}$ from FIR emission. Walter et al. (2002) compute a SFR of 0.05 M$_\\odot$ year$^{-1}$ based on extinction-corrected H$\\alpha$ measurements. Meier et al. (2001) estimate that the SFR is 0.4 M$_\\odot$ year$^{-1}$ from the 2.6 mm continuum flux, while Ott et al. (2003) estimate that the SFR is 0.6 M$_\\odot$ year$^{-1}$ based on the energy needed to create super gas shells. As noted by Ott et al. (2003), the high SFR inferred in their study may suggest that the SFR has dropped recently. While the present-day SFR in NGC 3077 is markedly lower than in M82 (e.g. Ott et al. 2003), the efficiency with which stars form out of molecular gas in both galaxies is similar (Walter et al. 2002).  The detection of a population of suspected young star clusters, which have ages log(t$_{yr}$) between 6.0 and 6.8 and masses log(M$_{\\odot}$) between 4.0 and 5.0 (\\S 4), is consistent with a very high recent SFR. The total integrated mass in the young clusters is $1 - 2 \\times 10^6$ M$_{\\odot}$. Given that the near-infrared SEDs of these clusters are suggestive of an age log(t$_{yr}) \\leq 6.6$, then the SFR needed to produce these objects is $0.25 - 0.50$ M$_{\\odot}$ year$^{-1}$, which falls within the range of estimates computed using other techniques. This is a lower limit to the total SFR, as it does not include stars that formed in clusters that have been disrupted, or clusters that are below the faint limit of these data.  The near-infrared SED of NGC 3077 does not preclude a modest contribution from a system having an SED like that of NGC 205. The presence of such a population is consistent with the detection of RGB stars in NGC 3077 by Sakai \\& Madore (2001) and Karachentsev et al. (2002), and the discovery of a healthy number of candidate old globular clusters (\\S 4). Thus, NGC 3077 is not a recently formed system, as may be the case for some members of the M81 group (Yun et al. 1994, Boyce et al. 2001).  The models used here to simulate the near-infrared SED of NGC 3077 include contributions from stars and gas emission. However, Hunt et al. (2002) investigated the infrared photometric properties of star-forming galaxies, and concluded that thermal emission from hot dust may contribute significantly to the light from very active star-forming systems at wavelengths longward of $2\\mu$m, and there are hints that emission from hot dust may contribute significantly to the infrared light from NGC 3077. One clue comes from the near-infrared SED. Thermal emission from dust with temperatures cooler than a few hundred K does not contribute significantly to light at wavelengths shortward of $2\\mu$m, and so the $J-H$ color of a system with significant dust emission will be the same as from a galaxy lacking this emission, while the $H-K$ color may be much redder. It is evident from Figure 3 that the majority of pixels near the center of NGC 3077 have $J-H$ colors that are consistent with those of dEs and BCDGs, but have $H-K$ colors that are redder than in these systems, as expected if emission from hot dust is present. Emission from hot dust will likely be concentrated in the dense central star-forming regions of NGC 3077, where the radiation field is most intense. In \\S 3.3 it was shown that the $J-H$ color measured near the center of NGC 3077 is similar to that measured at larger radii from 2MASS survey data, while the central $H-K$ color is much redder than at large radii, and so the radial color profile is also consistent with emission from hot dust. Finally, when compared with other galaxies, NGC 3077 has a very low mass of cool gas (Stickel et al. 2000), as might be expected if a large fraction of the dust is heated by young stars. Observations of NGC 3077 in the $3 - 5\\mu$m region will allow firmer constraints to be placed on any contribution made by thermal emission from hot dust. It is evident from Figure 4 that if significant emission from hot dust is present in NGC 3077 then the net result will be to allow for a larger contribution from a NGC 205-like SED.  \\subsection{The Evolution of NGC 3077}  Tidal interactions are common events in the local Universe, as there are debris trails in the halos of the Milky-Way (e.g. Ibata et al. 1995), M31 (e.g. Ibata et al. 2001), and in nearby groups (e.g. Boyce et al. 2001). Recognizing that tidal interactions can affect the structural characteristics of galaxies, and can also explain the morphology-density relation between dIrr and dSph systems, Mayer et al. (2001) suggested that the dSph and dE companions of the Milky-Way and M31 may have originally been gas-rich disky galaxies that were transformed by tidal interactions into spheroidal systems. Mayer et al. (2001) suggest that tides trigger bar instabilities that channel gas into the central regions of the progenitor, where star formation occurs. Feedback then heats the ISM, while the bar re-distributes angular momentum to the outer regions of the galaxy, which are subsequently stripped away. The bar eventually buckles, and the result is a system that has lost much of its angular momentum and has been transformed from a gas-rich disky to a gas-poor spheroidal morphology. Whether the final product is a dSph or a dE depends on the surface brightness of the progenitor, in the sense that dSphs evolve from low surface brightness dIrrs that experience multiple bursts of star formation, while dEs evolve from higher surface brightness dIrrs that experience one dominant, extended, star-forming episode that occurs over a $1 - 2$ Gyr period. Simulations indicate that the time scale for the transformation is a few Gyr, so there is a reasonable expectation of viewing this process at work in nearby galaxy groups.  While tidal interactions undoubtedly influence galaxy evolution, they do not provide a panacea for explaining dwarf galaxy morphology. Indeed, there are isolated dSph galaxies, such as Tucana, that likely have not been subjected to tidal interactions but still show remarkable similarities to dSphs in denser environments. The bar-driven transformation process described by Mayer et al. (2001) should have a major impact on radial population gradients, and yet the radial population behaviour of Tucana is similar to dSphs in denser environments (Harbeck et al. 2001). There are also galaxies in hierarchical systems that appear not to have been altered by tidal forces. In particular, the surface brightness profile and central black hole mass of M32, which is a galaxy that many have argued may be an extreme endpoint of tidal pruning (e.g. Faber 1973; Nieto 1990), suggest that the structural characteristics of this galaxy were imprinted early on and have not since been greatly altered (Graham 2002).  Tidal interactions are clearly affecting the properties of NGC 3077. Much of the atomic gas associated with NGC 3077 is in a tidal arm that is well offset from the main body of the galaxy (Walter et al. 2002; Yun et al. 1994), while there is a string of molecular complexes extending to the north and west of the center of NGC 3077 (Meier et al. 2001), the positioning of which may also be the result of tidal effects. While the interstellar medium of NGC 3077 is clearly being disrupted, when integrated over a large area, the HI mass to light ratio of NGC 3077 is more appropriate for a dIrr, rather than a dE (Walter et al. 2002), as might be expected if NGC 3077 is undergoing a morphological transformation. Of course, it is possible that some of the gas currently associated with NGC 3077 may have been stripped from another galaxy.  Simulations discussed by Mayer et al. (2001) suggest that the time scale for bar evolution in tidally influenced gas-rich dwarf galaxies is on the order of a few Gyr, and so there is a reasonable expectation of observing barred tidally interacting dwarf galaxies. In fact, stars in the bar of the LMC have an age of $\\sim 5$ Gyr (Smecker-Hane et al. 2002), supporting the notion that the bars in dIrr galaxies in hierarchical systems can be stable against buckling for long periods of time. The timescale for bar disruption is much longer than the time since the last major encounter between M81, M82, and NGC 3077, so if a bar formed in NGC 3077 after the last encounter then it should still be present. This being said, the near-infrared images of NGC 3077 do not show evidence of a bar, and the peaky light profile of NGC 3077 (\\S 3.3) is not consistent with a bar-dominated light distribution.  dIrr and dE galaxies have very different S$_N$'s (e.g. Harris 1991), and so the old globular cluster content of NGC 3077 may provide clues about the nature of the galaxy prior to its most recent encounter with M81. In particular, if NGC 3077 were a `typical' dE before encountering M81 then it should contain a rich population of classical globular clusters, while if it were recently a dIrr galaxy then it should contain only a modest population of such objects. A caveat is that the young clusters that form during vigorous star-forming episodes can have globular cluster-like masses, as appears to be the case in NGC 3077 (\\S 4.3). If these clusters are not disrupted then, when viewed in a few Gyr, the galaxy will contain an even larger globular cluster population, albeit spanning a range of ages. The S$_N$ may thus change with time over the course of the transformation process.  Spectroscopic observations, which will yield line strengths and radial velocities, will be essential to distinguish between actual clusters and faint field stars in NGC 3077. The globular clusters in NGC 3077 will likely be more metal-poor than the surrounding field population (e.g. da Costa \\& Mould 1988), and so will likely have weak absorption lines. Higher angular resolution images will also be useful for identifying clusters, as both globular clusters (Kundu \\& Whitmore 2001) and compact young clusters (Chandar et al. 2001; Harris et al. 2001) have characteristic sizes of a few parsecs, and so will have a non-stellar appearance when viewed with image qualities approaching 0.1 arcsec FWHM.  Although lacking spectra and high-resolution images, the present data still provide tantalizing hints into the nature of NGC 3077. In \\S 4.2 it was estimated that S$_N = 2.5$, suggesting that the specific frequency of globular clusters in NGC 3077 is similar to that of NGC 205. The specific globular cluster frequency is consistent with the structural characteristics of NGC 3077, which suggest that if star formation was terminated immediately then the galaxy would fade to become a dE,n. While the young central cluster in NGC 3077 is offset slightly from the isophotal center of the galaxy, this occurs in roughly 20\\% of dE,n (e.g. Binggeli, Barazza, \\& Jerjen 2000)."
    },
    "0402/astro-ph0402082_arXiv.txt": {
        "abstract": "{We use mid-infrared spectral decomposition to separate the 6$\\mu$m mid-infrared AGN continuum from the host emission in the ISO low resolution spectra of 71 active galaxies and compare the results to observed and intrinsic 2-10keV hard X-ray fluxes from the literature. We find a correlation between mid-infrared luminosity and absorption corrected hard X-ray luminosity, but the scatter is about an order of magnitude, significantly larger than previously found with smaller statistics. Main contributors to this scatter are likely variations in the geometry of absorbing dust, and AGN variability in combination with non-simultaneous observations. There is no significant difference between Type 1 and Type 2 objects in the average ratio of mid-infrared and hard X-ray emission, a result which is not consistent with the most simple version of a unified scheme in which an optically and geometrically thick torus dominates the mid-infrared AGN continuum. Most probably, significant  non-torus contributions to the AGN mid-IR continuum are masking the expected difference between the two types of AGN. ",
        "introduction": "The mid-infrared and hard X-rays are two of the regions of the electromagnetic spectrum that are of particular interest for the study of active galactic nuclei (AGN). Hard X-rays, unless extremely obscured in fully Compton-thick objects, can provide a direct view to the central engine, and are often believed to be a reasonable isotropic measure of the bolometric luminosity of the AGN. The nuclear infrared continuum in Seyferts, in contrast, is due to AGN emission reprocessed by dust, either in the putative torus or on somewhat larger scales, e.g. inside the Narrow Line Region. The observed mid-infrared emission is thus a function not only of the AGN luminosity but also of the distribution of the obscuring matter and of the viewing direction of the observer. In the most simple form this is due to the covering factor and distance of the obscuring dust, but much more complex radiative transfer effects may occur in high optical depth configurations (e.g. Pier et al. \\cite{pier92}). In general terms, measurements of the mid-infrared continuum in conjunction with hard X-ray observations can be thought of as testing unification scenarios for AGN. A tight relation between the two quantities has recently been reported by Krabbe et al. (\\cite{krabbe01}) on the basis of mid-infrared imaging of eight nearby Seyferts. Clavel et al. (\\cite{clavel00}) have found a large difference between the equivalent widths of the mid-infrared aromatic `PAH' emission features in Type 1 and 2 Seyferts. Under the assumption that the Type 1 and 2 subsamples are well matched in AGN luminosity and host properties, they interpret this as an orientation dependent depression of continuum in Seyfert 2s with respect to the host-related isotropic PAH emission.  The issue is plagued, however, by the technical difficulty of isolating the AGN mid-infrared continuum from the host galaxy emission. IRAS has readily detected large numbers of AGN, but the host contribution to these large beam infrared spectral energy distributions (e.g. Spinoglio et al. \\cite{spinoglio95}) is not easy to quantify and significant in all but the very powerful AGN. One way to address this difficulty is high spatial resolution imaging from groundbased telescopes, e.g. in the L-  and M-bands (e.g. Alonso-Herrero et al. \\cite{alonso01}) or the N-band (e.g. Maiolino et al. \\cite{maiolino95}; Krabbe et al. \\cite{krabbe01}). This method produces reliable results in cases of good surface brightness contrast between AGN and host, but can face ambiguities in cases where the AGN is surrounded by intense star formation on scales similar to the spatial resolution used, in particular if the observations are diffraction limited by a moderate size telescope (e.g. \\object{NGC\\,6240}, \\object{NGC\\,4945}; Krabbe et al. \\cite{krabbe01}).  We use the alternative approach of isolating the AGN mid-infrared continuum {\\em spectrally}, making use of the sizeable database of low resolution mid-infrared spectra of AGN that are a legacy of the Infrared Space Observatory ISO. Low resolution spectra of galaxies can be decomposed into three components (Laurent et al. \\cite{laurent00}): A component dominated by the aromatic `PAH' features arising in photodissociation regions or the diffuse interstellar medium of the host, an H\\,II region very small grain continuum rising steeply towards wavelengths beyond 10$\\mu$m, and for active galaxies a typically flatter thermal AGN dust continuum. Starlight is unimportant except for quiescent objects like ellipticals or nearby spirals with weak central star formation or AGN activity. The three components may also be obscured, with the additional complication of ice features (Spoon et al. \\cite{spoon02}). A full spectral decomposition accounting for all these effects can be attempted in cases of good signal-to-noise ratio and full wavelength coverage (e.g. Tran et al. \\cite{tran01}, Spoon et al. \\cite{spoon03}). Since most of our data are for the restricted ISOPHOT range (5.8 to 11.8 $\\mu$m) that limits the accuracy of separating silicate absorption and PAH emission, and since some spectra are of limited S/N, we follow a more straightforward approach.  In the range covered by the ISOPHOT spectra, the AGN emission is most easily isolated shortwards of the complex of aromatic emission features (Laurent et al. \\cite{laurent00}). We determine a continuum at 6$\\mu$m rest wavelength, and eliminate non-AGN emission that will in many cases be present in the fairly large beam. This is done by subtracting a star formation template scaled with the strength of the aromatic `PAH' features arising in the host or in circumnuclear star formation. This method does not require to spatially resolve the AGN from the contaminating star formation, and will face its limits only when trying to identify a weak AGN in the presence of strong star formation or a stellar continuum that can be detectable for the most nearby galaxies. Our sample of 71 AGN is then used to quantify the relation of mid-infrared and X-ray emission at significantly better statistics than previously possible. ",
        "conclusions": "We have used spectral decomposition of a large sample of ISO spectra of AGNs to isolate the AGN 6$\\mu$m continua. We compare these mid-infrared continua to intrinsic hard X-ray fluxes from the literature, assumed to be a fair isotropic measure of AGN luminosity. Due to this normalization, our comparison can test AGN properties and unification aspects without the level of sensitivity to selection biases (e.g. luminosities, ratio of Type 1 and Type 2 objects in the sample) that is found in comparisons of absolute quantitities or of normalizations to non-AGN quantitities. The main results are:  (1) Mid-infrared and intrinsic X-ray fluxes correlate, but with significant scatter. Ratios vary by more than an order of magnitude, the dispersion in the log is 0.4 for Seyfert 1 and 0.69 for Seyfert 2 galaxies. Main contributors to this spread likely include variations in the AGN spectral energy distribution and geometry of obscuring dust, as well as the effects of AGN variability.  (2) There is no significant difference between Type 2 and Type 1 in the average ratio of X-ray and mid-infrared continuum, in contrast to expectations from unified scenarios invoking an optically and geometrically thick torus emitting anistropically in the mid-IR. Most likely, this is due to a large contribution from extended dust components emitting more isotropically. This will dilute anisotropic emission that, however, may still be present at lower levels."
    },
    "0402/astro-ph0402561_arXiv.txt": {
        "abstract": "Accretion flows having low angular momentum and low viscosity can have standing shock waves. These shocks arise due to the presence of multiple sonic points in the flow. We study the region of the parameter space in which multiple sonic points occur in viscous flows in the absence of cooling. We also separate the parameter space in regions allowing steady shocks and oscillating shocks. We quantify the nature of two critical viscosities which separate the flow topologies. A post-shock region being hotter, it emits harder X-rays and oscillating shocks cause oscillating X-ray intensities giving rise to quasi-periodic oscillations. We show that with the increase in viscosity parameter, the shock always moves closer to the black hole. This implies an enhancement of the quasi-periodic oscillation frequency as viscosity is increased. ",
        "introduction": "In the standard theory of thin accretion flows around black holes (Shakura \\& Sunyaev, 1973, hereafter referred to as SS73) viscosity plays a major role. Viscosity transports angular momentum outwards and allows matter to sink into the potential well formed by the central compact object. In this model, the flow angular momentum is assumed to be Keplerian and this is the standard notion about how matter is  accreted. However, Chakrabarti \\& Molteni (1995, hereafter referred to as Paper I), and Lanzafame, Molteni \\& Chakrabarti (1998, hereafter referred to as Paper II), through extensive numerical simulations showed that the angular momentum distribution depends strictly on the viscosity parameter and the way the viscous stress is defined. They showed that close to a black hole, the disk does not have a Keplerian distribution. This is because the flow must be supersonic on the horizon (Chakrabarti, 1990a) whereas a Keplerian disk is always subsonic (SS73). In Papers I  and II, it was shown that for a large region of the parameter space, shocks may form in accretion flows and when viscosity is increased beyond a critical value (Chakrabarti, 1990ab; Chakrabarti 1996, hereafter C96a), the shocks disappear.  Paper I also improved the concept of viscosity parameter $\\alpha$ (SS73): it argued that in a generalized flow with significant radial velocity $\\vartheta$, the viscous stress $w_{\\phi r}$ should not be equated to $-\\alpha P$ as in SS73, where $P$ is the total pressure, but to $-\\alpha_\\Pi (P+\\rho \\vartheta^2)$, (actually, its vertically integrated value using a thin disk approximation) where, $\\rho$ is the density and a subscript $\\Pi$ is given to $\\alpha$ to distinguish it from the Shakura-Sunyaev viscosity parameter. The latter prescription naturally goes over to the original prescription when radial velocity is unimportant as in the case of a standard Keplerian disk model (SS73), however, when the radial velocity is important as in the transonic flow solutions (Chakrabarti 1990a), the latter definition preserves the angular momentum even across axisymmetric discontinuities, such as accretion shocks. The reason is, according to the Rankine-Hugoniot conditions (Landau \\& Lifshitz, 1959), in a steady flow, the sum of thermal pressure and ram pressure, i.e., $P+\\rho \\vartheta^2$ is continuous across discontinuities. This makes the viscous stress $w_{r\\phi}$ continuous across axisymmetric discontinuities as well.  In an earlier study, Chakrabarti (1989a, hereafter C89a) considered the transonic properties of isothermal accretion flows and showed that for a large region of the parameter space spanned by the specific angular momentum and the temperature of the flow, an accretion disk can have standing shock waves. The specific angular momentum of the disk was smaller than that of a Keplerian disk everywhere. This flows come about especially when the matter is accreted from the winds of a binary companion. Subsequently, Chakrabarti (1990b, hereafter C90b) showed that inclusion of viscosity reduces the region of the parameter space in that, at a sufficiently high viscosity, the Rankine-Hugoniot conditions which must be satisfied at a steady shock are not satisfied anywhere in the flow. Existence of standing shocks in sub-Keplerian inviscid accretion disks have been tested independently by several groups since then (Nobuta and Hanawa 1994; Yang and Kafatos, 1995; Lu and Yuan, 1997). Numerical simulations have also been carried out with several independent codes such as Smootherd Particle Hydrodynamics (SPH) and Total Variation Diminishing (TVD) and distinct standing shocks were found exactly at the predicted locations (Chakrabarti \\& Molteni, 1993; Molteni, Ryu \\& Chakrabarti, 1996).  In more recent years, it has become evident that the standing shocks may be very important in explaining the spectral properties of black hole candidates (Chakrabarti \\& Titarchuk, 1995, hereafter CT95) as the post-shock region behaves as the boundary layer where accreting matter dissipates its thermal energy and generates hard X-ray by inverse Comptonization. C96a considered unification of solutions of winds and accretion around compact objects. However, the cooling was treated in terms of a parameter and no parameter space was studied. The post-shock region is also found to be responsible to produce relativistic outflows (Chakrabarti, 1999; Chattopadhyay and Chakrabarti, 2002). Furthermore, numerical simulations indicated that the shocks may be oscillating at nearby regions of the parameter space in presence of cooling effects (Molteni, Sponholz \\& Chakrabarti, 1996) and the shock oscillations correctly explain intricate properties of quasi-periodic oscillations (Chakrabarti \\& Manickam, 2000). Recent observations do support the presence of sub-Keplerian flows in accretion disks (Smith et al. 2001; Smith, Heindl and Swank, 2002).  In view of the importance of the sub-Keplerian flows we plan to re-investigate the work done on isothermal flow by C89a and C90b by extending them to study {\\it polytropic flows} to check the properties of shock waves in viscous flows. What is more, unlike C89a and C90b, we investigate the behaviour of the solutions in the entire parameter space spanned by the specific energy, angular momentum and the viscosity. In C96a some work was done, the parameter space was not explored. We find very important results: even when the viscosity parameter is very high, the flow continues to have three sonic points: a prime condition to have a standing or oscillating shock waves. However, the parameter space for standing shock waves is gradually reduced with the increase of viscosity. On the other hand, we discover that the shock location itself is reduced with the increase in viscosity parameter.  We wish to emphasize that the problem at hand is by no means a trivial extension of previous works. In an accretion flow, where the flow is subsonic at a large distance and is necessarily supersonic on the horizon, the flow has to first become supersonic at a sonic point and then after the shock transition where it becomes sub-sonic, the flow must again pass through the inner sonic point before entering into the black hole. In studying flows with constant energy (Chakrabarti, 1989b, hereafter C89b) or an isothermal flows (C89a), both the sonic points were known when the so-called `eigen-values', namely, the specific energy (for polytropic flow) or temperature (isothermal flow) and the specific angular momentum, are supplied. In the present situation, neither of these two quantities is constant in the flow since the viscosity will heat up the gas, increase thermal energy and at the same time reduce the specific angular momentum as the flow proceeds towards the black hole. Thus, the inner sonic point, through which the flow will pass after the shock, is not known before the entire problem is actually solved. We have devised a novel way to solve the entire problem by iterating the location of the inner sonic point till the shock condition is satisfied. We have identified the topologies which are essential for shock formation. We have also identified the parameter space which will have solutions with three sonic points but need not have standing shocks. These solutions generally produce oscillating shocks as shown by Ryu, Chakrabarti and Molteni, (1997). In C96a, some effects of viscous heating were studied and cooling effect was chosen to be proportional to the heating effect for simplicity. No parameter space study was made. In the present paper, we ignore cooling completely. Exact effects of various cooling processes and their influence on the parameter space will be discussed elsewhere (Das \\& Chakrabarti, 2003).  The plan of the present paper is the following: in the next Section, we present the model equations. In \\S 3, we present the sonic point analysis. In \\S 4, we  study the global solution topology. In \\S 5, we classify the parameter space in terms of whether a global solution has triple sonic points or not. In \\S 6, we classify the region with triple sonic points further to indicate which region may allow standing shocks and which region may allow oscillating shocks in presence of viscosity. We showed in particular that matter with a very low angular momentum may allow shocks even when the viscosity parameter is very high. In C96a it was shown that topologies are changed with viscosity and there exists two critical viscosity parameters at which such changes take place.  In \\S 7, we quantify these critical viscosity parameters. Finally, in \\S 8, we discuss about the relevance of shock shocks in the context of quasi-periodic oscillations and make concluding remarks. ",
        "conclusions": "In this paper, we have extended our earlier results of the study of shock formation to include a very difficult yet more realistic case of viscous polytropic flows. Some of the results have been touched upon in C96a but the new results in our work include a detailed study of the parameter space in which shocks form even in presence of viscosity. We found a large number of important results: (a) That there exists two critical viscosity parameters which separate the region of the parameter space in three parts -- 1) in which the flow has a Bondi-type single sonic point; 2) in which there are three sonic points but no Rankine-Hugoniot relations are satisfied and 3) when Rankine-Hugoniot relations are satisfied. These critical viscosity parameters decrease with the increase of the specific angular momentum of the flow at the inner sonic point. (b) That at high viscosities, standing and oscillating shocks may form if the flow has very little angular momentum at the inner sonic point, while at low viscosities the situation is exactly the opposite. It is widely believed that accreting matter on galactic and extra-galactic black holes could be of very low angular momentum, especially when the central compact object is accreting winds from the nearby star or stars. This brings out the possibility that shocks may be active ingredients of an accretion flow. Our results, with a very plausible accretion flow models, indicate that the standing and oscillating shocks are produced even for large viscosity parameters. (c) That the shock location is reduced with enhancement of viscosity parameter. This, coupled to earlier results (Chakrabarti \\& Manickam, 2000) that the infall time is proportional to the period of quasi-periodic oscillation (QPO) of X-rays from black holes imply that the QPO frequency should increase as the viscosity is increased. This is consistent with the observational findings that QPO frequency is increased as the spectral slope softens, widely known to be due to increase in viscosity and accretion rate.  One of the questions we have not addressed here is the stability properties of these shocks. A number of authors pointed out that while the shocks are stable, they should undergo oscillations, either radially, or vertically, or non-axisymmetrically (Molteni, Toth \\& Kuznetsov, 1999; Gu \\& Foglizzo, 2003). We anticipate that our shock solutions in viscous flows would suffer similar types of oscillations, especially when the viscosity is low. In particular, Gu \\& Foglizzo (2003) while studying shocks in inviscid, isothermal flows found such instability and interpreted as due to cycles of acoustic waves between their co-rotation radius and the shock. In their interpretation this could be a form of Papaloizou-Pringle instability (Papaloizou \\& Pringle, 1984) which is known to destabilize accretion tori when the angular momentum gradient is less than a certain value. If so, such an instability could disappear at high enough viscosity. This could have a bearing on the quasi-periodic oscillations of observed X-rays in galactic and extra-galactic black hole candidates in that QPOs would cease to exist above a certain frequency. The interesting aspect is that these so-called 'instability' only causes oscillation of shocks and does not destroy the shock (Molteni, Toth \\& Kuznetsov, 1999).  This work is partly supported by a project (Grant No. SP/S2/K-15/2001) funded by Department of Science and Technology (DST), Govt. of India."
    },
    "0402/astro-ph0402269_arXiv.txt": {
        "abstract": "We present new stellar population models that include the contribution of the Thermally Pulsing Asymptotic Giant Branch (TP-AGB) phase also in the synthetic spectral energy distribution (SED). The TP-AGB phase is essential for a correct modeling of intermediate-age ($0.2~\\lapprox t/{\\rm Gyr} \\lapprox 1\\div2$) stellar populations, because it provides $\\sim~40$~\\% of the bolometric contribution, and up to $\\sim~80$~\\% of that in the $K$-band. These models are obtained by coupling the energetic of the TP-AGB phase as calibrated with data of Magellanic Clouds star clusters (\\cite{io98}), with empirical spectra of TP-AGB stars (\\cite{LM02}). Now that the Spitzer Space Telescope (SST) allows the sight of the rest-frame IR at high redshifts, these models provide the opportunity to use the TP-AGB phase as an age indicator also for high-redshift stellar populations. Here we focus on redshift $\\sim 3$ and provide predictions of the colours of various galaxy models as will be measured by means of the IRAC imaging instrument on board the SST. We find a sizable magnitude difference between TP-AGB-dominated high-redshift stellar populations and those being older or younger. The first releases of GOODS data should allow a check of these predictions. ",
        "introduction": "The epoch(s) of galaxy formation is constrained by dating the stellar populations at zero as well as at high redshift, because the timescales of stellar evolution are independent of cosmological models. Such a constraint provides an important check of current models of hierarchical galaxy formation, in the framework of which the assembly of massive galaxies appears to occur over a rather extended redshift range, with a substantial amount of star formation at redshift lower than 1 (see reviews by S.~White and R.~Sommerville, {\\it this volume}). Such prediction appears to be at odd with the old average age, and the small spread in ages, derived for local massive ellipticals (Es) using optical absorption features and chemical evolution arguments (Thomas, Maraston, Bender, {\\it this volume}; see also G. Gavazzi, {\\it this volume}; \\cite{kauffmann+03}). Also, the finding of galaxies already massive at high redshifts (see contributions by A. Cimatti; R. Genzel, {\\it this volume}; \\cite{saraccoetal03}) is difficult to accommodate in such models. On the other hand, there are also galaxies whose average ages appear to be rather young and could indeed be consistent with small formation redshifts ($z\\lapprox~1$), like the so-called k+a galaxies (\\cite{biancaetal}), or lenticulars and some low-mass Es in the field (see Thomas, Maraston, Bender, {\\it this volume}).  However, the dating of unresolved stellar populations by means of spectro-photometric indicators in the optical is limited by the well-known {\\it age/metallicity degeneracy} (\\cite{F72}, \\cite{W94}, \\cite{MT00}), i.e. the phenomenon that a low metallicity can mimick a low age, and vice versa. When a stellar population ages above $1\\div 2$~Gyr, the Red Giant Branch and the Main Sequence share almost equally the energy production (see e.g. Figure~3 in \\cite{io98}), and their contributions evolve very smoothly with age. At the same time there are no other stellar phases of short duration and relevant energetics that become important and could be used as age indicators. Therefore, at old ages the age/metallicity degeneracy works at its best in confusing the age determination. To trace back the beginning of the formation of a stellar system it would be ideal to recognize it before it becomes a few Gyr old.  A clear signpost of intermediate age ($t\\sim 1~\\rm Gyr$) stellar populations are Thermally-Pulsing Asymptotic Giant Branch stars (TP-AGB; \\cite{RB86}, \\cite{io98}). According to stellar evolution, the TP-AGB stellar phase, the brightest and the coolest on the HR diagram, becomes fully developed in stars with degenerate carbon-oxygen cores. The onset of such event in the life of a stellar population has been called the AGB-{\\it phase transition} (\\cite{RB86}). The observational evidence of the onset of the TP-AGB is a sizable jump in the V$-$K colour that increases from $\\sim~1.4$~to $\\sim~3.2$, as observed among the Magellanic Clouds globular clusters (see Section~2). The narrow interval in evolutionary mass ($1.5 -3~\\Msun$) constrains the whole duration of the TP-AGB dominance to be $\\sim$~1 Gyr (\\cite{io98}). Therefore, picking up the TP-AGB is a powerful way of dating a stellar population, and this technique has been applied with success to local stellar populations (\\cite{ioetal01}).  How can we extend this approach of age dating to $z>0$ ?  An early suggestion in this direction is due to \\cite{alvio92}, who indicated that the AGB phase transition potentially is an effective tool to date the high redshift formation of Es. Two factors has hampered the exploitation of this idea until now. First, the rest-frame IR at redshift $2\\div3$ is sampled by the observed frame around $8-10~\\mu$m. This window is only now available thanks to the advent of the S(pitzer)S(pace)T(elescope). Second, synthetic spectral energy distributions (SEDs) including the TP-AGB phase are clearly required, but usually evolutionary population synthesis models include only the early part of the AGB phase (so-called the Early-AGB), thereby missing the TP-AGB that is the one energetically important (see Section~2).  In this contribution we introduce model SEDs that include the TP-AGB phase (Section~3) and show the substantial effect on the integrated SEDs of intermediate-age stellar populations. In Section~4 we construct diagnostic colour-colour diagrams for the imaging instrument IRAC on board the SST for the illustrative redshift of 3. We further discuss the use of these models at every redshift. ",
        "conclusions": ""
    },
    "0402/astro-ph0402575_arXiv.txt": {
        "abstract": "We have analyzed the soft X-ray emission in a wide area of the Sculptor supercluster by using overlapping ROSAT PSPC pointings. After subtraction of the point sources we have found evidence for extended, diffuse soft X-ray emission. We have investigated the nature of such extended emission through the cross-correlation with the density of galaxies as inferred from the M\\\"unster Redshift Survey. In particular we have analyzed the correlation as a function of the temperature of the X-ray emitting gas. We have found a significant correlation of the galaxy distribution only with the softest X-ray emission (0.1--0.3~keV) and only for gas temperatures kT $< 0.5$ keV. We have excluded that this soft X-ray diffuse emission, and its correlation with the galaxy distribution, is significantly contributed by unresolved AGN, group of galaxies or individual galaxies. The most likely explanation is that the soft, diffuse X-ray emission is tracing Warm-Hot Intergalactic Medium, with temperatures below 0.5 keV, associated with the large-scale structures in the Sculptor supercluster. ",
        "introduction": "Cosmological simulations predict the formation at low redshifts ($\\rm{z}\\,< 1$) of a diffuse gas phase with temperatures of the order of $T \\sim 10^{5.5}\\div 10^7 \\,\\rm{K}$ and typical densities 10--30 times the mean baryonic density (although 30\\% of this gas can exceed overdensities greater than 60, and even greater than 100 in the proximity of clusters of galaxies). This gas phase should be distributed in large-scale filamentary structures connecting virialized structures \\citep{cen,dave}. Such Warm-Hot Intergalactic Medium (WHIM) has been identified as the main contributor to the missing matter in the baryonic census, i.e. $\\sim 36 \\pm 11$ per cent of the baryons \\citep{fukugita_last}  \\footnote{In this value are comprised both the low redshift Lyman~$\\alpha$ forest and the OVI absorbers whose indipendent contribution to the cosmic baryonic fraction is still subject to uncertainties due to the possible double counting of the absorbers, since both phases can coexist in the same systems.}. The formation of these warm gaseous filaments is due to the infall of baryonic matter onto the previously formed dark matter cosmic web. The gravitational potential of the dark matter heats the gas through shocks and triggers the formation of galaxies. The WHIM can be observed in the soft X-rays \\citep[below $\\sim2\\,\\rm{keV}$;][]{croft} as low surface brightness structures. The detection of its radiation is very difficult because of many Galactic foregrounds (such as the Local Hot Bubble --LHB-- and the Galactic halo) and extragalactic background due AGNs, groups of galaxies and clusters. Simulations and X-ray background studies have shown that the WHIM continuum emissivity below $2\\,\\rm{keV}$ is roughly of the same order of magnitude as the Galactic foregrounds. More specifically $\\rm F_{0.5-2keV}(WHIM)\\approx 7\\, keV\\, s^{-1}\\,cm^{-2}\\,sr^{-1}\\,keV^{-1}$ \\citep[][]{croft,kuntz} and $\\rm F_{0.2-0.3keV}(WHIM)\\approx 15\\, keV\\, s^{-1}\\,cm^{-2}\\,sr^{-1}\\,keV^{-1}$ (Croft, private communication). Within this context, \\citet{pierre} showed, from simulated observations that {\\em XMM} can observe strong filaments up to $\\rm{z}\\sim 0.5$ in the $0.4-4\\,\\rm{keV}$ energy band. \\\\ \\citet{cen1995} pointed out that this gas phase should also emit characteristic spectral lines mainly due to Oxygen, Neon and Iron ions. The level of emissivity of these spectral features is below the sensitivity and spectral resolution limits of the current X-ray instruments. However, cosmological simulations show that these lines will be detectable with the future generation of X-ray satellites \\citep{yoshikawa,fang}. \\\\ Various detections of (continuum) WHIM emission have been claimed, either obtained by observing soft X-ray structures in galaxy overdense regions \\citep{scharf,bagchi,zappacosta}, or by detecting a soft X-ray excess in clusters of galaxies \\citep[][]{kaastra,finoguenov}, or in their proximity \\citep{tittley,soltan}, or through shadowing effects \\citep{bregman}. These observations have been possible by means of X-ray satellites very sensitive to low energies ($<$ 1--2~keV), such as {\\it ROSAT} and XMM.\\\\ Theoretical works had predicted that the WHIM should be detectable through UV and X-ray absorption lines imprinted on the spectra of background QSOs \\citep{hellsten,perna}. The detectability of such absorption features does not depend on the brightness of the filaments but on their column density and on the brightness of the background QSO. So far several detections have been reported through X-ray and far-UV absorption lines probing the hot and cool phase of the WHIM \\citep{nicastro,tripp,mathur}.  \\begin{figure} \\begin{center} \\includegraphics[angle=270, width=0.45\\textwidth]{Fig1.ps} \\caption{The position of the 10 partially overlapping {\\it ROSAT} PSPC pointings in the region of the SSC. For each pointing the {\\it ROSAT} observation ID and the exposure time are shown. The position of the South Galactic Pole (SGP) is also shown. The three deepest pointings (that will be used for the subsequent quantitative analysis) are identified with thick circles.} \\label{exposures} \\end{center} \\end{figure}  Simulations show that the WHIM should be distributed in filamentary structures extending over several tens of Mpc and connecting clusters of galaxies. Therefore, superclusters (hosting several clusters) are optimal regions where WHIM is more likely to be detected. In this work we focus on the Sculptor supercluster \\citep[hereafter SSC,][]{schuecker2,seitter}. This is one of the richest local superclusters, comprising more than 20 clusters  of galaxies \\citep{einasto} spread over a projected length of more than 140 Mpc at a redshift $z \\sim 0.105$\\footnote{Here and in the rest of the paper we will assume a cosmology with $\\Omega_{m} = 0.3$, $\\Omega_{\\Lambda} = 0.7$ and $\\rm{H_{0}} = 70\\,Km\\,s^{-1}\\,Mpc^{-1}$}. It is located in the south Galactic pole, a region where the Galactic hydrogen column density is low enough ($N_H \\sim 1.5\\pm 0.2\\,\\times10^{20} \\,\\rm{cm^{-2}}$ ) to avoid significant effects of patchy absorption that could mimic a pattern of apparent X-ray structures \\citep[see][ for more details]{zappacosta}. The SSC has already been observed by \\citet{spiekermann} and \\citet{obayashi}, using {\\it ROSAT} and ASCA data, with the purpose of detecting large-scale X-ray diffuse emission. They did not find indications for emission extended in large-scale structures. However, \\citeauthor{spiekermann} focused the analysis to the relatively energetic bands at $0.5-3$~keV (without investigating the correlation with the galaxy distribution), whereas \\citeauthor{obayashi} observed in an even harder band ($0.8-10$~keV) and by using pointings centered on clusters with the goal of detecting hot diffuse emission in their outskirts.   In this paper we present evidence for a correlation between the galaxy distribution and soft X-ray emission in the central region of the SSC. In particular we show that galaxies and the softest X-ray emission ($< 0.3$~keV) correlate in regions with gas temperatures kT$< 0.5$~keV. This finding is interpreted as WHIM emission associated with the large-scale structures in the SSC. ",
        "conclusions": "We have investigated the emission from Warm-Hot Intergalactic Medium (WHIM) associated with large-scale structures in the central region of the Sculptor supercluster ($\\rm{z}\\,\\approx 0.1$). Ten overlapping {\\it ROSAT} PSPC fields, covering the central 8.3$\\times$6.4~deg$^2$ of the supercluster, were analysed. After removal of the point sources, the {\\it ROSAT} maps show indication of diffuse, filamentary structures, in some cases connecting known clusters of the Sculptor. The diffuse emission spans a wide range of X-ray spectral shapes: from relatively hard emission expected for clusters and unresolved AGNs, to very soft emission expected for WHIM.  To investigate the nature of the diffuse X-ray emission we have cross correlated the X-ray flux with the density of galaxies obtained from the M\\\"{u}nster redshift catalog (whose galaxies mostly belong to the Sculptor supercluster in this region). The correlation has been analyzed as a function of the gas temperature (or X-ray spectral shape). The most important result is the finding of a significant correlation between the diffuse soft (0.1--0.3~keV) X-ray flux and the density of galaxies at the coolest gas temperatures (kT$<$0.5~keV). Such a correlation is interpreted as emission by WHIM associated with the galaxy distribution.  We have also investigated the possible contribution to the diffuse soft X-ray emission, and to the correlation with galaxies, due to individual galaxies and by cold clusters. We have found that in both cases the contribution is negligible.  We have also detected a weak, marginal correlation between the harder X-ray flux (1.5~keV, R67 band) and the density of galaxies at apparently higher gas temperatures (kT$\\sim$1~keV). The latter correlation is ascribed to slightly obscured, unresolved AGNs."
    },
    "0402/astro-ph0402096_arXiv.txt": {
        "abstract": "{ We present here results obtained from three \\sax\\ observations of the accretion-powered X-ray pulsar SMC~X-1 carried out during the declining phases of its 40--60 days long super-orbital period. Timing analysis of the data clearly shows a continuing spin-up of the neutron star. Energy-resolved timing analysis shows that the pulse-profile of SMC~X-1 is single peaked at energies less than 1.0 keV whereas an additional peak, the amplitude of which increases with energy within the MECS range, is present at higher energies. Broad-band pulse-phase-averaged spectroscopy of the \\sax\\ data, which is done for the first time since its discovery, shows that the energy spectrum in the 0.1--80 keV energy band has three components, a soft excess that can be modeled as a thermal black-body, a hard power-law component with a high-energy exponential cutoff and a narrow and weak iron emission line at 6.4 keV. Pulse-phase resolved spectroscopy indicates a pulsating nature of the soft spectral component, as seen in a few other binary X-ray pulsars, with a certain phase offset with respect to the hard power-law component. Dissimilar shape and phase of the soft and hard X-ray pulse profiles suggest a different origin of the soft and hard components. ",
        "introduction": "The bright, eclipsing, accretion-powered binary X-ray pulsar SMC~X-1 was first detected during a rocket flight (Price et al. 1971). The discovery of X-ray eclipses with the $Uhuru$ satellite established the binary nature of SMC~X-1. The pulsar, with a pulse period of 0.71 s (Lucke et al. 1976), is orbiting a B0I super-giant (Sk 160) of mass $\\sim$ 19 M$_\\odot$ (Primini et al. 1977) with an orbital period of $\\sim$ 3.89 days (Schreier et al. 1972). Since its discovery, observations with various X-ray observatories clearly show a steady spin-up of the neutron star in the binary system. This makes SMC~X-1 an exceptional X-ray pulsar in which no spin-down episode has been observed (Wojdowski et al. 1998). An observed decay in the orbital period with a time scale of 3 $\\times$ 10$^{6}$ yr (Levine et al. 1993) is interpreted as due to tidal interaction between the neutron star and the binary companion. The later is presumed to be in the hydrogen shell burning phase of its evolution. A super-orbital period of 40--60 days in SMC~X-1, analogous to the well known X-ray pulsars Her~X-1 and LMC~X-4, was suggested by Gruber \\& Rothschild (1984) and was confirmed by recent observations with the $RXTE$/ASM and $CGRO$/BATSE (Clarkson et al. 2003). Varying obscuration by a precessing accretion disk provides a good explanation for the long term quasi-periodic intensity variations.  Although the continuum energy spectrum of accreting X-ray pulsars is described by a power-law component with an exponential cutoff (White et al. 1983), there are some binary X-ray pulsars which show the presence of a soft excess over the extended hard power-law component. The soft component is detectable only in pulsars which do not suffer from absorption by material along the line of sight (Paul et al. 2002 and references therein). Pulsations in the soft spectral component with a certain phase difference with respect to the hard component are also seen in a few X-ray pulsars (Her~X-1: Oosterbroek et al. 1997, 2000, Endo et al. 2000, SMC~X-1: Paul et al. 2002, LMC~X-4: Naik \\& Paul 2004). Apart from the hard and soft spectral components, iron emission line features are also seen in many of the X-ray pulsars. Iron K shell emission lines in X-ray pulsars are believed to be produced by illumination of neutral or partially ionized material in accretion disk, stellar wind of the high mass companion, material in the form of circumstellar shell, material in the line of sight, or in the accretion column. Pulse-phase-averaged and pulse-phase-resolved spectroscopy, therefore, provide important information in understanding these systems in more detail.  The hard X-ray spectrum (20--80 keV energy band) of SMC~X-1, obtained from the High Energy X-ray Experiment ($HEXE$) observations, was fitted with a thin thermal bremsstrahlung spectrum with a plasma temperature of $\\sim$ 14.5 keV (Kunz et al. 1993). Though a pure power law spectrum was rejected, a power law component modified with an exponential cutoff also provided a good fit to the $HEXE$ data. The broad-band X-ray spectrum (0.2 -- 37 keV) of SMC~X-1 was earlier studied by fitting combined spectra obtained from the $ROSAT$ and $Ginga$ observations (Woo et al. 1995). The energy spectrum is best fitted with a model consisting of a cutoff power-law type component, soft excess which is modeled as a single blackbody component, and a broad iron emission line. Pulsating hard X-rays and a non-pulsating soft X-rays were detected from observations made with $HEAO~1$ A-2 and $Einstein$~SSS (Marshall et al. 1983). However, pulse-timing analysis of the $ROSAT$ and $ASCA$ observations shows clear pulsations of the soft X-rays with a pulse profile different to that of the hard component (Wojdowski et al. 1998). Pulse-phase-resolved spectroscopy of $ASCA$ data in 0.5 -- 10.0 keV energy band shows a pulsating nature of the soft component with some phase difference compared to the hard X-rays (Paul et al. 2002) as is seen in some other binary X-ray pulsars. The nearly sinusoidal single peaked profile of the pulsating soft component contrasts with the double peaked profile seen at higher energies. As the $ASCA$ GIS spectrometers are not sensitive at energies where the soft excess dominates ($<$ 0.6 keV), it is interesting to probe the nature of the soft spectral component over the pulse period of the 0.7 s pulsar in SMC~X-1 with the \\sax\\ LECS.  In this paper, we present the broad band X-ray spectrum of SMC~X-1 over three decades in energy. We have carried out detailed timing and spectral analysis of three observations of SMC~X-1 with the Low Energy Concentrator Spectrometers (LECS), Medium Energy Concentrator Spectrometers (MECS) and the hard X-ray Phoswich Detection System (PDS) instruments of \\sax\\ in the energy band of 0.1--80.0 keV during decaying state of the 40--60 days super-orbital period of SMC~X-1. To examine nature of the soft excess, pulse-phase-resolved spectral analysis has been carried out for the observation with highest X-ray intensity. In the subsequent sections we give details of the observations, the results obtained from the timing and spectral analysis, followed by a discussion on the results obtained from these three \\sax\\ observations.  \\begin{figure}[t] \\vskip 6.2 cm \\special{psfile=asm_lc.ps vscale=58 hscale=42 hoffset=-45 voffset=310 angle = -90} \\caption{The RXTE-ASM light curve of SMC~X-1 from 1996 December 03 (MJD 50420) to 1997 June 21 (MJD 50620). The arrows mark the dates of the \\sax\\ observations which are used for the analysis.}\\label{long} \\end{figure} ",
        "conclusions": "\\subsection{Pulse period evolution of SMC~X-1}  Accurate pulse period measurement of a number of X-ray pulsars has been achieved over last three decades using various X-ray observatories (Bildsten et al. 1997). X-ray pulsars which accrete matter from the stellar wind of the companion star often show irregular spin rate changes on longer time scales, whereas the disk accreting pulsars generally show long-term systematic changes in spin period. On the shortest time scales, however, the change in spin period appears to be comparable in both the groups of X-ray pulsars.  In the standard accretion-disk model, a pulsar can spin at an equilibrium period if the spin-up torque given by the accreting matter is balanced by a braking torque due to the interaction of the magnetic field with the accretion disk outside the corotation radius (Ghosh and Lamb 1979). For a neutron star with given magnetic moment, the equilibrium spin period depends on the accretion rate and the pulsar is expected to spin-up or down as the accretion rate increases or decreases. The observed correlations between the pulse period and the luminosity of X-ray pulsars establish the consistency of the model. Using hydromagnetic equations, Ghosh \\& Lamb (1979) calculated the torque on the neutron star and found that for sufficiently high stellar angular velocities or sufficiently low mass accretion rates the rotation of the star can be braked while accretion continues.  The observed mean spin-up rate in SMC~X-1 system makes it unique among the close binary systems with supergiant companion in which mass accretion takes place from stellar wind. The period evolution of SMC~X-1 is quite different from other persistent HMXB pulsars. $BATSE$ observations (Bildsten et al. 1997) showed that accreting pulsars with massive companions (eg. Cen~X-3) show short term spin-up and spin-down episodes. Though Cen~X-3 shows 10--100 day intervals of steady spin-up and spin-down trend at a much larger rate, it also shows a long term spin-up trend which is the average of the frequent transition between spin-up and spin-down episodes (Finger et al. 1994). In SMC~X-1, however, the absence of spin-down (torque reversal) episodes in more than three decades makes it different from other pulsars which show long term spin-up trend. The monotonous decrease in the derived pulse period of SMC~X-1 with time suggests that the accretion flow has never slowed enough to allow any breaking in the neutron star rotation, and that SMC~X-1 is far from an equilibrium rotator. It can be noted that considering the low metallicity of the SMC/LMC, the wind of the supergiant companions alone cannot account for the large persistent X-ray luminosities of the pulsars like SMC~X-1 and LMC~X-4. Roche Lobe overflow as a partial accretion mechanism is a distinct possibility in these binary systems, which is also probable in Cen~X-3.  \\subsection{Broad band X-ray spectrum of SMC~X-1}  Since the detection of X-ray emission from SMC in 1971, the accretion powered high mass X-ray binary pulsar SMC~X-1 has been observed with many different X-ray observatories. Though the pulse phase averaged spectral studies of SMC~X-1 have been done in different energy bands using X-ray data from various instruments such as 20--80 keV from $HEXE$ observations (Kunz et al. 1993), 0.2--37 keV from $ROSAT$ and $Ginga$ observations (Woo et al. 1995), 0.1--10 keV from $Chandra$ observation (Vrtilek et al. 2001), 0.5-10 keV from $ASCA$ observations (Paul et al. 2002), broad band X-ray spectral study in 0.1--80 keV energy range is reported for the first time here. A thermal bremsstrahlung model, used to describe the 20--80 keV hard X-ray spectrum (Kunz et al. 1993) is ruled out while fitting the source spectrum in 0.1--80 keV energy range. A Comptonization continuum component, used to describe the spectrum of a few other accretion powered X-ray pulsars, is also found to be unsuitable for spectral fitting in comparison to the hard power-law continuum component. Simultaneous spectral fitting to the broad band X-ray spectrum of the source, therefore, shows significant improvement in understanding the accretion processes in the binary system.  Broad-band pulse-phase-averaged spectroscopy of SMC~X-1 shows the presence of a weak and narrow iron emission line with very low equivalent width ($\\sim$ 20 eV) and soft excess above the hard power-law continuum component as seen in several accreting pulsars. A detailed and systematic analysis of X-ray spectra of SMC~X-1 at different phases of its 40--60 days super-orbital period would establish the spectral variations of the source over the third period. A correlation between the hard X-ray continuum flux and the iron emission line flux, a highly variable nature of the otherwise constant iron equivalent width during the low intensity states have been found in other X-ray binary pulsars with super-orbital period (LMC~X-4 and Her~X-1, Naik \\& Paul 2003).  \\subsection{Nature of the soft excess}  Accreting X-ray binary pulsars which do not suffer from strong absorption by the material along the line of sight show soft excess over the hard power-law component. Her~X-1 (Endo et al. 2000), SMC~X-1 (Paul et al. 2002), EXO~053109--6609.2 (Haberl et al. 2003, Paul et al. 2004), and LMC~X-4 (Naik \\& Paul 2004) are the sources in which the difference in the pulse profiles at soft and hard X-ray bands along with the presence of a soft component over the dominating hard power-law component are already reported. Some of the sources also show  pulsations in the soft component. The pulsating nature of the soft blackbody component with a certain phase difference compared to the hard component and heterogeneous pulse profiles at different energy bands suggest different origin of emission of the soft and hard components. Endo et al. (2000) discussed about the origin of the soft and hard spectral components in Her~X-1 and suggested that the hard power-law component originates from the magnetic poles of the neutron star in the binary system whereas the origin of the soft blackbody component is believed to be the inner edge of the accretion disk. A blackbody or thermal bremsstrahlung type emission component fits the soft excess of SMC~X-1 and LMC~X-4 (Paul et al. 2002). However, from \\sax\\ observation of LMC~X-4 in the high state, Naik \\& Paul (2004) have established a pulsating nature of the soft component which rules out the bremsstrahlung model.  The soft spectral component derived from the \\sax\\ observations is entirely compatible with the results from the ASCA spectra (Paul et al. 2002). However, a short exposure of only 7.5 ks with the LECS during this SAX observation does not allow us to determine accurately the shape and phase of the soft component. Therefore, from the present observation we cannot rule out a nonvarying soft excess."
    },
    "0402/astro-ph0402605_arXiv.txt": {
        "abstract": "The sharp magnetic discontinuities which naturally appear in solar magnetic flux tubes driven by turbulent photospheric motions are associated with intense currents. \\citet{Par83} proposed that these currents can become unstable to a variety of microscopic processes, with the net result of dramatically enhanced resistivity and heating (nanoflares). The electric fields associated with such ``hot spots'' are also expected to enhance particle acceleration. We test this hypothesis by exact relativistic orbit simulations in strong random phase magnetohydrodynamic (MHD) turbulence which is forming localized super-Dreicer Ohm electric fields ($E_\\Omega/E_D$ = $10^2 \\, ... \\, 10^5$) occurring in 2..15 \\% of the volume.  It is found that these fields indeed yield a large amplification of acceleration of electrons and ions, and can effectively overcome the injection problem. We suggest in this article that nanoflare heating will be associated with sporadic particle acceleration. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0402/astro-ph0402119_arXiv.txt": {
        "abstract": "Half of the radio afterglows for which there is a good temporal coverage exhibit after 10 days from the burst a decay which is shallower than at optical frequencies, contrary to what is expected within the simplest form of the standard model of relativistic fireballs or jets. We investigate possible ways to decouple the radio and optical decays. First, the radio and optical emissions are assumed to arise from the same electron population and we allow for either a time-varying slope of the power-law distribution of electron energy or for time-varying microphysical parameters. Then we consider two scenarios where the radio and optical emissions arise in distinct parts of the GRB outflow, either because the outflow has an angular structure or because there is a long-lived reverse shock. We find that the only the last scenario is compatible with the observations. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0402/astro-ph0402433_arXiv.txt": {
        "abstract": "{ We analyse protostellar mass accretion rates $\\dot{M}$ from numerical models of star formation based on gravoturbulent fragmentation, considering a large number of different environments. To within one order of magnitude, $\\dot{M} \\approx M_\\mathrm{J} / \\tau_\\mathrm{ff}$ with $M_\\mathrm{J}$ being the mean thermal Jeans mass and $\\tau_\\mathrm{ff}$ the corresponding free-fall time. However, mass accretion rates are highly time-variant, with a sharp peak shortly after the formation of the protostellar core. We present an empirical exponential fit formula to describe the time evolution of the mass accretion and discuss the resulting fit parameters. There is a positive correlation between the peak accretion rate and the final mass of the protostar. We also investigate the relation of $\\dot{M}$ with the turbulent flow velocity as well as with the driving wavenumbers in different environments.  We then compare our results with other theoretical models of star formation and with observational data. ",
        "introduction": "\\label{sec:intro}  Stars are born in dense cores of interstellar molecular clouds. Despite recent observational and theoretical progress, the initial conditions and physical processes that determine the formation of stars are still not fully understood.  In the so-called ``standard theory of star formation'' (Shu \\etal\\ \\cite{sal87}) stars are formed by the inside-out collapse of a singular isothermal sphere that is initially in quasistatic equilibrium, supported against gravity by magnetic and thermal pressure and evolves only due to slow ambipolar diffusion processes. This model predicts protostellar mass accretion rates that are constant with time and only depend on the isothermal sound speed (Shu \\cite{shu77}).  This hypothesis, however, has been challenged from several sides (see Larson \\cite{larson03} or Mac~Low \\& Klessen \\cite{mor_ralf} for a summary). It only is applicable to isolated, single stars, while it is known that the majority of stars form in small aggregates or large clusters (Adams \\& Myers \\cite{adams_myers}; Lada \\& Lada \\cite{ladalada}). Furthermore, there is both observational evidence (Crutcher \\cite{crutcher}; Andr{\\'e} \\etal\\ \\cite{andre00}; Bourke \\etal\\ \\cite{bourke}) and theoretical reasoning (e.g.\\ Nakano \\cite{nakano98}) showing that most observed cloud cores do not have magnetic fields strong enough to support against gravitational collapse.  Similarly, the long lifetimes implied by the quasi-static phase of evolution in the model are difficult to reconcile, e.g., with observational statistics of cloud cores (Taylor \\etal\\ \\cite{tmw96}; Lee \\& Myers \\cite{lee_my}; Visser \\etal\\ \\cite{vrc02}) and with chemical age considerations (van Dishoeck \\& Blake \\cite{dis_blake}; Langer \\etal\\ \\cite{langer00}).  Molecular clouds appear to actually be transient objects with lifetimes of a few million years that form and dissolve in the larger scale turbulent flow of the Galactic disc (Ballesteros-Paredes \\etal\\ \\cite{bhv99}; Elmegreen \\cite{elmeg00}; Hartmann \\etal\\ \\cite{hbb01}; Hartmann \\cite{hartmann03}; V\\'azquez-Semadeni \\etal\\ \\cite{vsb03}).  Observations of self-similar structure in molecular clouds (e.g. Mac Low \\& Ossenkopf \\cite{maclow_oss00}; Ossenkopf \\& Mac Low \\cite{oss_maclow02}) indicate that interstellar turbulence is driven on scales substantially larger than the clouds themselves.  These large-scale turbulent flows compress and cool gas. At sufficiently high densities atomic gas is then quickly converted into molecular form (Hollenbach \\etal\\ \\cite{hws71}).  These same flows will continue to drive the turbulent motions observed within the newly formed cloud.  Some combination of turbulent flow, free expansion at the sound speed of the cloud, and dissociating radiation from internal star formation will then be responsible for their destruction on a timescale of 5--10~Myr.  The most likely source of such large-scale interstellar turbulence in the Milky Way is the combined energy and momentum input from supernovae explosions. They appear to overwhelm all other possibilities.  In the outer reaches of the Galaxy and in low surface brightness galaxies, on the other hand, the situation is not so clear, with magnetorotational or gravitational instabilities looking most likely to drive the observed flows (Mac~Low \\cite{maclow02}; Mac~Low \\& Klessen \\cite{mor_ralf}).  Modern star formation theory, therefore, considers supersonic interstellar turbulence as controlling agent for stellar birth, rather than mediation by magnetic fields (Mac~Low \\& Klessen \\cite{mor_ralf}). This turbulence typically carries sufficient energy to counterbalance gravity on global scales. On small scales, however, it may actually provoke localised collapse (Hunter \\& Fleck \\cite{hunter_fleck}; Elmegreen \\cite{elmeg93}; Padoan \\cite{padoan95}; Ballesteros-Paredes \\etal\\ \\cite{bvs99}; Klessen \\etal\\ \\cite{khm00}; Padoan \\& Nordlund \\cite{pado_nord99}, \\cite{pado_nord02}).  This apparent paradox can be resolved when considering that supersonic turbulence establishes a complex network of interacting shocks, where converging flows generate regions of enhanced density. The system becomes highly filamentary, with elongated structures being caused either by shear motions or by compression at the intersection of shocked layers of gas. At some locations the density enhancement can be sufficiently strong for gravitational instability to set in. The stability criteria for filaments and sheets have been derived and discussed in the context of star formation, e.g., by Larson (\\cite{larson85}), Lubow \\& Pringle (\\cite{lp93}), and Clarke (\\cite{clarke99}).  However, the same random flow that creates density enhancements may disperse them again.  For local collapse to result in stellar birth, it must progress fast enough for the region to `decouple' from the flow.   The efficiency of protostellar core formation, the growth rates and final masses of the protostars, and essentially all other properties of nascent star clusters then depend on the intricate interplay between gravity on the one hand side and the turbulent velocity field in the cloud on the other. The star formation rate is regulated not just at the scale of individual star-forming cores through ambipolar diffusion balancing magnetostatic support, but rather at all scales (Elmegreen \\cite{elmeg02}), via the dynamical processes that determine whether regions of gas become unstable to prompt gravitational collapse. The presence of magnetic fields does not alter that picture significantly (Mac~Low \\etal\\ \\cite{maclow98}; Stone \\etal\\ \\cite{sog98}; Padoan \\& Nordlund \\cite{pado_nord99}; Heitsch \\etal\\ \\cite{heitsch01}). In particular, it cannot prevent the decay of interstellar turbulence.  Clusters of stars build up in molecular cloud regions where self-gravity overwhelms turbulence, either because such regions are compressed by a large-scale shock, or because interstellar turbulence is not replenished and decays on short timescales.  Then, many gas clumps become gravitationally unstable synchronously to go into collapse. If the number density is high, contracting protostellar cores interact and may merge to produce new cores which now contain multiple protostars.  Close encounters drastically alter the trajectories of the protostars, thus changing their mass accretion rates. This has important consequences for the final stellar mass spectrum (Bonnell \\etal\\ \\cite{bonnell97}; Klessen \\& Burkert \\cite{kless_bur00}, \\cite{kless_bur01}; Bonnell \\etal\\ \\cite{bonnell01a}, \\cite{bonnell01b}; Klessen \\cite{klessen01b}; Bate \\etal\\ \\cite{bbb02}).  Inefficient, isolated star formation will occur in regions which are supported by turbulence carrying most of its energy on very small scales. This requires an unrealistically large number of driving sources and appears at odds with the measured velocity structure in molecular clouds which in almost all cases is dominated by large-scale modes (Mac Low \\& Ossenkopf \\cite{maclow_oss00}; Ossenkopf \\& Mac Low \\cite{oss_maclow02}).  In this paper we extend the analysis of protostellar mass accretion rates from a single case (Klessen \\cite{klessen01a}) to a large series of numerical models of turbulent molecular cloud fragmentation, which essentially cover the entire spectrum of observed star-forming regions, ranging from inefficient and isolated star formation to the fast and efficient build-up of stellar clusters.  These calculations, their numerical realisation, and the adopted parameters are described in Section \\ref{sec:models}.  In Section \\ref{sec:discussion} we discuss our findings. We investigate the mass growth history of all protostars in our set of models and present a simple analytic fit formula for the accretion rate $\\dot{M}$. We discuss our study in relation with previous analyses and observational data in Sections \\ref{sect:comp_mod} and \\ref{sect:obs}, respectively. Finally, in Section \\ref{sec:summary} we summarise our results.       \\section {The models} \\label{sec:models}  To adequately describe the fragmentation of turbulent, self-gravitating gas clouds, and the resulting formation and mass growth of protostars, it is prerequisite to resolve the dynamical evolution of collapsing cores over several orders of magnitude in density. Due to the stochastic nature of supersonic turbulence, it is not known in advance where and when this local collapse occurs. Hence, SPH ({\\em smoothed particle hydrodynamics}) is used to solve the equations of hydrodynamics. It is a Lagrangian method, where the fluid is represented by an ensemble of particles and flow quantities are obtained by averaging over an appropriate subset of the SPH particles (Benz \\cite{benz90}; Monaghan \\cite{monagh92}). The method is able to resolve large density contrasts as particles are free to move and so naturally the particle concentration increases in high-density regions.  We use the same smoothing procedure for gravity and pressure forces. This is one requirement to prevent artificial fragmentation (Bate \\& Burkert \\cite{bb97}). Because it is computationally prohibitive to treat the cloud as a whole, we concentrate on subregions within the cloud and adopt periodic boundary conditions (Klessen \\cite{klessen97}). Once the central region of a collapsing protostellar core exceeds a density contrast of $\\sim 10^5$, it is replaced by a ``sink'' particle (Bate \\etal\\ \\cite{bbp95}), which has the ability to accrete gas from its surrounding while at the same time keeping track of mass and linear and angular momentum. By adequately replacing high-density cores with sink particles we can follow the dynamical evolution of the system over many free-fall times.  \\begin{table*}[t] \\caption{Overview of our models (See text for details)} \\label{tab:models} {\\small \\begin{center} \\begin{minipage}{15cm} \\begin{tabular}{l c c r l r r r r  r r r r} \\hline \\hline Name & \\mach % & $k$ % & \\multicolumn{1}{c}{$n_\\mathrm{p}$\\footnote{number of particles in the simulation}} & $M_\\mathrm{min}$\\footnote{SPH resolution limit} & $M_\\mathrm{accr}$\\footnote{fraction of the total mass that has been accreted by the end of the simulation} & $n_*$\\footnote{total number of formed protostars} & $n_*$\\footnote{number of protostars that can be fitted by Eq.~(\\ref{eq:fit})} & $\\sigma_\\mathrm{mean}$\\footnote{mean deviation of the fits, calculated from Eq.~(\\ref{eq:sigma})} & \\multicolumn{4}{c}{$\\dot{M}_\\mathrm{mean}$ [$10^5 \\mathrm{M_{\\sun} yr^{-1}}$]}\\\\ & & & & \\multicolumn{1}{c}{\\scriptsize [M$_{\\sun}$]} & \\multicolumn{1}{c}{\\scriptsize [\\%]} &  & {\\scriptsize fitted} & & {\\scriptsize bin1} & {\\scriptsize bin2} & {\\scriptsize bin3} & {\\scriptsize bin4} \\\\ \\hline G1    & --& --&  50\\,000 & 0.44  & 93.1 & 56 & 31 & 0.49 & 1.18 & 1.30 & 1.70 & 2.93 \\\\ G2    & --& --& 500\\,000 & 0.044 & 84.9 & 56 & 52 & 0.43 & 0.94 & 1.40 & 2.09 & 3.51\\vspace{0.2cm}\\\\ M01k2 & 0.1 & 1..2 & 205\\,379 & 0.058 & 74.9 & 95 & 91 & 0.43 & 0.77 & 1.76 & 3.04 & --    \\\\ % M01k4 & 0.1 & 3..4 & 205\\,379 & 0.058 & 27.2 &  3 &  3 & 0.81 & 2.83 & --   & --   & 57.98 \\\\ % M01k8 & 0.1 & 7..8 & 205\\,379 & 0.058 & 85.9 &  3 &  3 & 0.47 & --   & 1.37 & 13.25& 59.16 \\\\ % M05k2 & 0.5 & 1..2 & 205\\,379 & 0.058 & 37.2 & 23 & 22 & 0.49 & 1.63 & 5.05 & 4.88 & 12.95 \\\\ % M05k4 & 0.5 & 3..4 & 205\\,379 & 0.058 & 77.9 & 48 & 48 & 0.39 & 1.31 & 2.49 & 2.56 & 6.30  \\\\ % M05k8 & 0.5 & 7..8 & 205\\,379 & 0.058 & 59.5 & 48 & 48 & 0.40 & 1.34 & 2.22 & 3.79 & 7.77  \\\\ % M2k2  &  2  & 1..2 & 205\\,379 & 0.058 & 75.1 & 68 & 62 & 0.41 & 0.86 & 1.38 & 2.54 & 4.32  \\\\ % M2k4  &  2  & 3..4 & 205\\,379 & 0.058 & 47.9 & 62 & 62 & 0.44 & 1.35 & 1.92 & 2.43 & 3.84  \\\\ % M2k8  &  2  & 7..8 & 205\\,379 & 0.058 & 66.2 & 42 & 40 & 0.42 & 0.87 & 1.29 & 1.72 & 3.38  \\\\ % M3k2  & 3.2 & 1..2 & 205\\,379 & 0.058 & 79.7 & 65 & 65 & 0.46 & 1.31 & 1.86 & 2.98 & 3.78  \\\\ % M3k4  & 3.2 & 3..4 & 205\\,379 & 0.058 & 82.1 & 37 & 35 & 0.55 & 1.01 & 1.13 & 1.15 & 1.86  \\\\ % M3k8  & 3.2 & 7..8 & 205\\,379 & 0.058 & 60.2 & 17 & 17 & 0.41 & 0.51 & 1.09 & 1.74 & 5.84  \\\\ % M6k2a &  6  & 1..2 & 205\\,379 & 0.058 & 85.4 &100 & 97 & 0.49 & 0.79 & 1.69 & 1.96 & 3.39  \\\\ % M6k4a &  6  & 3..4 & 205\\,379 & 0.058 & 62.4 & 98 & 93 & 0.44 & 0.38 & 0.83 & 0.89 & 1.10  \\\\ % M6k2b &  6  & 1..2 & 195\\,112 & 0.062 & 34.5 & 50 & 50 & 0.42 & 1.02 & 1.81 & 2.50 & --    \\\\ % M6k4b &  6  & 3..4 &  50\\,653 & 0.24  & 29.7 & 50 & 47 & 0.43 & 0.72 & 1.66 & 1.87 & --    \\\\ % M6k8b &  6  & 7..8 &  50\\,653 & 0.24  & 35.7 & 25 & 25 & 0.44 & 0.35 & 0.61 & 1.38 & 2.47  \\\\ % M6k2c &  6  & 1..2 & 205\\,379 & 0.058 & 75.8 &110 & 97 & 0.43 & 0.83 & 1.32 & 1.50 & 1.23  \\\\ % M6k4c &  6  & 3..4 & 205\\,379 & 0.058 & 61.9 & 53 & 46 & 0.54 & 1.31 & 1.46 & 2.05 & 1.97  \\\\ % M6k8c &  6  & 7..8 & 205\\,379 & 0.058 &  6.4 & 12 & 10 & 0.43 & 0.50 & 0.62 & 0.58 & --    \\\\ % M10k2 & 10  & 1..2 & 205\\,379 & 0.058 & 56.5 &150 &146 & 0.44 & 1.08 & 2.62 & 2.09 & --    \\\\ % M10k8 & 10  & 7..8 & 205\\,379 & 0.058 & 32.4 & 54 & 44 & 0.53 & 0.26 & 0.64 & --   & --    \\\\ % \\end{tabular} \\end{minipage} \\end{center} } \\end{table*}  The suite of models consists of two globally unstable models that contract from Gaussian initial conditions without turbulence (for details see Klessen \\& Burkert \\cite{kless_bur00}, \\cite{kless_bur01}) and of 22 models where turbulence is maintained with constant rms Mach numbers $\\cal M$, in the range $0.1 \\le {\\cal M} \\le 10$. We distinguish between turbulence that carries its energy mostly on large scales, at wavenumbers $1 \\le k \\le 2$, on intermediate scales, i.e.\\ $3 \\le k \\le 4$, and on small scales with $7 \\le k \\le 8$. The corresponding wavelengths are $\\ell = L/k$, where $L$ is the total size of the computed volume.  The models are labelled mnemonically as M$\\cal M$k$k$, with rms Mach number $\\cal M$ and wavenumber $k$, while G1 and G2 denote the two Gaussian runs.  The main parameters are summarised in Table~\\ref{tab:models}. To have well defined environmental conditions given by $\\cal M$ and $k$, $\\cal M$ is required to be constant throughout the evolution. However, turbulent energy dissipates rapidly, roughly on a free-fall timescale (Mac Low \\etal\\ \\cite{maclow98}; Stone \\etal\\ \\cite{sog98}; Padoan \\& Nordlund \\cite{pado_nord99}). We therefore apply a non-local driving scheme that inserts energy at a given rate and at a given scale $k$. We use Gaussian random fluctuations in velocity. This is appealing because Gaussian fields are fully determined by their power distribution in Fourier space.  We define a cartesian mesh with $64^3$ cells, and for each three-dimensional wave number $\\vec{k}$ we randomly select an amplitude from a Gaussian distribution around unity and a phase between zero and $2\\pi$.  We then transform the resulting field back into real space to get a ``kick-velocity'' in each cell.  Its amplitude is determined by solving a quadratic equation such to keep $\\cal M$ constant (Mac~Low \\cite{maclow99}; Klessen \\etal\\ \\cite{khm00}). The ``kick-velocity'' is then simply added to the speed of each SPH particle located in the cell. We adopted this method for mathematical simplicity. In reality, the situation is far more complex. Still, our models of large-scale driven clouds contain many features of molecular clouds in supernovae driven turbulence (e.g. Ballesteros-Paredes \\& Mac~Low \\cite{bpml02}; Mac~Low \\etal\\ \\cite{mak03}). Conversely, our models of small-scale turbulence bear certain resemblance to energy input on small scales provided by protostellar feedback via outflows and winds.  Our models neglect the influence of magnetic fields, because their presence cannot halt the decay of turbulence (Mac~Low \\etal\\ \\cite{maclow98}; Stone \\etal\\ \\cite{sog98}; Padoan \\& Nordlund \\cite{pado_nord99}) and does not significantly alter the efficiency of local collapse for driven turbulence (Heitsch \\etal\\ \\cite{hmk01}). More importantly, we do not self-consistently consider feedback effects from the star formation process itself (like bipolar outflows, stellar winds, or ionising radiation from new-born O or B stars). Our analysis of protostellar mass accretion rates solely focuses on the interplay between turbulence and self-gravity only. This is also the case in the Shu (\\cite{shu77}) theory of isothermal collapse. Hence, our findings can be directly compared to the ``standard theory of star formation''.  The models are computed in normalised units using an isothermal equation of state. Scaled to physical units we adopt a temperature of 11.3$\\,$K corresponding to a sound speed $c_{\\rm s} = 0.2\\,$km$\\,$s$^{-1}$, and we use a mean density of $n({\\rm H}_2) = 10^5\\,$cm$^{-3}$, which is typical for star-forming molecular cloud regions (e.g.\\ in $\\rho$~Ophiuchi, see Motte \\etal\\ \\cite{man98}). The total mass contained in the computed volume in the two Gaussian models is 220$\\,$M$_{\\sun}$ and the size of the cube is $0.34\\,$pc. This corresponds to 220 thermal Jeans masses. The turbulent models have a mass of 120$\\,$M$_{\\sun}$ within a volume of ($0.28\\,{\\rm pc})^3$, equivalent to 120 thermal Jeans masses\\footnote{We use a spherical definition of the Jeans mass, $M_{\\rm J} \\equiv 4/3\\,\\pi \\rho (\\lambda_{\\rm J}/2)^3$, with density $\\rho$ and Jeans length $\\lambda_{\\rm J}\\equiv \\left(\\frac{\\pi{\\cal R}T }{G \\rho}\\right)^{1/2}$ and where $G$ and $\\cal R$ are the gravitational and the gas constant. The mean Jeans mass $\\langle M_{\\rm J} \\rangle$ is then determined from the average density in the system $\\langle \\rho \\rangle$. }. The mean thermal Jeans mass in all models is thus $\\langle M_{\\rm J} \\rangle = 1\\,$M$_{\\sun}$, the global free-fall timescale is $\\bar{\\tau}_{\\rm ff} = 10^5\\,$yr, and the simulations cover a density range from $n({\\rm H}_2) \\approx 100\\,$cm$^{-3}$ in the lowest density regions to $n({\\rm H}_2) \\approx 10^9\\,$cm$^{-3}$ where collapsing protostellar cores are identified and converted into ``sink'' particles in the code.  This coincides in time with the formation of the central protostar to within $\\sim10^3\\,$yr (Wuchterl \\& Klessen \\cite{wucht_kless01}). The resolution limit for each model, requiring that the local Jeans mass is always resolved by at least 100 gas particles (Bate \\& Burkert \\cite{bb97}), is given in Col.~5 of Table~\\ref{tab:models}.  In the subsequent protostellar phase of evolution, we determine accretion rates $\\dot{M}$ by measuring the amount of mass as function of time that falls into a control volume defined by each ``sink'' particle.  Its diameter is fixed to $560\\,$AU. Entering gas particles pass through several tests to check if they remain bound to the ``sink'' particle (Bate \\etal\\ \\cite{bbp95}) before they are considered accreted. As all gas particles have the same  mass and as accretion events occur at random times, the resulting accretion rates are mass-binned and we smooth over a few consecutive accretion events to get a description of the time evolution of $\\dot{M}$. We cannot resolve the evolution in the interior of the control volume. Because of angular momentum conservation most of the matter that falls in will assemble in a protostellar disc. There it is transported inwards by viscous and possibly gravitational torques. The latter will be provided by spiral density waves that develop when the disc becomes too massive, which happens when mass is loaded onto the disc faster than it is removed by viscous transport alone. Altogether, the disc will not prevent or delay material from accreting onto the protostar for long. It acts as a buffer and smoothes eventual accretion spikes. For the mass range considered here also feedback effects are too weak to halt or delay accretion.  With typical disc sizes of order of several hundred AU, the control volume therefore fully encloses both, star and disc, and the measured core accretion rates are good estimates for the actual stellar accretion rates.  Deviations may be expected only if the protostellar core forms a binary star, where the infalling mass must then be distributed between two stars, or if very high-angular momentum material is accreted, where a certain mass fraction may end up in a circumbinary disc and not accrete onto a star at all.  In the prestellar phase, i.e.\\ before the central protostar forms, we determine the accretion history by computing the change of mass inside the control volume centered on the SPH particle that turns into a ``sink'' during the later evolution. Turbulent compression leads to mass growth, i.e.\\ $\\dot{M}>0$, while expansion will result in mass loss and $\\dot{M}<0$. Appreciable mass growth, however, is only achieved when gravity takes over and the region goes into collapse. ",
        "conclusions": "\\label{sec:discussion}  \\subsection{First approximation to $\\dot{M}$}   The entire process of molecular cloud collapse and build-up of the stellar cluster lasts several global free-fall times ($\\bar{\\tau}_\\mathrm{ff} = 10^5$~yr). Likewise, the accretion process of a protostellar core takes place on a timescale of a few $\\bar{\\tau}_\\mathrm{ff}$, comparable to most other models of star formation.  A simple approximation to the accretion rate can be achieved by dividing the local Jeans mass by the local dynamical timescale: \\begin{equation} \\dot{M} \\approx M_\\mathrm{J}/\\tau_\\mathrm{ff} \\label{eq:jeans} \\end{equation} By substituting \\begin{equation} M_\\mathrm{J} = \\frac{\\pi^{5/2}}{6} \\rho_0^{-1/2} \\left(\\frac{\\mathcal{R} T}{G}\\right)^{3/2} =  \\frac{\\pi^{5/2}}{6} \\rho_0^{-1/2} G^{-3/2} c_\\mathrm{s}^3, \\end{equation} where $\\rho_0$ denotes the initial density, $T$ the temperature and $c_\\mathrm{s}$ the iso\\-thermal sound speed, and \\begin{equation} \\tau_\\mathrm{ff} = \\sqrt{\\frac{3\\,\\pi}{32\\,G \\rho_0}} \\end{equation} Eq.~(\\ref{eq:jeans}) can be written as \\begin{equation} \\dot{M} \\approx \\frac{M_\\mathrm{J}}{\\tau_\\mathrm{ff}} = \\sqrt{\\frac{32}{3}} \\frac{\\pi^2}{6} \\frac{c_\\mathrm{s}^3}{G} = 5.4\\,\\frac{c_\\mathrm{s}^3}{G}, \\end{equation} depending only on the isothermal sound speed (or temperature). For a sound speed $c_\\mathrm{s} = 0.2~\\mathrm{km\\,s^{-1}}$ we obtain $\\dot{M} = 10^{-5}~\\mathrm{M_{\\sun}\\,yr^{-1}}$. This is higher than the accretion rate for the collapse of a singular isothermal sphere: $\\dot{M} = 0.975\\,c_\\mathrm{s}^3/G$ (Shu \\cite{shu77}). However, the accretion rates in our models vary with time. Typical peak accretion rates are roughly in the range $(3-50)\\,c_\\mathrm{s}^3/G$ or $5 \\times 10^{-6}$ to $10^{-4}~\\mathrm{M_{\\sun}\\,yr^{-1}}$. The values exceed the approximated value $M_\\mathrm{J}/\\tau_\\mathrm{ff}$ due to external compression in the turbulent flow.   \\begin{figure*}[t] \\centering \\includegraphics[width=17cm]{paper_accrcurve_both.eps} \\caption{Mass accretion rates of nine randomly selected protostellar cores of three different models. Left panel: $\\dot{M}$ versus time for a Gaussian collapse (G2; upper row), a turbulent model driven on a large scale (M6k2a; middle row), and a turbulent model driven on a small scale (M6k8b; lower row). The thin line represents the actual simulation, the thick line the fit as described in the text. The deviation $\\sigma$ as given by Eq.~(\\ref{eq:sigma}) is indicated for each object. The dotted line shows the constant accretion rate that would be expected from the classical Shu (\\cite{shu77}) scenario. The dashed line stands for the assumed transition from Class~0 to Class~I. The right panel shows the same protostellar cores as on the left side plotted versus the ratio of accreted to final mass. The final masses (in M$_{\\sun}$) are also given.} \\label{accrcurve} \\end{figure*}  \\subsection{Time-varying mass accretion rates}   We analyse the full mass growth history of all protostellar cores in our models and we find that mass accretion rates from gravoturbulent fragmentation are highly time-variable. Several examples of the accretion rate $\\dot{M}$ are displayed in Fig.~\\ref{accrcurve}, plotted versus time (left panel) and the ratio of accreted to final mass (right panel), respectively. The maximum accretion rate is reached rather rapidly % and is then followed by a somewhat slower decline. In some cases this decline is interrupted by one or more secondary peaks. As shown above, the maximum accretion rate is significantly higher than the constant rate predicted by the classical isothermal collapse model (% plotted as dotted line in Fig.~\\ref{accrcurve}), but it falls below that value in later stages. Due to the dynamical interaction and competition between protostellar cores, the mass accretion rates of cores in a dense cluster are different from those of isolated cores. In the first stage a core accretes local gas from its immediate vicinity. Once the local reservoir is depleted, the core may accrete fresh gas streaming in from farther away or by encounters with non-collapsed gas clumps (see discussion in Klessen \\& Burkert \\cite{kless_bur00}). This results in secondary accretion peaks that are also visible in the right panel of Fig.~\\ref{accrcurve}, where one would expect a single bump in the case of an isolated core. For example, the central graph of the right panel of Fig.~\\ref{accrcurve} nicely shows that this particular protostar accretes only about half of the final mass from its direct environment (first bump), while the rest stems from later accretion events.  The transition phase between Class~0 and Class~I protostars is believed to take place when about half of the final mass has been accumulated (Andr\\'{e} et al.\\ \\cite{andre00}). This time is indicated by the dashed line in Fig.~\\ref{accrcurve}. Typically it takes place during or at the end of the peak accretion phase. It determines the lifetime of Class~0 objects, which will be discussed below.  We define a mean accretion rate \\mmean\\ by averaging $\\dot{M}$ in the mass range $0.1 \\le M/M_\\mathrm{end} \\le 0.8$, with $M_\\mathrm{end}$ being the final mass of the protostar.  This phase typically lasts only a few $10^4$ years. This is short compared to the full accretion history. The bulk of stellar material is therefore accumulated in the short time interval while the system is close to maximum accretion.  \\begin{figure*} \\centering \\includegraphics[width=17cm]{paper_meanaccr.eps} \\caption{Mean accretion rates \\mmean\\ versus final mass (upper panel) and versus time of core formation (lower panel) for the same models as in Fig.~\\ref{accrcurve}. The zero point of the timescale coresponds to the time when gravity is ``switched on''. Note the different timescales on which the formation of the cluster takes place.}   \\label{meanaccr} \\end{figure*}  In Fig.~\\ref{meanaccr}, we plot the mean accretion rates versus final star mass $M_\\mathrm{end}$ and versus time of core formation $t_\\mathrm{form}$, respectively, for the same models as in Fig.~\\ref{accrcurve}.  Not surprisingly, \\mmean\\ increases with increasing stellar mass, and decreases when the core forms later, although this second correlation is not that clear.  In other words, more massive stars have higher mass accretion rates and start to form first. They can grow large, because on average they form in the high-density regions of the cluster centre where they are able to maintain relatively high accretion rates over a long time as more and more gas falls in from the cluster outskirts.  \\subsection{An empirical fit formula for $\\dot{M}$}  One of our aims is to find a simple-to-use fit formula to approximate the accretion process.  The protostellar mass growth history in our models can be fitted empirically in the lin-log diagram by the function \\begin{equation} \\log \\dot{M}(t) = \\log \\dot{M}_0\\,\\frac{\\mathrm{e}}{\\tau}\\,t\\,\\mathrm{e}^{-t/\\tau}% \\label{eq:fit} \\end{equation} with time $t$ and the fit parameters $\\log \\dot{M}_0$ and $\\tau$. This holds for the following conditions: We shift the ordinate by $\\Delta \\log \\dot{M}/({\\rm M_{\\sun}\\,yr^{-1}}) = +7$ and we consider accretion when $\\log \\dot{M} \\ge -7$. The zero point of the timescale is determined once the accretion rate exceeds $\\log \\dot{M} = -7$. The fitted curves are plotted as thick lines in Fig.~\\ref{accrcurve}. Note that the ordinate displays the original values without the applied shift.  If there are secondary accretion peaks, they are typically ignored and levelled out by the routine. The overall quality of the fit can be estimated by the standard deviation \\begin{equation} \\sigma = \\sqrt{\\frac{1}{n-1} \\sum_{t=0}^{n} \\left[ \\dot{M}_\\mathrm{fit}(t) - \\dot{M}(t) \\right]^2} \\label{eq:sigma} \\end{equation} where $\\dot{M}(t)$ is the actual value of $\\dot{M}$ at the time $t$ from our simulation, while $\\dot{M}_\\mathrm{fit}(t)$ denotes $\\dot{M}$ calculated using Eq.~(\\ref{eq:fit}) for the same time. The mean value of $\\sigma$ for each model is given in Col.~9 of Table~\\ref{tab:models}. Prestellar cores where the fit routine fails or where $\\sigma > 1$ are not taken into account in our subsequent analysis. This concerns a wide variety of cores, there is no correlation with the final mass or the time of formation. However, they represent only a small fraction of the total number of objects. The actual numbers of fitted cores are listed in Col.~8 of Table~\\ref{tab:models}.   When interpreting the fit parameter $\\log \\dot{M}_0$, the applied shift has to be taken into account. Thus, $\\log \\dot{M}_\\mathrm{max}^\\mathrm{fit} = \\log (\\dot{M}_0) - 7$ gives the real value of the peak accretion. This parameter is plotted for all protostellar cores and all models versus the respective final mass (Fig.~\\ref{a0_a1}).  A correlation with $M_\\mathrm{end}$ is obvious.  We apply a linear fit in the log-log diagram, which is indicated by the straight line. The fitted peak accretion rates show the same behaviour as the mean accretion rates \\mmean.  The parameter $\\tau$ indicates the time of the maximum of the accretion curve. It is plotted for all protostellar cores in Fig.~\\ref{tau}. In almost all models $\\tau$ shows a correlation with the final mass. The parameter indicates how fast the gas falls in onto the core, therefore we expect it to be related to the local free-fall time and, thus, to the local density at the onset of collapse. It lies in the range $10^4 \\lesssim \\tau \\lesssim 10^5$~yr, which is less than the global free-fall time $\\bar{\\tau}_\\mathrm{ff}$. If we take an average value $\\langle \\tau \\rangle \\approx \\bar{\\tau}_\\mathrm{ff}/3$, this suggests an initial overdensity of almost a factor of ten in the collapsing regions.    \\begin{figure*} \\centering \\includegraphics[width=17cm]{paper_a0_a1_incl_gauss.eps} \\caption{Peak accretion rates ($\\dot{M}_\\mathrm{max}^\\mathrm{fit}$) versus $M_\\mathrm{end}$ for all our models, sorted by Mach number \\mach\\ (top to bottom) and wave number $k$ (left to right). The straight line shows the applied linear fit. Details of the models can be found in Table~\\ref{tab:models}.} \\label{a0_a1} \\end{figure*}  \\begin{figure*} \\centering \\includegraphics[width=17cm]{paper_tau.eps} \\caption{The time of maximum accretion $\\tau$ for all models, arranged analogous to Fig.~\\ref{a0_a1}.} \\label{tau} \\end{figure*}  \\subsection{Class~0 lifetimes and the effect of the turbulent medium}   \\begin{figure*} \\centering \\includegraphics[width=17cm]{paper_class01b.eps} \\caption{The assumed duration of Class 0 phase versus $M_\\mathrm{end}$ for all models, arranged analogous to Fig.~\\ref{a0_a1}. The dotted lines confine the range of this parameter according to observations (Andr\\'{e} et al.\\ \\cite{andre00}).} \\label{class01} \\end{figure*}   \\begin{figure*} \\centering \\includegraphics[width=17cm]{paper_meanaccr_mach_4bins.eps} \\caption{Averaged mean accretion rates \\mmean$_\\mathrm{mean}$ of all models versus Mach number.   For the sake of clarity the upper panel is split into mass bins, while the lower panel shows the same, separated according to the wave numbers.} \\label{meanaccr_mach} \\end{figure*}  We calculate the transition times from Class~0 to Class~I, assumed as described above. This gives the duration of Class~0 phase for each protostar, the values are plotted versus the corresponding final masses in Fig.~\\ref{class01}. The duration of Class~0 phase increases with increasing final mass. Thus, a massive star is marked not only by a higher peak accretion rate but also by a longer time spent in Class~0 phase.   The mean accretion rates \\mmean\\ of all individual protostellar cores of one model are averaged in four mass bins: $0 < M_\\mathrm{end}/{\\rm M_{\\sun}} \\le 0.7$ (bin1), $0.7 < M_\\mathrm{end}/{\\rm M_{\\sun}} \\le 1.5$ (bin2), $1.5 < M_\\mathrm{end}/{\\rm M_{\\sun}} \\le 3$ (bin3), and $M_\\mathrm{end}/{\\rm M_{\\sun}} > 3$ (bin4). The values are given in Cols.~10 to 13 of Table~\\ref{tab:models}. Figure~\\ref{meanaccr_mach} shows the relation of the averaged mean mass accretion rate \\mmean$_\\mathrm{mean}$ to the Mach number for all models, split into mass bins and wave numbers, respectively. Three conclusions can be drawn from the figure: Firstly, there is a trend that \\mmean$_\\mathrm{mean}$ decreases with increasing Mach number. That means that the mean accretion rate is lower, when the rms velocity dispersion (i.e. the turbulent Mach number) of the medium is increased. The stronger support of the turbulent medium against gravitational collapse typically results in a lower mass accretion rate. Secondly, \\mmean$_\\mathrm{mean}$ is higher for higher mass bins. This is consistent with the findings for the individual mean and maximum accretion rates (\\mmean\\ and $\\dot{M}_\\mathrm{max}^\\mathrm{fit}$) discussed above. Finally, though, there is no correlation of \\mmean$_\\mathrm{mean}$ with the wavenumber. Apparently the scale of the driving energy has no influence on the accretion rate.      \\label{sec:summary}  We have studied protostellar mass accretion rates from numerical models of star formation based on gravoturbulent fragmentation. Twenty-four models covering a wide range of environmental conditions from low to high turbulent velocities and different driving scales with a total number of 1325 protostellar cores have been investigated. Our main results may be summarised as follows:  \\begin{enumerate} \\item An order-of-magnitude estimate for mass accretion rates resulting from gravoturbulent fragmentation is given by $\\dot{M} \\approx M_\\mathrm{J} / \\tau_\\mathrm{ff}$ with $M_\\mathrm{J}$ being the mean thermal Jeans mass and $\\tau_\\mathrm{ff}$ the corresponding free-fall time. \\item However, protostellar mass accretion is a highly time-variant process. It can be approximated by the empirical function $\\log \\dot{M}(t) = \\log \\dot{M}_0\\,(\\mathrm{e}/\\tau)\\,t\\,\\mathrm{e}^{-t/\\tau}$. The peak accretion rate is reached during Class 0 stage, shortly after the formation of the core; its value ranges between about $5 \\times 10^{-6}$ and $10^{-4}\\ {\\rm M_{\\sun}\\,yr^{-1}}$. The maximum accretion rate is approximately one order of magnitude higher than the constant rate predicted by the collapse of a classical singular isothermal sphere.  \\item Around the peak accretion phase the mass accretion rates are roughly constant. The mean accretion rates % increase with increasing final mass. More massive stars have higher mass accretion rates and tend to form first.  \\item The same applies to the fitted peak accretion rates, which are also proportional to the final stellar mass.  \\item There is a similar correlation between the duration of Class~0 phase (assuming that half of the final mass is accreted in this phase) and the final mass. \\item \\mmean$_\\mathrm{mean}$ decreases with increasing Mach number of the turbulent environment, but is not correlated with the driving wavenumber. \\end{enumerate}  \\noindent Our results agree well with many other models concerning the time evolution of the mass accretion process and the value of the peak accretion rate. In particular, the accretion rates from our models show an exponential decline, as it is also proposed by Bontemps et al.\\ (\\cite{bontemps}), Myers et al.\\ (\\cite{myers98}) and Smith (\\cite{smith99}, \\cite{smith00}). They also match observational findings like the supposed decline of the mass accretion rate from Class~0 to Class~I phase. We conclude that a theory of star formation based on gravoturbulent fragmentation of molecular clouds is an adequate approach to describe stellar birth in the Milky Way."
    },
    "0402/hep-ph0402276_arXiv.txt": {
        "abstract": "Modifications to the Friedmann equation in brane cosmology can have important implications for early universe phenomena such as inflation and the generation of the baryon asymmetry. We study a simple scenario of chaotic brane inflation where, in a minimal supersymmetric seesaw model, the scalar superpartner of a heavy singlet Majorana neutrino drives inflation and, simultaneously, generates the required lepton asymmetry through its direct out-of-equilibrium decays after the inflationary era. For a gravitino mass in the range $m_{3/2} \\simeq$ 100~GeV~-~1~TeV, we find that successful nucleosynthesis and leptogenesis in this framework require that the 5D Planck mass is in the range $M_{5} \\simeq 10^{10}-10^{13}$~GeV and the reheating temperature $T_{rh} \\simeq 10^{6}-10^{8}$~GeV. ",
        "introduction": "\\label{one}  Today there is a wide consensus that the early universe underwent a period of cosmological inflation \\cite{Lyth:1998xn}. Inflationary era can be regarded as a necessary stage, responsible not only for the observed flatness, homogeneity and isotropy of the present universe, but also for the origin of the density fluctuations as observed by the Cosmic Background Explorer (COBE) and, more recently, the Wilkinson Microwave Anisotropy Probe (WMAP) satellites~\\cite{Bennett:2003bz}. At the end of inflation, the universe was in a cold and low-entropy state and it must has been subsequently reheated to become a high-entropy and radiation-dominated universe. Such a reheating process could occur, for instance, through the coherent oscillations of the inflaton field about the minimum of the potential until the age of the universe equals the lifetime of the inflaton. The latter decays into ordinary particles, which then scatter and thermalize. Besides entropy creation, the right abundance of baryons must be created after the inflationary epoch. This usually poses serious problems in constructing particle physics models which lead simultaneously to a successful inflationary and baryogenesis scenario. In particular, the reheating temperature is typically too low when compared with the grand unification scale, at which baryogenesis is expected to take place in the simplest GUTs. Moreover, any preexisting baryon asymmetry would be erased by the anomalous sphaleron processes \\cite{Kuzmin:1985mm} unless an initial $B-L$ asymmetry is generated.  Another major obstacle in constructing viable supergravity-inspired cosmological models is the overproduction of gravitinos. In conventional scenarios, the gravitino mass is expected to be comparable to the masses of the supersymmetric partners of the standard model particles and, therefore, $m_{3/2}\\lesssim$ a few TeV in order to solve the gauge hierarchy problem. Since the gravitino coupling to matter is suppressed by the Planck mass $M_{P}$ , its lifetime is $\\tau_{3/2} \\sim M_{P}^{2}/m_{3/2}^{3} \\sim 10^{8}(100$~GeV $/m_{3/2})^{3}$~s. During the reheating phase gravitinos can be thermally produced through scatterings in the plasma. However, if they are overproduced after inflation, their decay products could put at risk the successful predictions of primordial nucleosynthesis~\\cite{Khlopov:pf,Cyburt:2002uv}. Since in standard cosmology their abundance is proportional to the reheating temperature, $T_{rh}\\,$, constraints from big bang nucleosynthesis (BBN) yield a stringent upper bound on the allowed $T_{rh}$ after inflation: $T_{rh} \\lesssim 10^{7}-10^{10}$~GeV for 100~GeV $\\lesssim m_{3/2} \\lesssim 1$~TeV \\cite{Cyburt:2002uv}.  Among the current chaotic inflationary scenarios in supersymmetric seesaw theories~\\cite{Murayama:1992ua,Hamaguchi:2001gw,Ellis:2003sq}, inflation driven by the scalar superpartner of the right-handed Majorana neutrino is one of the simplest and most economical ones. In this context, the heavy singlet neutrinos are naturally invoked to give masses to the light neutrinos through the seesaw mechanism~\\cite{seesaw}. Moreover, their superpartners - the sneutrino fields - can play the role of the inflaton. Also, if $CP$ is violated the sneutrino decays will create a lepton asymmetry, which is then converted into a baryon asymmetry by the electroweak sphalerons. This is indeed an appealing scenario, since cosmology and particle physics merge together to make predictions about the early universe and the low-energy physics that we test today.  There is, however, a drawback in the above-mentioned framework. In the usual chaotic inflation scenario based on standard cosmology, super-Planckian inflaton field values $\\sim 3 M_{P}$ are typically required to allow for a sufficiently long period of inflation (the so-called $\\eta$ problem). Thus one expects nonrenormalizable quantum corrections of the order of $\\mathcal{O}[(\\phi /M_{P})^{n}]$ (with $n > 4$) to destroy the flatness of the potential necessary for successful inflation. A possible way out of this situation is to consider, for instance, higher-dimensional cosmological models, where our four-dimensional world is viewed as a 3-brane embedded in a higher-dimensional bulk.  A remarkable feature of brane cosmology is the modification of the expansion rate of the universe $H$ before the nucleosynthesis era \\cite{Binetruy:1999ut}. While in standard cosmology the expansion rate scales with the energy density $\\rho$ as $H\\propto\\sqrt{\\rho}$, this dependence becomes $H\\propto\\rho$ at very high energies in brane cosmology. This behavior, which appears to be quite generic and not specific to Randall-Sundrum braneworld scenarios~\\cite{Randall:1999vf}, may have drastic consequences on early universe phenomena such as inflation and the generation of the baryon asymmetry. In particular, modifications to the Friedmann equation not only ease the conditions for slow-roll inflation but also enable the simplest chaotic inflation models to inflate at field values far below $M_{P},$ thus avoiding well-known difficulties with higher-order nonrenormalizable terms. Another important difference between standard and brane cosmologies is in the predictions for gravitino production. For a given value of the brane tension or, equivalently, of the 5D Planck mass, $M_{5}$, the gravitino abundance in the brane decreases as $T_{rh}$ increases. Therefore, in contrast to standard cosmology, BBN constraints in the brane scenario imply a lower (rather than an upper) bound on the reheating temperature.  The aim of this paper is combines the above ideas, i.e. chaotic inflation and direct leptogenesis through sneutrino decays in the braneworld context~\\cite{Dvali:1999gf}. More precisely, we study a simple scenario of chaotic brane inflation where, in a minimal supersymmetric seesaw model, the scalar superpartner of a heavy singlet Majorana neutrino drives inflation and, simultaneously, generates the required lepton asymmetry through its direct out-of-equilibrium decays after the inflationary era. This requires the reheating temperature $T_{rh}$ to be smaller than the sneutrino inflaton mass, $M_1\\,$. In this framework, there exists a direct connection between the brane inflationary era, the reheating of the universe, leptogenesis from sneutrino decays and the light neutrino properties, which allows us to strongly constrain the fundamental 5D Planck mass and the reheating temperature of the universe.  We shall not consider here the case where leptogenesis is purely thermal, i.e. when $T_{rh} > M_1$. In this case, any lepton asymmetry generated through the sneutrino inflaton decays is erased by thermal effects, and therefore, leptogenesis is driven by the out-of-equilibrium decays of the heavy singlet Majorana neutrinos and sneutrinos thermally created. ",
        "conclusions": "We now put together the various constraints we have derived sofar. In Figure \\ref{fig1}, we show the dependence of the 5D Planck mass $M_5$ on the effective leptogenesis phase $\\delta_L$, for two values of the gravitino mass, $m_{3/2}= 100$~GeV and 1~TeV. The lower bound on $M_5$ (solid lines) comes from the direct leptogenesis condition, i.e. $T_{rh} < M_1$, together with Eqs.~(\\ref{mnum}) and (\\ref{ntrhb}). The upper bound (dot-dashed lines) is obtained from Eqs.~(\\ref{b1t}) and (\\ref{trhomega}). The shaded area is the allowed region, which is clearly bigger for larger $m_{3/2}$. From this figure we see that the minimum allowed values for the effective leptogenesis phase are $\\sin \\delta_L \\gtrsim 0.17$ and $\\sin \\delta_L \\gtrsim 6.4 \\times 10^{-3}$, for $m_{3/2}= 100$~GeV and 1~TeV, respectively.  In Figure \\ref{fig2}, we plot the reheating temperature $T_{rh}$ as a function of $M_5$, including the bounds from gravitino production (dot-dashed) (cf. Eqs.(\\ref{b1t}) and (\\ref{trhomega})), the direct leptogenesis bound (solid) and the bound from Eq.~(\\ref{trhlow}), obtained from the WMAP result for $\\eta_B$ (dotted). The shaded area corresponds to the allowed region. From these figures we conclude that the allowed range for $M_5$ and $T_{rh}$ is \\begin{eqnarray} 3.6 &\\times& 10^{10}~\\mbox{GeV} \\lesssim M_5 \\lesssim 2.1 \\times 10^{11}~\\mbox{GeV}~, \\nonumber\\\\ 1.6 &\\times& 10^6~\\;\\mbox{GeV} \\lesssim T_{rh} \\lesssim 9.6 \\times 10^6\\;~\\mbox{GeV}~, \\end{eqnarray} if $m_{3/2}= 100$~GeV, while for $m_{3/2}= 1$~TeV we find \\begin{eqnarray} 3.6 &\\times& 10^{10}~\\mbox{GeV} \\lesssim M_5 \\lesssim 5.7 \\times 10^{12}~\\mbox{GeV}~,\\nonumber\\\\ 1.6 &\\times& 10^6\\;~\\mbox{GeV} \\lesssim T_{rh} \\lesssim 2.6 \\times 10^8\\;~\\mbox{GeV}~. \\end{eqnarray} Finally, the bounds on the lightest right-handed Majorana neutrino mass, $M_1$, are the same as the ones on the reheating temperature.  Braneworld cosmology is a rich subject. During the past few years there has been renewed activity and interest in this domain. Modifications to the expansion rate of the universe, as is typically the case in braneworld scenarios, can have profound implications for the processes that took place in the early universe. In this paper, we have considered the possibility that two of these phenomena, namely, chaotic inflation and the generation of the baryon asymmetry of the universe through leptogenesis, occurred during the nonconventional era in the brane. We have studied a minimal supersymmetric seesaw scenario where the lightest singlet sneutrino field not only plays the role of the inflaton but also produces a lepton asymmetry through its direct decays. Taking into account the BBN constraints on the gravitino production and the observed baryon asymmetry of the universe, we were able to strongly constrain the fundamental 5D Planck mass scale, and consequently, the lightest sneutrino mass, as well as the reheating temperature of the universe. The effective leptogenesis phase is also bounded in this framework."
    },
    "0402/astro-ph0402595.txt": {
        "abstract": "{For the nonthermal radio emission of the Galactic Center Arc in situ electron acceleration is imperative. The observed radio spectrum can be modeled by a transport equation for the relativistic electrons which includes particle acceleration by electric fields, momentum diffusion via scattering by magnetohydrodynamical turbulence and energy losses by synchrotron radiation. The accelerating electric fields can be regarded as a natural consequence of multiple reconnection events, caused by the interaction between a molecular cloud and the Arc region. The radio spectrum and even the recently detected 150 GHz emission, explicitely originating from the interaction regions of a molecular cloud with the magnetized Arc, can be explained in terms of quasi-monoenergetically distributed relativistic electrons with a typical energy of about 10 GeV accelerated in stochastically distributed magnetic reconnection zones. ",
        "introduction": "The Galactic Center (GC) Arc is a unique nonthermal radio continuum structure in the Galaxy. It is located at a projected distance of about 30 pc from the Sgr A complex and consists of a network of magnetic filaments filled with relativistic electrons running almost perfectly perpendicular to the Galactic plane (e.g. \\cite{Yus84}). The nonthermal nature of these filaments has been proven beyond any doubt by radio emission which exhibits a degree of linear polarization which is close to the intrinsic value at high frequencies of 60\\% \\cite{reich88, lesch92}. We emphasize the nonthermal character of the Arc since its radio spectrum is a very remarkable one. A decomposition of its radio spectrum between 843 MHz and 43 GHz by \\cite{reich88} shows a spectral index $\\alpha = + 0.3$ (we use the conventional relation that the observed flux $S_\\nu$ scales as $\\nu^{\\alpha}$). This finding of an increasing radio spectral index has been confirmed by high resolution VLA-observations \\cite{anan91}.  Such an inverted radiation spectrum with an index of +0.3 is expected from a quasi-monoenergetic electron distribution or an energy distribution with a well-defined low energy cutoff, respectively \\cite{lesch88}. Observations at high frequencies (between 32 GHz \\cite{lesch92} and 43 GHz \\cite{sof87b}) seem to indicate a spectral turn over, i.e. a fading of the Arc towards higher frequencies. However, Reich et al. (2000) detected an enhanced emission at 150 GHz slightly offset relative to the most intense vertical nonthermal filaments seen at lower frequencies. This emission originates from the apparent interacting areas of dense molecular material with the Arc. 150 GHz emission was also detected south of a molecular cloud, where the vertical filaments of the Arc cross a weak filamentary structure.  The spectrum of the emission is inverted relative to 43 GHz and compatible with an origin from quasi-monoenergetic electrons or an electron distribution with a low energy cutoff, but not with optically thin emission from cold dust. Reich et al. (2000) conclude \"that the coincidence of enhanced emission with regions of interacting molecular gas strongly suggests that high-energy electrons are accelerated in those places where the magnetic field is compressed\". They calculated a Lorentz factor of the electrons to be $\\gamma\\simeq 2\\times 10^4$ emitting synchrotron radiation at 150 GHz within a magnetic field of 1mG.  To summarize, in the Galactic center several areas appear to be filled with relativistic electrons whose energy must be distributed quasi-monoenergetically around a few GeV.  We note that even the very center of our Galaxy Sgr A$^{*}$ exhibits a radio to infrared spectrum which is in surprisingly good agreement with optically thin synchrotron emission of a quasi-monoenergetic electron distribution with a typical energy of 80 MeV, i.e. Lorentz factor of 160 \\cite{dusch94}.   In a former paper about particle acceleration in the Arc \\cite{lesch92} we considered the magnetohydrodynamical interaction of a molecular cloud with the Arc filaments. Especially we investigated the role of the moving gas cloud as a trigger mechanism for magnetic field amplification accompanied by dissipation of the magnetic energy via magnetic reconnection.  Magnetic reconnection takes place when magnetic field lines with antiparallel directions encounter.  Such a situation corresponds to the formation of an electric current sheet.  Typically astrophysical plasmas are ideal electric conductors, i.e. their electrical conductivity is very high. In such media the magnetic field is frozen into the motion of the conducting fluid. Any plasma velocity distortion onto the magnetic field lines is automatically related to an electric field which is perpendicular to the plasma velocity and the magnetic field. This electric field is described by the ideal Ohm's law \\be \\vec{E} + \\frac{1}{c} \\vec{v} \\times \\vec{B} = 0. \\ee In a plasma where the magnetic field is strongly distorted by shear flows, radial explosive flows or stochastic motions, it is unavoidable that field lines with antiparallel directions encounter. Consequently, the ideal conducting plasma switches into a nonideal medium with a finite, localized electrical conductivity, i.e. the ideal form of Ohm's law is violated. The strong spatial gradients of the magnetic field in such interaction zones represent an energetic \"crisis\" which is relaxed by partial dissipation of the magnetic energy via the formation of current sheets in the resistive medium (e.g. \\cite{pri00} and references therein). The nonideal Ohm's law is given by \\be \\vec{E} + \\frac{1}{c}\\vec{v} \\times \\vec{B} = \\vec{R} \\neq 0, \\ee where $\\vec{R}$ is some yet unspecified nonideal term, i.e. the plasma resistance.  In other words, magnetic reconnection corresponds to a magnetic field aligned electric field which is associated with a generalized electric potential \\cite{sch88, sch91} \\be V=-\\int{E_s \\; \\rm{ds}}=-\\int{R_s \\; \\rm{ds}} \\ee where the integral is evaluated along the magnetic field lines that penetrate the reconnection region. Of course, such a potential drop offers the possibility to accelerate particles very efficiently \\cite{sch91, black96, lesch97, lesch98, sch98, lit99, nod03}.  Instead of investigating the acceleration in one reconnection zone, we consider in this contribution the acceleration of relativistic electrons in numerous reconnection regions, i.e. current carrying filaments driven by the interaction of a molecular cloud with the magnetic field in the Galactic Center Arc \\cite{lesch92, ser94}.  This scenario arises quite naturally, since the necessary energy source is represented by moving molecular gas which encounters the poloidal magnetic fields in the Arc which is proven to be present by the observed very high polarization of the radio emission up to 60\\%.  Since we have many acceleration regions we apply a dynamical description of the energy distribution function of relativistic electrons based on systematic momentum gains by reconnection, losses by synchrotron losses and momentum diffusion due to scattering on magnetohydrodynamical turbulence, i.e. Alfv\\'en waves and magnetosonic waves. ",
        "conclusions": "In the filaments of the Galactic Center Arc electrons are accelerated to considerably high energies of about 10 GeV. This is obvious from the detected polarized radio emission exhibiting a very high degree of polarization (60\\%) and a rising spectrum with a spectral index of +0.3 up to 150 GHz. Such a spectral behavior is due to an energy distribution function of the radiating relativistic electrons which either has a low-energy cutoff or which is quasi mononenergetic. No obvious energy sources for particle acceleration are present in the neighborhood of the filaments. The very center of the galaxy, SgrA${^*}$ is a source for monoenergetic relativistic particles but on significant lower energies of about 50 MeV \\cite{dusch94}. However, the magnetic filaments interact with the plasma of a molecular cloud which moves with velocities of some 10 km ${\\rm s}^{-1}$ relative to the magnetic field.  This interaction can be interpreted in terms of the induced Lorentz force inducing a convective electric field which in the first place is oriented perpendicular to the magnetic field and the local plasma velocity of a molecular cloud.  The motion of the plasma locally distorts the magnetic field necessarily in such a way that antiparallel directed field lines encounter. Provided that some violation of ideal Ohm's law occurs, e.g. due to microturbulence, magnetic reconnection will occur. The magnetic field energy will be partly converted to particle energization. In the reconnection regions a magnetic field-aligned electric field component forms with a magnitude of some fraction of the convective electric field.  This parallel electric field represents a perfect candidate for efficient particle acceleration. The required electric field for the electron energization are considerably weaker than the magnetic field which is in accordance with the observational fact that the cloud plasma moves with velocities much lower than the speed of light.  Since reconnection is a localized phenomenon which depends on the local properties of the plasma, we consider a scenario characterized by multiple reconnection sheets in which particles are accelerated and scattered by magnetohydrodynamical fluctuations.  Such scattering leads to diffusion in energy space.  Moreover, we take into account the energy losses by the observed synchrotron radiation.  By means of a relatively simply transport equation for the distribution function of the radiating relativistic electrons we could show that an injected monoenergetic distribution function with 75 MeV (coming from the Galactic center) evolves into a quasi-monoenergetic distribution function at 10 GeV.  The resulting radiation spectra increase with frequency up to some hundred GHz according to the observations, i.e. $S_\\ind{\\nu} \\sim \\nu^{1/3}$.  A major assumption of our study is the complete isotropy of the distribution function which is based on calculations by \\cite{acha90}, who show that the isotropisation time $t_\\ind{iso} = c/v_\\ind{A}(B/\\delta B)1/\\Omega$, where $\\Omega$ is the gyro frequency of the protons. In this application $t_\\ind{iso}$ is about $4\\cdot10^6$ s which is much shorter than the studied time scale of the evolution of the distribution function. Some non-isotropic particle population would lead to an overlap of a second radiation component in a specific energy range that we can not handle in our model. However, the present observations seem to give no hint to such a component.   \\begin{figure} \\subfigure[Temporal evolution of $N(\\gamma)$. The solid, dotted and dashed lines show the injected particle spectrum and the evolution after $t=1\\,t_0$ and $t = 2\\,t_0$, respectively.]{ \\includegraphics[angle=0,width=0.55\\linewidth,keepaspectratio]{0489fig1.eps} } \\subfigure[The corresponding synchrotron spectrum $I_{\\nu}$ ]{ \\includegraphics[angle=0,width=0.55\\linewidth,keepaspectratio]{0489fig2.eps} } \\subfigure[Temporal evolution of $N(\\gamma)$. The solid and dotted lines show the evolution after $t=10\\,t_0$ and $t=20\\,t_0$, respectively.]{ \\includegraphics[angle=0,width=0.55\\linewidth,keepaspectratio]{0489fig3.eps} } \\subfigure[The corresponding synchrotron spectrum $I_{\\nu}$]{ \\includegraphics[angle=0,width=0.55\\linewidth,keepaspectratio]{0489fig4.eps} } \\caption[] {Temporal evolution of $N(\\gamma)$ calculated from equation 6. The values of the physical parameters, which we have used for our calculations are: $\\zeta = 1.5\\cdot 10^{-22}$ dyne, $k = 1.4\\cdot 10^3\\;\\rm{g^{-1}\\,cm^{-1}}$ and $D = 5.4\\cdot 10^{-18} \\rm{s^{-1}}$. To evaluate the synchrotron emission spectrum (see equation 12) we have used a magnetic field strength of $B = 3$ mG.} \\label{} \\end{figure}   \\begin{figure} \\subfigure[Temporal evolution of $N(\\gamma)$. The solid and dotted lines show the evolution after $t=30\\,t_0$ and $t=50\\,t_0$, respectively.]{ \\includegraphics[angle=0,width=0.55\\linewidth,keepaspectratio]{0489fig5.eps} } \\subfigure[The corresponding synchrotron spectrum $I_{\\nu}$]{ \\includegraphics[angle=0,width=0.55\\linewidth,keepaspectratio]{0489fig6.eps} } \\subfigure[Temporal evolution of $N(\\gamma)$. The solid, dotted and dashed lines show the evolution after $t=80\\,t_0$, $t=100\\,t_0$ and $t=150\\,t_0$, respectively.]{ \\includegraphics[angle=0,width=0.55\\linewidth,keepaspectratio]{0489fig7.eps} } \\subfigure[The corresponding synchrotron spectrum $I_{\\nu}$]{ \\includegraphics[angle=0,width=0.55\\linewidth,keepaspectratio]{0489fig8.eps} } \\caption[] {Temporal evolution of $N(\\gamma)$ towards a quasi-monoenergetic distribution function (a), (c). The resulting synchrotron emission spectrum (d) shows an inverted spectrum with a spectral index of $\\alpha = 1/3$ ($I_\\nu \\sim \\nu^{\\alpha}$) and a cutoff-frequency of some hundred GHz after $1.5\\cdot10^{11}$ s.} \\label{fig2} \\end{figure}"
    },
    "0402/astro-ph0402163_arXiv.txt": {
        "abstract": "It has been suggested that the peculiar properties of the luminous Type Ic supernova SN\\,1998bw and its low-energy gamma-ray burst GRB\\,980425 may be understood if they originated in a standard gamma-ray burst explosion viewed far from the axis of the relativistic jet.  In this scenario, strong radio emission is predicted from the jet on a timescale 1 to 10 years after the explosion as it decelerates and spreads into our line of sight.  To test this hypothesis we have carried out late-time radio observations of SN\\,1998bw at $t=5.6$ years, yielding upper limits which are consistent with the continued fading of the supernova. We find these limits to be consistent with an off-axis jet only if the progenitor mass loss rate is $\\dot{M}\\lesssim 4 \\times 10^{-7}$ M$_\\odot$ yr$^{-1}$ (for a wind velocity $v_w=1000$ km s$^{-1}$) or the fraction of the shock energy in magnetic fields is $\\epsilon_B \\lesssim 10^{-3}$.  These values are low relative to those inferred for cosmological GRBs.  We combine the SN\\,1998bw measurements with existing observations for a sample of 15 local Type Ibc supernovae to estimate that at most 6\\% produce collimated, relativistic outflows. ",
        "introduction": "GRB\\,980425 was the first gamma-ray burst (GRB) to be identified with a supernova (SN), SN\\,1998bw \\citep{gvv+98,paa+00}.  However, this association was not universally accepted since the spatial and temporal coincidence of these events could only be approximated to $\\pm 8^\\prime$ and $\\pm$2 d, respectively. Recently, the detection of spectroscopic features in the light curve of GRB\\,030329, similar to those seen in SN\\,1998bw \\citep{smg+03,hsm+03}, has strengthened the hypothesis of the SN\\,1998bw/GRB\\,980425 association. These results have given a powerful impetus to the ``collapsar'' model \\citep{mw99,zwm03} in which a Wolf-Rayet progenitor undergoes core collapse, producing a rapidly rotating black hole surrounded by an accretion disk which injects energy into the system and thus acts as a ``central engine''. The energy extracted from this system supports a quasi-spherical Type Ibc SN explosion and drives collimated jets through the stellar rotation axis which produce the prompt gamma-ray and afterglow emission (see review by \\citealt{zm03}).  Despite this progress, some unresolved issues remain.  Foremost among these is understanding the connection between local SNe Ibc and cosmological GRBs.  Estimates of the fraction of Type Ibc SNe that produce a GRB range from $10^{-5}$ \\citep{psf03}, to 0.5\\% \\citep{fks+01}, and even approaching 100\\% \\citep{ldg03}.  This uncertainty in the relative rates of Ibc SNe and GRBs is due to different assumptions about the geometry and energetics of GRBs.  As the only local SN observed in association with a GRB, SN\\,1998bw is the key to our understanding.  At a distance of 38 Mpc, this Type Ic SN was unusually luminous at optical and radio wavelengths \\citep{kfw+98,gvv+98}.  The broad absorption lines seen in early spectra of SN\\,1998bw implied expansion velocities of $\\ge 30,000\\rm~km s^{-1}$ and (isotropic) kinetic energies of $\\sim 3 \\times 10^{52}$ erg \\citep{imn+98,wes99,tmm01}. Likewise, the bright, early-peaked radio emission implied a significant amount of energy ($\\sim 10^{50}$ erg) coupled to mildly ($\\Gamma\\approx 2$) relativistic ejecta \\citep{kfw+98}. Detailed modeling by \\citet{lc99} confirmed this result and also showed the need for a second energy injection episode, indicating the presence of a central engine, similar to the model inferred for GRBs.  However, in contrast to the large energies inferred from the optical and radio emission from SN\\,1998bw, GRB\\,980425 was a sub-energetic gamma-ray burst. The prompt emission had an (isotropic) energy release of only $8\\times 10^{47}$ erg \\citep{paa+00} -- 4 to 6 orders of magnitude below typical GRBs \\citep{fks+01}.  One popular explanation posits that SN\\,1998bw/GRB980425 was a typical GRB viewed away from the jet axis \\citep{nak98,cen98,el99,wes99,sal01,yyn03}. This hypothesis can been described by two different scenarios based on the ratio of the off-axis angle, $\\theta_{\\rm oa}$, to the opening angle of the GRB jet, $\\theta_j$. In Case 1, $\\theta_{\\rm oa} \\sim 3\\theta_j$ so the jet emission is detected at early time.  Here, the $\\gamma$-rays originate from the edge of the relativistic jet, causing the inferred (isotropic) energy to be suppressed \\citep{gpk+02}. In Case 2, $\\theta_{\\rm oa}$ is large ($\\gg 10$ $\\theta_j$) so the jet emission is only detectable at late time ($t\\sim 1-10$ years) when the jet has reached spherical symmetry. In this case, the prompt $\\gamma$-ray emission may be due to Compton scattering of photons into our line-of-sight \\citep{wax03}.  \\citet{gpk+02} have investigated Case 1, finding the optical and $\\gamma$-ray emission to be consistent with an off-axis angle of $\\theta_{\\rm oa}\\approx 4^o$ (for $\\Gamma \\approx 100$).  This scenario, however, may have difficulty explaining the X-ray and radio evolution. \\citet{wax03} has recently investigated Case 2, predicting a late-time rise in the observed luminosity which is most easily detectable at radio frequencies \\citep{pac01,tot03,gl03}.  The possibility of Case 2 has given rise to unification models which imply some fraction of local Ibc SNe can also be described as GRBs viewed off-axis \\citep{wax03,ldg03}.  While the off-axis jet model provides a convenient framework in which to unite the GRB and SN phenomena, confirmation of the model requires observational evidence for a GRB jet within a local SN.  In this paper we carry out late-time radio observations of SN\\,1998bw in an effort to detect the putative off-axis jet from GRB\\,980425.  We combine this measurement with existing observations for a sample of nearby Type Ibc SNe to place constraints on the parameters of the off-axis jet and the fraction of local Ibc SNe which could be associated with relativistic jets similar to those seen in cosmological GRBs. ",
        "conclusions": "Late-time observations of SN\\,1998bw ($t\\simeq$5.6 yrs) have allowed us to test the hypothesis that GRB\\,980425 was a standard GRB viewed far away from the jet axis. Our measured upper limits at 1384 and 2368 MHz are consistent with the continued power-law decay of the SN emission.  These limits imply an off-axis jet is only plausible if the normalized mass loss rate of the progenitor star is $\\dot{M}_{-5}/v_{w,3} \\le 0.04$ (for $\\epsilon_B \\ge 0.1$). This is $\\sim 20-200$ times smaller than the observed mass loss rates for local Wolf-Rayet stars \\citep{cgv03} and is below the range typically observed in GRBs. Larger mass loss rates are possible but only if the energy fraction in magnetic fields is low ({\\em i.e.}, $\\epsilon_B \\lesssim 10^{-3}$).  Even tighter constraints are derived for the off-axis jet model when we examine a larger sample of local Ibc SNe. The low luminosity limits derived for this sample require values of $\\dot{M}_{-5}/v_{w,3}\\approx 0.01-0.1$ or  $\\epsilon_B \\lesssim  10^{-3}$ which are below values for typical GRBs. The absence of any late-time radio emission can therefore be used to put a limit on the fraction of core-collapse SNe that produce collimated, relativistic outflows. Our results imply that off-axis jets from nearby SNe are rare ($\\lesssim $6\\%) with the possible exception that the radio emission from SN\\,2001em is due to a GRB jet.  This conclusion complements the findings of \\citet{bkf+03} who constrained the GRB/SN fraction through a radio survey of local Ibc SNe at early time. \\citet{bkf+03} used early, bright radio emission as a proxy for relativistic ejecta, as in the case for SN\\,1998bw. After studying 33 local SNe with detection limits $10^3$ times fainter than SN\\,1998bw, \\citet{bkf+03} found no evidence for relativistic ejecta in any of the SNe observed, thereby constraining the GRB/SN fraction to $\\lesssim 3\\%$.  Taken together, these results support a view that SN\\,1998bw was a rare and unusually energetic SN -- distinct from local SNe and GRBs. In this scenario, the characteristics of SN\\,1998bw/GRB\\,980425 are not dictated by the observer's viewing angle, but rather by the properties of its central engine. SN\\,1998bw was an engine-driven explosion \\citep{lc99}, in which 99.5\\% of the kinetic energy ($\\sim 10^{50}$ erg) was coupled to mildly ($\\Gamma\\approx 2$) relativistic ejecta \\citep{kfw+98}, while a mere 0.5\\% was detected in the ultra-relativistic ($\\Gamma\\approx 100$) flow.  In contrast, GRBs couple most of their energy to relativistic $\\gamma$-rays. The observed diversity of cosmic explosions (SNe, X-ray flashes, and GRBs) may therefore be explained with a standard energy yield, but with a varying fraction of that energy given to relativistic ejecta \\citep{bkp+03}.  We thank Edo Berger, Sarah Yost and Eli Waxman for helpful discussions.  AMS is supported by the NSFGRFP."
    },
    "0402/astro-ph0402480_arXiv.txt": {
        "abstract": "{  In 2001, using a large spectroscopic dataset from an extensive monitoring campaign, we discovered that the peculiar Of star \\hd\\ displayed extreme line variations. This strange behaviour could be attributed to a variety of models, and an investigation of the high energy properties of \\hd\\ was needed to test the predictions from these models. Our dedicated \\x\\ observation of \\hd\\ shows that its spectrum is well represented by a two temperature thermal plasma model with $kT_1\\sim0.2$~keV and $kT_2\\sim1.4$~keV. In addition, we find that the star does not display any significant short-term changes during the \\x\\ exposure. Compared to previous $Einstein$ and \\ro\\ detections, it also appears that \\hd\\ does not present long-term flux variations either. While the line variations continue to modify \\hd's spectrum in the optical domain, the X-ray emission of the star appears thus surprisingly stable: no simple model is for the moment able to explain such an unexpected behaviour.  Thanks to its high sensitivity, the \\x\\ observatory has also enabled the serendipitous discovery of 57 new X-ray sources in the field of \\hd. Their properties are also discussed in this paper. ",
        "introduction": "Some O-type stars are still challenging astronomers many decades after their discovery. \\hd\\ is a good example of such an object. In the past, this star has been the target of several investigations, but apparently, none could arrive at a consensus on the exact nature of this peculiar star. Consequently, various theories have successively been proposed: \\hd\\ has been classified as a short-term binary (having even survived a supernova event according to some authors); it has been proposed to be a single star experiencing wind variability and/or harbouring a disc and jets; and finally it was suggested to be a long-term binary (see Naz\\'e et al. 2001, hereafter Paper I, and references therein). In Paper I, we presented a 30 year campaign of optical spectroscopy dedicated to \\hd. We showed that the behaviour of \\hd\\ was not that of a classical short- or long-term SB1 binary. The star also did not present any short-term variations. But our extensive campaign clearly indicated the peculiar characteristics of this star: tremendous line variations on the timescale of decades. The hydrogen and \\hei\\ lines change from strong P Cygni profiles to simple absorptions, and a few other emission lines apparently follow the same behaviour. Comparing with all data available, we noted that these variations were recurrent, on a timescale of a few decades. Such a continuous decline of \\hi\\ and \\hei\\ lines was recently discovered by Walborn et al. (\\cite{wal03}) in another Of?p star, HD\\,191612. The line variations of this star are strikingly similar to \\hd, but occur on a shorter timescale. In this context, we note that HD\\,191612 has now returned to a high emission state: a spectrum of this star covering the \\ha\\ line, taken in October 2003 at the Observatoire de Haute-Provence, is almost identical to that of August 1997 presented by Walborn et al. (\\cite{wal03}).  The nature of these Of?p stars is still unknown. The most popular models to explain their peculiar behaviour involve the presence of a compact companion on an elliptical orbit, or require the star to be a single rapid rotator (see e.g. the review in Walborn et al. \\cite{wal03}). It was also suggested that these stars could be transition objects, explaining their small number (only 3 are known in our Galaxy). In this context, X-ray observations represent an important opportunity to better understand these objects, since stars reveal at high energies the most exotic processes taking place in their vicinity. X-ray data are especially well suited to test the possibility of accretion processes linked to the presence of a compact object or a colliding wind interaction in a binary. We thus decided to observe \\hd\\ with the \\x\\ satellite, in the hope that its high sensitivity could provide us with definitive answers on the nature of \\hd.  In this paper, we describe in Sect.~2 the observations used in this study. The optical and X-ray data of \\hd\\ will then be successively analysed in Sect.~3. The remarkable sensitivity of \\x\\ also enabled the serendipitous discovery of many fainter sources during our observation. These sources detected in the field of \\hd, their possible counterparts, their hardness ratios (HRs), their variability, and their spectral characteristics will be discussed in Sect.~4. Finally, we will conclude in Sect.~5. ",
        "conclusions": "We have obtained an \\x\\ observation of \\hd\\ and its surroundings. The peculiar star \\hd\\ was found to present a two-temperature spectrum, and did not show any significant short-term variations during the exposure. These observations are also compatible with a stable X-ray emission since 1979. In parallel, we have continued our extensive optical monitoring of the star, and discovered that \\hd\\ continues to present dramatic line variations in the optical domain. The lack of significant changes in the X-ray emission compared to the optical data is a puzzle. In such a context, no simple model is for the moment capable of explaining the overall behaviour of one of the most peculiar O stars of the sky: any variation of the physical properties (magnetic field, mass-loss rate) of \\hd\\ is expected to have an impact on the whole wind, not only on the formation region of the optical emission but also on that of the X-rays. Long-term monitoring of \\hd, preferentially in a multi-wavelength campaign, is thus necessary to eventually understand this star.  In addition, 57 new X-ray sources were also discovered in the field of \\hd, and we present here their characteristics (count rate, HRs, plus lightcurve analysis and spectral fits for the brightest ones). Only two correspond to rather bright stars, both of them being late-type stars with coronal X-ray emission. With large HRs and/or large absorbing columns and power law fits, several sources might be background objects. On the other hand, none of the B stars of the field was convincingly detected in X-rays."
    },
    "0402/astro-ph0402449_arXiv.txt": {
        "abstract": "Sunyaev-Zeldovich (SZ) effect from a cosmological distribution of clusters carry  information on the underlying cosmology as well as the cluster gas physics. In order to study either cosmology or clusters one needs to break the degeneracies  between the two. We present a toy model showing how complementary informations from SZ power spectrum  and the SZ flux counts, both obtained from upcoming SZ cluster surveys,  can be used to mitigate the strong cosmological influence (especially that of $\\sigma_8$) on the SZ fluctuations.  Once the strong dependence of the cluster SZ power spectrum on $\\sigma_8$ is diluted, the cluster power spectrum can be used as a tool in studying cluster gas  structure and evolution. The method relies on the ability to write the Poisson contribution to the SZ power spectrum in terms  the observed SZ flux counts.  We test the toy model by applying the idea to simulations of SZ surveys. ",
        "introduction": "In the last decade, X-Ray observations of galaxy clusters have been used extensively to determine  the cosmological matter density parameter and the amplitude of density fluctuations (Henry 1997, Bahcall \\& Fan 1998,  Viana \\& Liddle 1999). In near future, advent of new mm and sub-mm high-sensitivity experiments will open  a new window for studies of galaxy clusters through the SZ    effect surveys.  These surveys  (for example SPT\\footnote{http://astro.uchicago.edu/spt/}, ACT\\footnote{http://www.hep.upenn.edu/$\\sim$angelica/act/act.html}, APEX-SZ\\footnote{http://bolo.berkeley.edu/apexsz/}) would be capable,  through wide area coverage and high sensitivity, of detecting  thousands of galaxy clusters up to  redshift $z > 1$. The cosmological possibilities of such a large data sets are  enormous and will allow to carry out independent estimations of the  cosmological parameters (Diego et al. 2002, Levine et al. 2002, Weller 2002, Majumdar  \\& Mohr 2003a,2003b, Hu \\& Kravtsov 2003, Lima \\& Hu 2004) which could be compared with those  obtained from  other observations (CMB, SNIa, Ly-$\\alpha$ forest, etc).  Specifically, the cluster abundance with redshift, $dN\\over  dz$,  would provide constraints on the dark energy equation  of state parameter $w\\equiv p/\\rho$ (Wang \\& Steinhardt 1998).  Since SZ effect is redshift independent, to get the redshifts of the clusters detected in these SZ surveys, one would need to have followup observations in the optical/IR bands. In the absence of redshifts, a large yield SZ survey have cluster number counts as a function  of SZ flux, $\\mathcal{N}(S)$ would be available. This is also sensitive to cosmological parameters (Barbosa et al. 1996, 1998, Bartlett 2000) but weaker in comparison to the redshift abundance. It is interesting then to explore alternative methods which exploit  the new SZ data sets without any need of $z$ information. This is the main essence of the paper.  In the absence of any redshift information one would also be able to estimate  the SZ power spectrum from the temperature maps in addition to the flux counts. The statistics of SZ maps have also been well studied, both analytically (e.g., Cooray 2000, Zhang \\& Pen 2001, Komatsu \\& Seljak 2002) and through  numerical simulations (e.g., Refregier et al. 2000, Seljak et al. 2001, Springel et al. 2001, da Silva et al. 2001, Zhang, Pen \\& Wang 2002). Comparison made between the two approaches (Refregier \\& Tessier 2002, but also see Zhang et al. 2004)  show that a simple halo model (Cooray \\& Sheth 2002) description of the clusters reproduces results from detailed N-body simulations. Some differences remain, especially at very high $\\ell$'s where presumably the effect from substructures creep in. All these studies come to  similar conclusions about the potential of using the cluster SZ power spectrum to constrain the background cosmology, especially $\\sigma_8$ (Bond et al. 2002). Apart from cosmology, it has also been shown that with some knowledge of cosmological parameters, the SZ power spectrum can be used as a powerful probe of cluster gas physics (Majumdar 2001, Holder \\& Carlstrom 2001, Sadeh \\& Rephaeli 2004a).  It is clear from all these studies that the possibilities of doing science with the SZ power spectrum are enormous.   It is, however, well known that cosmological studies based on cluster  data will show degeneracies in the cosmological parameters and  cluster scaling relations. One of the best known is the degeneracy between  $\\sigma _8 - \\Omega _m$ which can only be weakened if one can measure the  evolution of the cluster population as a function of redshift.  Moreover, imperfect knowledge of cluster structure and evolution can weaken cosmological constraints obtained from these surveys. In the presence of any such evolution of gas physics mimicking cosmological evolution of the cluster number density, one would have to resort to other options  (Majumdar \\& Mohr 2003b, Lima \\& Hu 2004) to tighten the constraints. In the absence of  any redshift information, when dealing with flux counts or power spectrum, one encounters similar degeneracies  between cosmology and cluster physics. One can try to break these degeneracies by directly adding constraints from flux counts to those from power spectrum resorting to time intensive joint-fits to the two data sets. This has the advantage in that it not only gives us the constraints on the different parameters but also the correlations between them. However, it is instructive to  look whether at all one can disentangle cosmology from cluster physics before resorting to more time consuming techniques. It is in this `exploratory' sense that we try to look at the possibility of mitigating the cosmological influence by combining flux counts to SZ power spectrum such that it becomes possible to probe gas physics. The intention is not to estimate precise constraints achievable from upcoming SZ surveys but to look at the feasibility of studying cluster gas physics with these surveys. This approach towards studying cluster properties would be worthwhile since the present uncertainties in determining cluster properties from targetted observations remains large (like upto many tens of percents for any redshift evolution of the cluster scaling relations).  In this work, we propose to combine the cluster number counts as a function  of SZ effect flux, $\\mathcal{N}(S)$, and the power spectrum of the sample, C$_{\\ell}$,  in such a way such as to dilute the dependency of the SZ power spectrum on cosmological parameters.  Each depends on the background cosmology and gas physics in slightly different way than the other (for example, the SZ power spectrum depends on the detailed distribution of the gas in a cluster whereas the SZ flux only depends on the total thermal content of the cluster). Although both $\\mathcal{N}(S)$ and $C_\\ell$ depend on the total fluxes of the clusters, the latter also depends on the shape of the clusters, the size of cluster being inversely proportional to its angular distance. Thus the cosmology-gas physics degeneracies are different for the two quantities. Let us note, once again, at this point that both  $\\mathcal{N}(S)$ and $C_\\ell$ are quantities where the implicit redshift dependence is integrated over unlike that from $\\frac{dN}{dz}$. Moreover, single SZ survey would give us both the quantities and hence no extrapolation is needed between observations at different wavelengths.   The rest of the paper is organized as follows: in \\S \\ref{sec_formalism} we lay down the basic formalism and then in \\S  \\ref{sec_modelling} we describe the toy model of the clusters. The standard SZ power spectrum and its dependences on different parameters are shown in \\S \\ref{sec_powspec}. We present the simulations in \\S \\ref{sec_simulations} and in \\S \\ref{sec_clusterphys} we show the use of hybrid power spectrum as a tool for probing gas physics and discuss the systematics. Finally, in  \\S \\ref{sec_summary} we discuss and summarize the method and results. ",
        "conclusions": "\\label{sec_summary}  In the previous sections we have demonstrated that cosmology as well as cluster gas structure  and evolution simultaneously shapes the thermal SZ power spectrum from clusters of galaxies.  The power spectrum depends strongly on certain cosmological parameters (like $\\sigma_8$) and  on any evolution, if present, thus leading to a cosmology-gas physics degeneracy. One can take priors from external observations, for example fix the cosmological model from primary CMB,  and use the power spectrum to probe gas physics. However, it is possible to combine complimentary  information from the same survey itself to mitigate the influence of cosmology in order to study  gas physics.  To do so, we have formulated a different way of looking at the SZ power spectrum by constructing  the hybrid power spectrum. Given that any future large yield SZ survey would also detect many  thousands of clusters, we would have a flux counts of the objects detected in the survey. At the  same time, temperature fluctuations from all the survey pixels would be used to construct the  total SZ power spectrum. We have shown that one can rewrite the Poisson part of the SZ power  spectrum using the information available from the SZ flux counts. Since the Poisson power spectrum  dominates, in general, over the clustering power spectrum, the resultant hybrid power spectrum  represents the total SZ power spectrum sufficiently well up to $\\ell \\sim 2000$. We also used results from numerical simulations to test and verify our analytical results. Since both the  SZ power spectrum and the flux counts represent the same underlying cosmology, the hybrid power  spectrum is, in essence, normalized to the background cosmology. Once the cosmology-gas physics  degeneracy is diluted, we are left with modeling of the gas physics to understand the observed  total SZ power spectrum at higher multipoles.  We have used a simple model of the cluster gas structure and evolution as well analytical results  for the mass function in our analytical model. This is sufficient for the present work which is {\\it exploratory} in nature. Needless to say, practical application of the method would need better  modeling of the complex cluster structure using N-body hydro simulations or with more complicated  analytical models (suitable to fit simulation results). Cosmological simulations of large yield of  clusters mimicking upcoming surveys are still naive in the sense that they still need some  sort of phenomenological approach to model the gas physics (especially non-adibaticity) and has to make simplifying assumptions  as to the cluster structure. At the same time simulating any non-standard cluster evolution is  non-trivial and either done too simplistically or is put in by-hand in the simulations. It is this  very cluster structure and evolution that we are able to probe by studying the SZ power spectrum  at high $\\ell$-values.  Note, that in figure (\\ref{fig_Clsimul}), at multipoles  ($\\ell > 2000$) the hybrid power spectrum defers from the actual power spectrum. This, of  course, should be the case since the cluster structure (c.f. equations (\\ref{eqn_TM}) \\& (\\ref{eqn_RM})) used in constructing $C^H_l$ is very different from the gas physics used in the SZ N-body simulations.  However, it must  be kept in mind that a part of the difference may arise from increase in power  at high $\\ell$'s due to presence of substructures in the clusters which are naturally captured  in any N-body simulations but not modeled analytically.  The ability to construct $C^H_l$ depends crucially on our capability to get the cluster flux counts. The main source of systematic error in the construction of the hybrid power spectrum  comes from our partial knowledge of the ${\\mathcal{N}}(S)$ curve. This partial  knowledge may be due to the lack of enough statistics or a poor sensitivity  in the instrument. It is also important to understand that processing of the raw data can introduce a selection function in the cluster catalog. Usually, one should fit the ${\\mathcal{N}}(S)$ curve only using those  data points for which the selection function is well understood and then extrapolate  the fit to lower (and/or higher) fluxes. This extrapolation will have little effect in the constructed hybrid power spectrum if the survey is large enough such that it contains enough bright clusters and if its sensitivity is good enough to recover clusters with fluxes as low as a few mJy (say, at 145 GHz). Future SZ cluster surveys (like ACT/SPT/APEX-SZ) satisfy both these conditions such that the hybrid power spectrum can be a useful tool in these surveys.   The main conclusion of this work can be summarized by comparing figures (\\ref{fig_hybridcosmo} \\& \\ref{fig_hybridphys}) with  figures (\\ref{fig_totalcosmo} \\& \\ref{fig_totalphys}). By using the hybrid power spectrum instead of the standard power spectrum, one can dilute the dependency of the power spectrum on the cosmological model and concentrate on studies  of the cluster physics. The precision  in the determination of the cosmological parameters has increased dramatically  in recent times. However, much uncertainties remain in studies of  cluster physics. Better handle on cluster structure and evolution becomes even more challenging when one takes in to account the cosmology-gas physics degeneracy in any such studies.  The hybrid power spectrum helps to soften this problem by diluting much of this degeneracy.   The prospect of probing gas physics using SZ power spectrum would depend on our ability to measure  the SZ power spectrum as very high $\\ell$-values. This sub arc-min scale has the distinct advantage  of the primary CMB contribution being negligible. However one has to be careful in eliminating other  possible sources of secondary CMB anisotropies that may introduce further systematic uncertainties. In multifrequency experiments, an appropriate SZ reconstruction algorithm should be  able to recover the true SZ power spectrum up to the resolution limit  of the experiment  Finally, all future high sensitivity SZ observations have to tackle the noise introduced by unresolved  point sources. The spectral dependence of the thermal SZ effect would be helpful in eliminating many  of such systematics. Although not studied in this work, another interesting application of the hybrid power  spectrum is to perform consistency checks on the excess in power in single frequency  CMB experiments.  By this we mean that future CMB experiments will observe an excess in the power spectrum  of CMB fluctuations at $\\ell > 2000$. This excess will be due in part to galaxy clusters,  and in part to non-removed point sources (see figure 2 in White \\& Majumdar 2003). Now, at $\\ell \\sim 2000$ the hybrid power spectrum is a decent representation of the total SZ power spectrum without worrying much about cluster physics. If there is an estimation of the $\\mathcal{N}(S)$ curve for the survey, then one can estimate SZ power due to clusters at $\\ell \\sim 2000$.  Any excess power can be interpreted as due to unresolved point sources.  In spite the stringent requirements needed to make SZ observations from the upcoming surveys probe of  cluster physics, the rewards reaped would undoubtedly be great. In addition to new high  resolution targeted observations, the statistical study of the SZ sky from cosmological distribution  of clusters may provide the next leap in the study of the structure and evolution of galaxy clusters."
    },
    "0402/astro-ph0402213_arXiv.txt": {
        "abstract": "Fast variability studies of accreting black holes in the Galaxy offer us a unique opportunity to measure the spins of black holes and test the strong-field behavior of general relativity. In this review, I summarize the arguments often used in attempts of measuring the spins of black holes, concentrating on their theoretical foundations. I also argue that X-ray studies of accreting black holes will be able to provide in the future strong constraints on deviations from general relativity in the strong-field regime. ",
        "introduction": "Astrophysical black holes in general relativity are characterized by two quantities, their masses and spins, which determine uniquely the properties of their gravitational fields. As a result, both can, in principle, be measured by experiments involving test-particle orbits in their exterior spacetimes.  This approach has been very successful in measuring black-hole masses both in galactic systems (McClintock \\& Remillard 2003) and in the centers of galaxies (Sch{\\\" o}del et al.\\ 2002). However, the imprint of the spins of black holes on their spacetimes, i.e., the dragging of inertial frames, is very weak at the large distances from the horizons, where the observed orbits reside. It is expected that the detection of gravitational waves from close, inspiraling compact objects with LIGO and LISA will allow for a complete mapping of the spacetimes of the objects and hence for the measurement of black-hole spins (see, e.g., Hughes 2003). However, only a small fraction of mostly supermassive black holes exist in the near universe in configurations that will allow such studies.  Most of the black holes we observe today are visible because they accrete matter from their companions or the surrounding medium. The intense X-ray radiation we detect is generated in a region only a few Schwarzschild radii around the black-hole horizons. As a result, these X-ray photons carry with them the signatures of the strong gravitational fields in which they are produced and hence information regarding the masses and spins of the black holes.  In this review, I discuss the potential of measuring black-hole spins and confirming the predictions of general relativity, using their rapid-variability properties. In particular, I concentrate on the various methods of inferring black-hole spins that are based on the observations of constant-frequency quasi-periodic oscillations from galactic black-hole binaries (QPOs; see Remillard, this volume). The frequencies of these QPOs depend very weakly on the observed X-ray flux and, for this reason, it is believed that they are determined mostly by gravity and not by the hydrodynamic properties of the accretion flows, such as their temperatures and densities. ",
        "conclusions": "Studies of black-hole variability probe the strongest field regime of gravity that is possible for observers outside the horizon of the black hole. They offer the potential of proving the existence of black holes in the universe, measuring their spins, and testing gravity theories in regimes that is unattainable by local experiments and can complement cosmological probes.  \\begin{figure} \\includegraphics[height=.3\\textheight]{psaltisd_f7.ps} \\caption{The solid lines are contours of constant lifetime for black-holes of stellar mass in theories with a large extra dimension, as a function of the size of the extra dimension. Table-top experiments provide a bound of $L\\le 0.2$~mm. The dot represents the upper bound on the size of the large extra dimension imposed by inferring an age of at least 240 Myr for the $6.8\\pm 0.4 M_\\odot$ black hole in XTE~J1118$+$480 (after Emparan et al.\\ 2002).} \\end{figure}    \\begin{theacknowledgments} It is a pleasure to thank Tomaso Belloni, Chi-Kwan Chan, Simon DeDeo, Mike Nowak, Feryal \\\"Ozel, Martin Pessah, and Michiel van der Klis for all the discussions over the last several years that have helped me appreciate the various ways of using compact object variability in measuring black hole spins and in testing strong-field gravity. \\end{theacknowledgments}"
    },
    "0402/astro-ph0402025_arXiv.txt": {
        "abstract": "We report here the first polarimetric measurements of the pulsars in the J0737--3039 binary neutron star system using the Green Bank Telescope. We conclude both that the primary star (\\pAn) has a wide hollow cone of emission, which is an expected characteristic of the relatively open magnetosphere given its short spin period, and that \\A has a small angle between its spin and magnetic dipole axes, $4\\pm 3$ degrees. This near alignment of axes suggests that \\An's wind pressure on \\pBn's magnetosphere will depend on orbital phase.  This variable pressure is one mechanism for the variation of flux and profile shape of \\pB with respect to the orbital phase that has been reported.  The response of \\pB to the \\A wind pressure will also depend on the particular side of its magnetosphere facing the wind at the spin phase when \\B is visible. This is a second  possible mechanism for variability. We suggest that \\pB may have its spin axis aligned with the orbital angular momentum owing to \\An's wind torque that contributes to its spindown.  Monitoring the pulsars while geodetic precession changes spin orientations will provide essential evidence to test detailed theoretical models. We determine the Rotation Measures of the two stars to be $-112.3\\pm 1.5$ and $-118\\pm 12$ rad m$^{-2}$. ",
        "introduction": "\\label{sec-intro} The double pulsar system, PSR J0737--3039, recently reported by Burgay et al. (2003) and Lyne et al. (2004) is likely to unlock many mysteries concerning isolated neutron star magnetospheres and winds as well as the nature of the pulsar emission mechanism. A 22.7-ms pulsar (\\An) and a 2.77 sec pulsar (\\Bn) revolve about each other in a 2.4h, nearly edge-on orbit. The two stars show interesting and contrasting emission properties. Pulsar \\A exhibits a complex double profile structure. Apart from an eclipse when \\A is behind \\B (with respect to our sight line), this pulsar does not show significant variation in its flux. The flux and profile structure of \\B show remarkable variations as a function of orbital phase (Lyne et al. 2004).   A simple pressure balance calculation leads to the conclusion that the MHD wind of \\pA will surpress a large fraction of the quasi-static \\B magnetosphere.  The eclipse duration of \\pA is roughly consistent with this idea (Lyne et al. 2004; Kaspi et al. 2004). Detailed calculations of the \\An-wind/\\Bn-magnetosphere interactions are required for comparisons with observations (Arons et al. 2004). In this model a bow shock, magnetosheath, magnetopause and magnetotail structures, which are relativistic analogs of solar wind/Earth magnetosphere structures, are established around \\pBn.  In this work, we first present polarimetric measurements on these two stars. These observations were conducted at the National Radio Astronomy Observatory\\footnote{The National Radio Astronomy Observatory (NRAO) is owned and operated by Associated Universities, Inc under contract with the National Science Foundation.} Green Bank Telescope (GBT). After describing our observational setup in \\S\\ref{sec-observation}, we present our results in \\S\\ref{sec-result}. Our observations place strong constraints on the geometrical orientation of {\\bf A}'s rotation and magnetic axes. We also present our Rotation Measure (RM) measurements for the two stars. In \\S\\ref{sec-model} we discuss implications of our polarimetric results within the context of the Arons et al. model. ",
        "conclusions": "\\label{sec-model}   A remarkable feature of both the Lyne et al. (2004) and our GBT observations of  J0737--3039B is the dependence of  flux and pulse profile, including polarization properties, on orbital phase, but not on radio frequency. We concur in general with Lyne et al. that the variable influence of \\pA on the magnetosphere of \\B is the likely source of the changes. In a companion paper (Arons et al. 2004) we develop a detailed model in which \\An's MHD wind confines the \\B magnetosphere. We outline this model below, and make connections between observational features that both stimulated the model making and provide directions for future work.  In our model the MHD wind from \\A decelerates through a relativistic bow shock that envelops \\Bn's magnetosphere. Between this shock and the boundary of \\Bn's closed magnetosphere (\\Bn's magnetopause) the shocked \\A wind plasma forms a relativistic hot layer whose optical depth to synchrotron absorption at 500 MHz is at least 100, and may be as high as 5000. This layer (the magnetosheath) is the likely cause of the eclipse of \\A by \\Bn. In this model we further interpret the faintest region of \\B emission near its superior conjunction as the result of absorption by the same magnetosheath layer that in these phases is between the observer and the pulsar emission region deep within the \\B magnetosphere. In Kaspi et al. (2004) we report that the \\A eclipse is asymmetric with slower flux decrease on ingress and a more rapid recovery on egress. As discussed above the strongest and faintest regions of \\B emission are also asymmetric: both precede times of conjunction by $\\sim 30^\\circ$. In our model these are attributed to the prograde rotation of \\B that leads to an asymmetric magnetopause and magnetotail. The frequency independence of the eclipse light curve over our observing bands is attributed to sharp boundaries, high optical depths and partial covering in space and/or time.  We favor the interpretation of our polarization observations  of \\pA that has the dipole field axis nearly aligned with the spin axis, which itself is oblique to the orbital plane and the line of sight. Pulsar \\B would then experience a significant difference in wind pressure and content around the orbit.  This variable pressure will affect the size of the polar cap and current structure that could lead to the \\B variations.   The \\A wind, asymmetric or not, will produce a propeller torque on \\B owing to \\Bn's rotation.  This torque contributes to the observed $\\dot{P}$ of \\Bn.  In this model the \\B emission results from voltage and $e\\pm$-avalanche current deep within its magnetosphere similar to that in normal pulsars. The ``normalcy'' of \\B emission is supported by these polarization observations as well as the single pulse study described elsewhere (Ramachandran et al.  2004).  When we observe \\pBn, it presents a different face of its closed and confined magnetosphere to the \\A wind depending on the orbital phase. This variable orientation will also change the \\B magnetosphere and, we suggest, its polar cap and the beamed emission we detect. At this point we cannot distinguish between these two possible sources of modification of the secondary star -- variable \\A wind pressure impinging on \\Bn, and variable \\B internal structure at the rotation phase of observation. They both are seemingly consistent with the frequency independence of the reported phenomena. Both cases will lead to variations in the polar cap size and relativistic current structure of \\Bn, which will affect the observed flux and its pulse morphology.   In the case of the variable wind pressure, the phase of modulation of \\B is expected to precess around the orbit as the spin of \\A undergoes geodetic precession. Geodetic precession of \\B will have a more complex effect on its internal response to a wind pressure at the spin phase of observation owing to the multiplicity of angles involved.   Our proposed geometry of \\An's beam will be testable as geodetic precession moves the observer through the cone and into a region of invisibility during the 70-year cycle (Lyne et al. 2004). Is it improbable that we are seeing both pulsars at the same time? We suggest that the \\A wind torque on \\B has aligned pulsar \\Bn's spin axis with the orbital angular momentum over time. In this case \\B would need to have its dipole axis at $\\alpha\\sim 90^\\circ$, and therefore we will continue to view \\B independent of geodetic precession. An interesting question is whether the torque of the \\A wind can misalign the dipole while aligning the spin. The full history of the evolution of the system and the dueling magnetospheric winds remains to be written. In its infancy \\B might have had a short spin period and strong magnetic field. If so, then the early wind from \\B could have dominated the \\A magnetosphere and altered its spin and magnetic dipole.  \\subsection{Summary} \\label{sec:sum} Our polarimetric observations indidate that pulsar {\\bf A}'s spin and magnetic axes are nearly aligned. This leads us to consider two possible mechanisms for the variability of \\B with orbital phase -- (1) a pole to equator asymmetry of the \\A wind, and (2) a rotational asymmetry of the force balance radius of the \\B magnetosphere as it adjusts to the \\A wind. The interaction also contributes to the spin-down torque on the \\B pulsar which may also lead to alignment of the \\B spin and orbital momentum vectors. Geodetic precession will provide an important means of separating the relative importance of the various effects in the coming years. Theoretical calculations are underway to explore these ideas quantitatively."
    },
    "0402/astro-ph0402281_arXiv.txt": {
        "abstract": "Non-helical hydromagnetic turbulence with an externally imposed magnetic field is investigated using direct numerical simulations. It is shown that the imposed magnetic field lowers the spectral magnetic energy in the inertial range. This is explained by a suppression of the small scale dynamo. At large scales, however, the spectral magnetic energy increases with increasing imposed field strength for moderately strong fields, and decreases only slightly for even stronger fields. The presence of Alfv\\'en waves is explicitly confirmed by monitoring the evolution of magnetic field and velocity at one point. The frequency $\\omega$ agrees with $v_{\\rm A}k_1$, where $v_{\\rm A}$ is the Alfv\\'en speed and $k_1$ is the smallest wavenumber in the box. ",
        "introduction": "Turbulent magnetic fields are seen in many astrophysical settings \\cite{Beck_etal96,BH98,Biskamp03}. Such magnetic fields usually result from the conversion of kinetic energy into magnetic energy, i.e.\\ from dynamo action. Numerical simulations show that a dynamo-generated magnetic field can be of appreciable strength even when there is no kinetic helicity \\cite{CV00a,Scheko02}. Simulations have recently also shown that at scales smaller than about five times the energy carrying scale the magnetic energy spectrum seems to enter an inertial subrange where the magnetic spectral energy exceeds the kinetic spectral energy \\cite{HBD03}. This means that over any subvolume, whose scale is within the inertial range, there is always a larger scale component of the field with significant strength. This raises the question whether one can model the small scale properties of such turbulence simply by imposing a magnetic field.  A lot of work has already been devoted to studying hydromagnetic turbulence in the presence of an external field \\cite{CV00,MG01,CLV02}. Nevertheless, the super-equipartition magnetic energy seen in simulations without imposed field has never been seen in simulations with imposed field. An exception is when the magnetic Prandtl number is large\\cite{CLV02b}. However, the super-equipartition is then seen between the viscous and the resistive cutoff -- not in the inertial range. It is one of our goals to elucidate this puzzle. Likewise, although dynamos with helicity can produce substantial super-equipartition on the scale of the system, they too are not able to produce super-equipartition in the inertial range \\cite{B01}. In that sense the difference between dynamos with and without imposed field is similar to the difference between helical and nonhelical dynamos.  The views on the effects of external fields are divided. A common scenario that applies when the conditions for dynamo action are not met (e.g.\\ if the magnetic Reynolds number is too small), is one where a local magnetic field can be enhanced simply by winding up an external magnetic field. Possible candidates where this may be the case are Io and Ganymede, in which convection interacts with the field of Jupiter leading to local field enhancement \\cite{Schubert96,Khurana97,Sarson97}. A similar possibility may also apply to the solar convection zone where the large scale field of the 11 year solar cycle is primarily located at the bottom of the convection zone \\cite{SW80}, but the overlying convection zone may shred the field to produce small scale field \\cite{Sch84}. Another possibility that has been discussed more recently is that the small scale field at the solar surface could be generated locally by a small scale dynamo operating near the surface \\cite{Cat99}.  In hydromagnetic turbulence theory magnetic and kinetic energy are assumed to cascade from large to small scales, similar to the hydrodynamic case, although recent work has established a strong intrinsic anisotropy \\cite{GS95}, which has no counterpart in the hydrodynamic case. However, this theory does not address the possibility of dynamo action. It remains therefore an open question as to what is the nature of the interaction resulting from imposed and dynamo-generated magnetic fields. In particular, we shall present evidence that the imposed magnetic field does not enhance dynamo action. Instead, the external field does actually suppress dynamo action, albeit in a subtle way because the rms turbulent velocity is generally {\\it not} decreased by a modestly strong magnetic field. We show that the suppression can be associated with the work term resulting from the Lorentz force due to the imposed field. It turns out that this term changes sign above a certain field strength such that a certain fraction of magnetic energy flows backwards to enhance the kinetic energy instead. ",
        "conclusions": "The present studies have shown that a uniformly imposed magnetic field has two important effects on the magnetic field that is induced at finite wavenumbers ($k\\neq0$). First, the magnetic field is slightly enhanced at and around the forcing wavenumber (corresponding to the energy carrying scale). Second, the magnetic field is quenched with increasing $B_0$ at all larger wavenumbers corresponding to the inertial and diffusive subranges.  The enhancement and suppression at the two different wavenumber ranges is associated with a corresponding wavenumber dependence of the work term, $\\BB_0\\cdot\\bra{\\uu\\times\\jj}$. The suppression of the magnetic field in the inertial range is quite opposite to the behavior without imposed field when there is instead a significant enhancement of the magnetic energy spectrum over the kinetic energy spectrum. We therefore refer to this effect as a suppression of the dynamo by the imposed field.  The suppression of dynamo activity might be a consequence of the tendency toward two-dimensionalization of the turbulence by the large scale field \\cite{MBG03}. Such an effect is well-known for low-$\\Rm$ hydromagnetic turbulence \\cite{Schumann76}, and it is a mathematical theorem that there can be no dynamo action in two dimensions \\cite{Zeldovich57}. Of course, the turbulence does not really become two-dimensional, but instead the correlation length along the field becomes large. This type of anisotropy is a crucial ingredient of the Goldreich-Sridhar theory \\cite{GS95}.  The Goldreich-Sridhar theory also predicts that Alfven waves should be present in the system. This has been confirmed by inspecting velocities and magnetic fields at a single point in the middle of the simulation box. These Alfv\\'en waves have the expected frequency $\\omega_A=v_{\\rm A}k_1$. Furthermore, we do not find that there is equipartition between magnetic and kinetic energy spectra in the inertial range for large imposed field strengths. The absence of equipartition may be a consequence of the inertial range being still too short (or absent). In runs where $B_0=B_{\\rm eq}$, on the other hand, there is clear evidence that kinetic and magnetic energy spectra fall on top of each other throughout the dissipation subrange. This is also in agreement with earlier results of Cho and collaborators \\cite{CLV02}, who considered the case where the imposed field had equipartition strength.  Whether or not models with imposed field can reproduce the situation in small sub-domains of simulations with no overall imposed field is still unclear. At first glance the answer seems to be no, because none of the simulations with imposed field have ever been able to produce super-equipartition in the inertial range, as it is seen in the nonhelical simulations without imposed field \\cite{HBD03}. However, the reason for this may well lie in the still insufficient resolution of the simulations with no imposed field -- even though they do already have a resolution of $1024^3$ meshpoints. It is indeed possible that, even though the kinetic and magnetic energy spectra are approximately parallel to each other over a certain range of wavenumbers and offset by a factor of about 2.5, they may actually converge at still larger wavenumbers. Preliminary indications of this have now been seen in simulations using hyperviscosity and hyper-resistivity with no imposed field. However, a general difficulty with hyperviscosity and hyper-resistivity is that certain aspects of the physics of such systems are significantly modified \\cite{BS02}. It is therefore equally important to assess the features that are likely not to be altered by this manipulation. A detailed discussion of this will be the subject of a forthcoming paper."
    },
    "0402/astro-ph0402248_arXiv.txt": {
        "abstract": "Star formation plays an important role in the fate of interacting galaxies.  To date, most galactic simulations including star formation have used a density-dependent star formation rule designed to approximate a Schmidt law.  Here, I present a new star formation rule which is governed by the local rate of energy dissipation in shocks. The new and old rules are compared using self-consistent simulations of NGC\\,4676; shock-induced star formation provides a better match to the observations of this system. ",
        "introduction": "Numerical simulations of galaxy formation and interactions often include rules for star formation.  In many cases, the local rate of star formation, $\\dot{\\rho}_{*}$, is related to the local gas density, $\\rho_{\\rm g}$, by a power law: \\begin{equation} \\dot{\\rho}_{*} = C_{*} \\, \\rho_{\\rm g}^n \\, , \\label{sfr-power-law} \\end{equation} where $C_{*}$ is a constant.  This prescription has been justified both empirically and theoretically.  From an empirical perspective, (\\ref{sfr-power-law}) resembles the ``Schmidt law'' \\citep{S59}, which relates star formation per unit surface area, $\\dot{\\Sigma}_{*}$ to gas surface density, $\\Sigma_{\\rm g}$; a recent determination of the Schmidt law \\citep{K98} is \\begin{equation} \\dot{\\Sigma}_{*} = (2.5 \\pm 0.7) \\times 10^{-4} \\, \\frac{{\\rm M}_{\\sun}}{{\\rm yr} \\, {\\rm kpc}^{2}} \\, \\left( \\frac{\\Sigma_{\\rm g}}{1 {\\rm M}_{\\sun} {\\rm pc}^{-2}} \\right) ^{1.4 \\pm 0.15} \\, . \\label{schmidt-kennicutt-law} \\end{equation} \\citet{MRB91} adopted a rule equivalent to (\\ref{sfr-power-law}) with $n = 2$ as an approximation to a Schmidt law with index $\\sim 1.8$, while \\citet{MH94a} took $n = 1.5$, and presented numerical tests showing this gave a reasonable match to a Schmidt law with index $\\sim 1.5$.  A more theoretical approach, adopted by \\citet{K92} and \\citet{S00}, sets $\\dot{\\rho}_{*} = \\rho_{\\rm g} / t_{*}$, where $t_{*}$, the time-scale for star formation, is basically proportional to the local collapse time of the gas, $(G \\rho_{\\rm g})^{-1/2}$. This yields (\\ref{sfr-power-law}) with $n = 1.5$ and $C_{*} = G^{1/2} C'_{*}$, where $C'_{*}$ is a dimensionless constant.  Despite the apparent convergence of observation and theory on the index $n = 1.5$, it's unlikely that (\\ref{sfr-power-law}) really captures the process of star formation.  For one thing, only about $1$~per cent of the gas actually forms stars per collapse time $(G \\rho_{\\rm g})^{-1/2}$; in other words, consistency with the observations implies that $C'_{*} \\simeq 10^{-2}$.  Current theories of star formation don't offer any straightforward way to calculate this quantity.  Gravitational collapse is evidently not the limiting factor which sets the rate of star formation; ``feedback'' processes are important in determining the fraction of available interstellar material which ultimately becomes stars.  While several groups have now devised simulations including star formation regulated by various forms of feedback (e.g.~\\citealt{GI97}; \\citealt{S00}), this approach still has some way to go.  Moreover, models based on (\\ref{sfr-power-law}) don't reproduce the large-scale star formation seen in many interacting galaxies. \\citet{MRB92} found that most of the star formation was confined to the central regions of their model galaxies, while \\citet{MBR93} noted that (\\ref{sfr-power-law}) underestimated the rate of star formation in regions where interstellar gas exhibits large velocity dispersions and gradients, or where the gas appears to be undergoing strong shocks.  \\citet{MH94b}, \\citet{MH96}, and \\citet{S00} have modeled the central bursts of star formation seen in ultraluminous infrared galaxies \\citep[][and references therein]{SM96}.  But nuclear starbursts, while a necessary stage in the transformation of merger remnants into elliptical galaxies \\citep[e.g.][]{KS92}, may not be sufficient to accomplish this transformation.  For example, merger-induced starbursts create massive young star clusters \\citep[e.g.][]{WS95} which may subsequently evolve into globular clusters, but these clusters will be confined to the nuclei of remnants {\\it unless\\/} the starbursts are spatially extended.  In view of these considerations, it's worth examining alternatives to (\\ref{sfr-power-law}).  One long standing idea is that collisions of molecular clouds trigger of star formation (e.g.~\\citealt{SSC86}). This has been implemented in ``sticky particle'' schemes which model the interstellar medium as a collection of discrete clouds undergoing inelastic collisions (e.g.~\\citealt{OK90}; \\citealt{N91}).  However, molecular clouds have relatively long mean free paths, and direct collisions between clouds at velocities of $10^2 {\\rm\\,km/s}$ or more may result in disruption rather than star formation.  \\citet{JS92} proposed a model of shock-induced star formation in interacting galaxies; specifically, they suggested that fast collisions between extended H{\\small{I}} clouds create a high-pressure medium which compresses pre-existing molecular clouds and thereby induces bursts of star formation.  High-pressure regions, and especially large-scale shocks in colliding galaxies, may favor the formation of bound star clusters \\citep{EE97}.  Recent observations at optical, infrared, and radio wavelengths continue to reveal large-scale star formation in interacting galaxies, including NGC\\,4038/9, Arp\\,299, and NGC\\,4676 (\\citealt{WS95}; \\citealt{V+96}; \\citealt{M+98}; \\citealt{HY99}; \\citealt{AH+00}; \\citealt{X+00}; \\citealt{dG+03}).  Several of these studies explicitly invoke shock-induced star formation in discussing the observational data.  In this paper I focus on NGC\\,4676, a strongly interacting pair of disk galaxies with long tidal tails (\\citealt{TT72}, hereafter TT72; \\citealt{T77}), as a test-case for models of shock-induced and density-dependent star formation.  \\S~2 describes the star formation algorithms.  \\S~3 presents simulations of NGC\\,4676, and contrasts the results of shock-induced and density-dependent star formation rules. Conclusions are given in \\S~4. ",
        "conclusions": "Density-dependent and shock-induced models of star formation yield qualitatively and quantitatively different results.  These differences arise despite the fact that the gas density $\\rho_{\\rm g}$ and the dissipation rate $\\dot{u}$ are not completely independent variables; indeed, they can be explained in terms of the global relationship between $\\rho_{\\rm g}$ and $\\dot{u}$.  Dissipation is a necessary precursor for any significant increase in central gas density, since gas can only accumulate in galactic centers as a result of an irreversible process.  The rate of increase in gas density closely tracks the net energy radiated away in shocks (\\citealt{BH96}, Fig.~6); globally, $\\dot{u}$ and $d\\rho_{\\rm g}/dt$ are strongly correlated.  Thus, models of shock-induced star formation can respond promptly to external disturbances, while little activity occurs in density-dependent models until $\\rho_{\\rm g}$ has had time to build up.  The larger spatial extent of star formation in shock-induced models is, in part, a corollary of the earlier onset of activity in such models, since the gas is more widely distributed at earlier times.  Density-dependent rules may be used to implement unified models of star formation in normal and interacting galaxies.  The disks of normal galaxies {\\it and\\/} the central regions of starburst galaxies fit onto the same power law \\citep{K98}.  Thus, by setting $C_{*}$ to match the baseline rate of star formation in unperturbed disks, \\citet{MH96} were able to produce starbursts comparable to those inferred in ultraluminous infrared galaxies \\citep{SM96}; this represents a real success for the model defined by (\\ref{sfr-power-law}) with $n = 1.5$.  But there is abundant evidence for large-scale star formation in interacting galaxies.  This includes ongoing star formation in the ``overlap regions'' of systems like NGC\\,4038/9 and Arp\\,299, as well as the H$\\alpha$ emission from the tails of NGC\\,4676.  It also includes the A-star spectra seen in the tails of NGC\\,4676 and throughout the body of NGC\\,7252.  Density-dependent star formation can't easily explain these observations; interactions funnel gas from the disks to the central regions of galaxies, thereby promoting rapid star formation in galactic nuclei but reducing the supply of gas needed for star formation elsewhere.  Moreover, violent relaxation ceases long before binding energies are effectively randomized (e.g.~\\citealt{W79}; \\citealt{B92}), so merging is ineffective at transporting the products of star formation outward from nuclei to the bodies of merger remnants.  Thus, density-dependent rules offer at best an incomplete description of star formation in interacting galaxies.  In simulations with shock-induced rules, the star formation rate depends on dynamical circumstances.  Shocks in unperturbed disks are associated with transient spiral patterns, and matching the baseline level of star formation depends on reproducing the ``right'' level of spiral structure.  This, in itself, is a tough problem, and it's not clear that shock-induced models will soon yield a unified description of star formation in normal {\\it and\\/} interacting galaxies.  Until this becomes possible, unperturbed disks can't be used to set $C_{*}$, so the amplitude of a shock-induced starburst can't be predicted {\\it a priori\\/}.  Nonetheless, shock-induced star formation is an important element of galactic collisions, and its implementation in numerical simulations is a useful step toward increased realism.  These simulations have contrasted two limits of (\\ref{sfr-two-power-law}): density-dependent star formation, with $n > 1$ and $m = 0$, and shock-induced star formation, with $n = 1$ and $m > 0$.  These cases were chosen as instructive examples, but (\\ref{sfr-two-power-law}) is general enough to accommodate additional possibilities.  First, setting $n > 1$ and $m$ just slightly larger than $0$ would yield a modified law in which star formation is proportional to $\\rho_{\\rm g}^n$ but occurs {\\it only\\/} in regions with $\\dot{u} > 0$; this resembles the rule adopted by \\citet{K92}, who basically took $\\dot{\\rho}_{*} \\propto \\rho_{\\rm g}^{3/2}$ but restricted star formation to regions with convergent flows.  Notice that cases with $0 < m \\ll 1$ are {\\it not\\/} continuous with the case $m = 0$; in the latter, star formation is strictly independent of $\\dot{u}$.  Second, setting $n > 1$ and $m > 0$ would yield hybrid rules in which star formation depends on both $\\rho_{\\rm g}$ and $\\dot{u}$.  The consequences of such rules can sometimes be inferred from the limiting cases considered here.  For example, in the simulations of NGC\\,4676 with shock-induced star formation, the gas involved in burst B$_3$ has densities $\\sim 10^2$ times higher than the gas involved in burst B$_1$; setting $n > 1$ would boost B$_3$ relative to B$_1$ by a factor of $\\sim 10^{2 (n - 1)}$, assuming that neither burst was limited by the supply of gas.  Perhaps most exciting are the new avenues for research created by an alternative description of star formation in galaxy interactions:  \\begin{itemize}  \\item In contrast to density-dependent star formation, which seems fairly insensitive to most details of galactic encounters \\citep{MH96}, different encounter geometries may yield a variety of shock-induced star formation histories.  A survey of different encounters, along the lines of other surveys without star formation (e.g.~\\citealt{B02}), might establish if widespread star formation at first passage is generic or limited to a subset of close encounters, and also determine if shock-induced star formation can account for ultraluminous infrared galaxies.  \\item Shock-induced star formation makes definite predictions about the {\\it timing\\/} of starbursts triggered by galactic encounters.  As already mentioned, models including the photometric evolution of starburst populations could sharpen the comparison between simulations and observations of systems like NGC\\,4676.  In addition, these models might predict the ages of embedded young star clusters; such predictions could be checked by multi-object spectroscopy of the clusters.  \\item In the context of merger simulations, shock-induced star formation also yields predictions for the spatial distribution and kinematics of starburst populations.  It would be quite interesting to compare these predictions with observations of the distribution and kinematics of metal-rich globular clusters in elliptical galaxies, which may have been formed in merger-induced starbursts (e.g.~\\citealt{ZA93}).  \\item Encounters between gas-rich galaxies at redshifts of a few seem a natural setting for shock-induced star formation.  While many high-$z$ galaxies have peculiar morphologies (e.g.~\\citealt{vdB+96}), the bridges and tails characteristic of low-$z$ encounters (TT72) are not very evident.  Instead, the optical morphology of these objects may be dominated by rest-frame UV from widespread starbursts; if so, the photometric modeling approach described above could help interpret existing and forthcoming observations of interacting galaxies at intermediate and high redshifts.  \\end{itemize}"
    },
    "0402/astro-ph0402138_arXiv.txt": {
        "abstract": "The observed size distribution of Kuiper belt objects (KBOs)---small icy and rocky solar system bodies orbiting beyond Neptune---is well described by a power law at large KBO sizes. However, recent work by \\cite{bernstein03} indicates that the size spectrum breaks and becomes shallower for KBOs smaller than about 70~km in size. Here we show that we expect such a break at KBO radius $\\sim$40~km since destructive collisions are frequent for smaller KBOs. Specifically, we assume that KBOs are rubble piles with low material strength rather than solid monoliths. This gives a power-law slope $q\\simeq 3$ where the number $N(r)$ of KBOs larger than a size $r$ is given by $N(r) \\propto r^{1-q}$; the break location follows from this slope through a self-consistent calculation. The existence of this break, the break's location, and the power-law slope we expect below the break are consistent with the findings of \\cite{bernstein03}. The agreement with observations indicates that KBOs are effectively strengthless rubble piles. ",
        "introduction": "The Kuiper belt, a population of small bodies moving beyond the giant planets, was discovered when its first member was found in 1992 \\citep{jewitt93}.  As of late 2003, $\\sim$800 KBOs have been discovered. Due to KBOs' faintness, however, the size distribution of KBOs is well determined observationally only for bodies larger than $\\sim$100~km \\citep{trujillo01,gladman98,chiang99}; their size spectrum is consistent with a power law $N(r)\\propto r^{-4}$ \\citep{bernstein03}. Numerical studies concluded that the differential size spectrum below $\\sim$100~km should follow a power law with the slightly shallower $N\\propto r^{-2.5}$ due to the effects of destructive collisions \\citep{farinella96,davis97,kenyon02}.  The results seemed consistent with loose observational constraints available on the number of $\\sim$20~km and $\\sim$2~km KBOs \\citep{cochran95,holman93}.  In this context, the deficit in small KBOs observed by \\cite{bernstein03} was a surprise. Using the Advanced Camera for Surveys recently installed on the Hubble Space Telescope, they found just 3 KBOs of size $\\sim$25--45~km where they expected $\\sim$85 such bodies based on an extrapolation of the accepted best-fit large-KBO spectrum at the time \\citep{trujillo01}. While this observed decrement of more than an order of magnitude in the number of small KBOs clearly indicates a break between 45 and 100~km, the exact break position and slope below the break may well be refined by future data on small KBOs. Still, the results of \\cite{bernstein03} are inconsistent with the previously expected small-end spectrum $N\\propto R^{-2.5}$, or $q=3.5$, at better than 95\\% confidence.  This paper describes a simple self-consistent analytic calculation of the break location and the slope below the break. Note that using the $N(r)\\propto r^{-4}$ size spectrum obtained by \\cite{bernstein03} for large KBOs, we can estimate the size below which collisions between equal size bodies should be frequent to be $\\sim$ 1~km---well below the observed break location. However, this estimate needs two modifications.  First, due to the large velocity dispersion in the Kuiper belt, small bodies can shatter much larger objects. Since there are more small than large bodies, destructive collisions will occur frequently even for objects much larger than 1~km. Second, when collisions are important, they reduce the number of small bodies; this in turn decreases the frequency of collisions. Therefore, calculations of the effects of collisions and the size below which collisions are important must be done in a self-consistent manner. ",
        "conclusions": "We have derived a self-consistent size spectrum $23/8<q<22/7$ for a collisional population of bodies whose binding energy is dominated by gravity. We emphasize that this spectrum does not truly represent a steady state; instead, the number density of bodies decreases slowly compared to the collision timescale. For the case of the Kuiper Belt, the spectrum's small-end power-law slope $q\\simeq 3$ and break radius $r_{\\rm break}\\sim 40\\;\\km$ agree well with those found observationally by \\cite{bernstein03}. Since the power-law slope derived in the steady-state approximation depends heavily on the particular criterion for catastrophic destruction adopted for the bodies, observations of the KBO size spectrum provide a direct constraint on the bodies' internal structure.  The close agreement between this slope and break radius and the best-fit values found by \\cite{bernstein03} suggests that large KBOs are virtually strengthless bodies held together mainly by gravity.  Further surveys of small KBOs between $\\sim$10 and $\\sim$70~km in size would better constrain both the exact position of the actual break in the size spectrum and the power-law slope below the break. Data of this kind would thus confirm or refute our analysis. Such surveys would also allow more detailed comparison of the break locations in the classical and scattered KBO populations, which should reflect differences in the surface densities and velocity dispersions for those two groups.  \\noindent {\\em Acknowledgements.} We thank Oded Aharonson and Andrew McFadyen for useful discussions. MP is supported by an NSF graduate research fellowship."
    },
    "0402/astro-ph0402624_arXiv.txt": {
        "abstract": "{We present the results of our analysis of a high resolution ($R\\simeq30000$) infrared spectrum of V4334~Sgr (Sakurai's Object) around the first overtone CO bands, obtained in 1998 July. The $^{12}$CO and $^{13}$CO bands are well-resolved, and we compute synthetic hydrogen-deficient model atmosphere spectra to determine the \\CDC\\ ratio. We find \\CDC\\ $\\simeq4\\pm1$, consistent with the interpretation of V4334~Sgr as an object that has undergone a very late thermal pulse. ",
        "introduction": "V4334~Sgr (Sakurai's Object) is a low mass ($\\sim0.6$-$0.8\\Msun$; \\citealt{herwig}) star undergoing a very late thermal pulse (VLTP). After traversing the `knee' on the post-asymptotic giant branch (AGB) evolutionary track, the star experiences a final He flash as it evolves down the white dwarf cooling track: it becomes a `born-again' AGB star (see \\citealt{herwig, lawlor}).  A key factor in confirming this scenario is the determination of elemental abundances and isotopic ratios in the surface material. \\cite{asplund99} have used optical spectroscopy to determine the abundance of a range of species; they found V4334~Sgr to be hydrogen-deficient, with other elemental abundances close to those of the R~CrB stars; R~CrB stars. They also used the C$_2$ (1-0) and (0-1) Swan bands at 4740 and 5635\\AA\\ respectively to determine that the \\CDC\\ ratio was somewhere in the range 1.5--5, although they did not give an actual value.  \\cite{pav-trg} modelled the spectral energy distribution of V4334~Sgr in the near-IR ($\\lambda\\lambda$ 1 - 2.5 \\mic). They broadly confirmed the results of \\cite{asplund99}, but found some evidences of higher C abundance, by a factor $\\sim2$. \\citeauthor{pav-trg} also noted the presence of hot dust longward of $\\sim2\\mic$.  The \\CDC\\ ratio is of fundamental importance in understanding the VLTP process, as it is a reflection of the relative importance of $^{12}$C, $^{13}$C processing and dredge-up. In this Letter we use echelle observations of the first overtone vibration-rotation features of CO, obtained in 1998 July while the star was still visible, to determine the \\CDC\\ ratio in V4334~Sgr. ",
        "conclusions": "We have carried out a spectral synthesis to determine the \\CDC\\ ratio in V4334~Sgr. Given a range of C abundance, we determine that \\CDC\\ $\\simeq4\\pm1$, consistent with the VLTP interpretation of V4334~Sgr. The relative dust contribution at 2.3\\mic\\ is $\\epsilon_0\\simeq2\\pm0.2$. Our value of the \\CDC\\ ratio is quite well constrained, and will doubtless be useful in fine-tuning the VLTP scenario."
    },
    "0402/astro-ph0402468_arXiv.txt": {
        "abstract": "In recent years, indirect evidence has emerged suggesting that many nearby non-active galaxies harbor quiescent supermassive black holes. Knowledge of the frequency of occurrence of black holes, of their masses and spins, is of broad relevance for studying black hole growth and galaxy and AGN formation and evolution. It has been suggested that an unavoidable consequence of the existence of supermassive black holes, and the best diagnostic of their presence in non-active galaxies, would be occasional tidal disruption of stars captured by the black holes. These events manifest themselves in form of luminous flares powered by accretion of debris from the disrupted star into the black hole.  Candidate events among optically non-active galaxies emerged in the past few years. For the first time, we have looked with high spatial and spectral resolution at one of these most extreme variability events ever recorded among galaxies. Here, we report measuring a factor $\\sim 200$ drop in luminosity of the X-ray source RX\\,J1242--1119 with the X-ray observatories {\\it Chandra} and {\\it XMM-Newton}, and perform key tests of the favored outburst scenario, tidal disruption of a star by a supermassive black hole. We show that the detected ``low-state'' emission has properties such that it must still be related to the flare. The power-law shaped post-flare X-ray spectrum indicates a ``hardening'' compared to outburst. The inferred black hole mass, the amount of liberated energy, and the duration of the event favor an accretion event of the form expected from the (partial or complete) tidal disruption of a star. ",
        "introduction": "The X-ray luminous nuclei of active galaxies (AGN) are believed to be powered by accretion of gas onto supermassive black holes. There is now growing evidence that many nearby {\\em non-active} galaxies harbor quiescent, weakly or non-accreting black holes (see reviews by Kormendy \\& Gebhardt 2001; Richstone 2002). Studies of the abundance of black holes, of their masses and their spins, shed light on the mechanisms of black hole growth and of galaxy and AGN formation and evolution. Possibly the most direct means of detecting supermassive black holes at the centers of galaxies, and an unavoidable consequence of their existence, would be occasional tidal disruption of stars and subsequent accretion of their debris by these supermassive black holes (e.g., Hills 1975; Gurzadyan \\& Ozernoi 1979; Carter \\& Luminet 1982; Rees 1988, 1990; Wang \\& Merritt 2004). The events would appear as luminous flares of radiation emitted when the stellar debris is accreted by the black hole.  Stellar capture and disruption is -- apart from accretion of gas and black hole merging -- one of the three major processes studied in the context of black hole growth (e.g., Frank \\& Rees 1976). The relative importance of these  three processes in feeding black holes is still under investigation. Zhao, Haehnelt, \\& Rees (2002) and Merritt \\& Poon (2003) recently pointed out that tidal capture may play an important role in explaining the notable $M_{\\bullet}-\\sigma$ relation of galaxies. Given the intense theoretical attention the topic of stellar tidal disruption receives (see \\S3 of Komossa 2002, and references therein), it is of great interest to see whether these events do occur in nature, and how frequent they are.   Making use of the unique capability provided by the All-Sky Survey (RASS, Voges et al. 1999) of the X-ray satellite {\\it ROSAT\\/}, giant X-ray flares from the directions of a few nearby galaxies were discovered (Komossa \\& Bade 1999; Komossa \\& Greiner 1999; Grupe et al. 1999; Greiner et al. 2000). These flares were characterized by huge peak X-ray luminosities reaching $\\sim 10^{44}$ ergs~$s^{-1}$, large amplitudes of decline, ultrasoft X-ray spectra, and absence of Seyfert activity in ground-based optical spectra (see Komossa 2002 for a review). The target of the present study, RX\\,J1242--1119, was first detected by {\\it ROSAT\\/} in 1992 (Komossa \\& Greiner 1999) during a pointed observation with the Position Sensitive Proportional Counter (PSPC). At that time it showed a very soft X-ray spectrum with $kT_{\\rm bb} \\simeq 0.06$~keV, and an X-ray luminosity of $L_{\\rm x} \\simeq 9 \\times 10^{43}$ ergs~$s^{-1}$, which is exceptionally large given the absence of any signs of Seyfert activity of RX\\,J1242--1119 in ground-based (Komossa \\& Greiner 1999) and {\\it Hubble Space Telescope} ({\\it HST}, Gezari et al. 2003) optical spectra. The association with one of the two, previously unstudied non-active galaxies at redshift $z=0.05$ located in the X-ray position error circle (named RX\\,J1242--1119A and RX\\,J1242--1119B for lack of a better designation) however, remained uncertain because of the large {\\it ROSAT\\/} positional uncertainty, in this case 40$^{\\prime\\prime}$ (Figure 1).  Considering the extraordinary properties of this and a few similar X-ray events, particularly their enormous luminosity output, it is of utmost importance to understand what is happening in these systems. For the first time, we have now looked at one of these flare events with high spatial and spectral resolution, in order to confirm the counterpart, follow the long-term temporal behavior, study the spectral evolution and measure the post-flare spectrum, and to use these results to test the favored outburst model: tidal disruption of a star by a supermassive black hole. Among the few known X-ray flaring galaxies, RX\\,J1242--1119 was our target of choice for follow-up X-ray observations because it flared most recently, so the probability of catching the source in the declining phase was highest. Also, because there were two galaxies in the {\\it ROSAT\\/} X-ray error box, the correct counterpart needed to be confirmed. A Hubble constant of 50 km~s$^{-1}$~Mpc$^{-1}$ is used throughout this paper. ",
        "conclusions": "The results presented here demonstrate the effectiveness of combining the abilities of the two most powerful X-ray observatories in orbit to study one of the most extreme variability events ever recorded among galaxies. In our interpretation of these observations, we are seeing the post-disruption phase of a close encounter of a star with a central supermassive black hole, into which some of the tidal debris is accreting. Continued X-ray monitoring of RX\\,J1242--1119A will enable us to follow the expected further decline in luminosity.  Future X-ray (all)-sky surveys planned for missions like {\\it DUO\\/}, {\\it ROSITA\\/}, {\\it LOBSTER\\/}, and {\\it MAXI\\/} will be very useful in finding new flare events, while their detailed study will become possible with future high-throughput X-ray missions like {\\it XEUS\\/} and {\\it Con-X\\/}. We may then be able to probe the regime of strong gravity, since the temporal evolution of the flare X-ray emission is expected to depend on relativistic precession effects in the Kerr metric.  The results presented here will impact further model calculations of the tidal disruption process, which are still complex and challenging, and are expected to motivate new studies of the flare host galaxies as well as an expanded search for similar flare events in existing and future {\\it Chandra} and {\\it XMM-Newton} archives, including the deep fields.  Such observations of flares open up a new window to study supermassive black holes and their environment, and the physics of accretion events in otherwise inactive galaxy nuclei."
    },
    "0402/astro-ph0402142_arXiv.txt": {
        "abstract": "{ We report EVN, MERLIN and VLBA observations at 18\\,cm, 6\\,cm and 3.6\\,cm of the Seyfert galaxies \\object{NGC~7674}, \\object{NGC~5506}, \\object{NGC~2110} and \\object{Mrk~1210} to study their structure and proper motions on pc scales and to add some constraints on the many possible causes of the radio-quietness of Seyferts. The component configurations in \\object{NGC~7674} and \\object{NGC~2110} are simple, linear structures, whereas the configurations in \\object{NGC~5506} and \\object{Mrk~1210} have multiple components with no clear axis of symmetry. We suggest that \\object{NGC~7674} is a low-luminosity compact symmetric object. Comparing the images at different epochs, we find a proper motion in \\object{NGC~7674} of $(0.92\\pm0.07)~c$ between the two central components separated by 282~pc and, in \\object{NGC~5506}, we find a $3~\\sigma$ upper limit of $0.50~c$ for the components separated by $3.8~{\\rm pc}$. Our results confirm and extend earlier work showing that the outward motion of radio components in Seyfert galaxies is non-relativistic on pc scales. We briefly discuss whether this non-relativistic motion is intrinsic to the jet-formation process or results from deceleration of an initially relativistic jet by interaction with the pc or sub-pc scale interstellar medium.  We combined our sample with a list compiled from the literature of VLBI observations made of Seyfert galaxies, and found that most Seyfert nuclei have at least one flat-spectrum component on the VLBI scale, which was not seen in the spectral indices measured at arcsec resolution. We found also that the bimodal alignment of pc and kpc radio structures displayed by radio galaxies and quasars is not displayed by this sample of Seyferts, which shows a uniform distribution of misalignment between $0^{\\circ}$ and $90^{\\circ}$. The frequent misalignment could result from jet precession or from deflection of the jet by interaction with gas in the interstellar medium.\\\\ ",
        "introduction": "Active galaxies are conventionally divided into radio-quiet and radio-loud categories. \\cite{1989AJ.....98.1195K} found that quasars separated into these two classes when the ratio of their radio to optical continuum was considered, and a similar result has been found by most (e.g. \\citealt{1986MNRAS.218..265P}, \\citealt{1990MNRAS.244..207M}, \\citealt{2002AJ....124.2364I}) but not all (\\citealt{2000ApJS..126..133W}) workers since. Members of the radio-loud category include BL~Lac objects, radio galaxies and radio-loud quasars. Radio-loud objects have large-scale radio lobes and compact luminous cores which often exhibit apparent superluminal motion. Radio-quiet objects include Seyfert galaxies and radio-quiet quasars; most of their radio emission is confined to the sub-kpc scale.  Recent Very Long Baseline Interferometry (VLBI) has shown that many Seyferts also contain compact nuclear radio components with brightness temperatures $\\gg10^7~{\\rm K}$ despite having radio luminosities over a factor of $10^8$ less than the most powerful active galaxies (e.g. \\citealt{1998ApJ...496..196U}; \\citealt{1999ApJ...516..127U}; \\citealt{1999ApJ...517L..81U}; \\citealt{2000evn..proc....7R}; \\citealt{2000ApJ...529..816M}), and that they harbour compact components moving at non-relativistic speeds. In all cases, the central masses are inferred to be black holes of millions to billions of solar masses, so the central potentials are relativistically deep. Since most accretion energy is released close to the hole, jets might be expected to be relativistic in all cases. Why then are the radio sources in Seyferts small, weak and slow compared to those in radio-loud objects?  Various answers to this question have been proposed, involving either ``intrinsic'' differences in the central engine or ``extrinsic'' differences in the surrounding medium. Intrinsic differences that have been suggested include (i) systematically lower black hole masses in Seyferts (\\citealt{2000ApJ...543L.111L}), (ii) lower black hole spins (\\citealt{1995ApJ...438...62W}), (iii) a ``magnetic switch'' that was identified by \\cite{1997Natur.388..350M} during numerical modeling of jets, (iv) the production in Seyferts of buoyant plasmons that bubble up through the density gradient of the narrow line region (NLR) instead of a collimated relativistic jet (\\citealt{1985MNRAS.214..463P}, \\citealt{1986MNRAS.222..189W}, \\citealt{1989MNRAS.240..487T}), (v) a large thermal plasma fraction in the jet in Seyferts (\\citealt{1998ApJ...495..680B}), or (vi) radiative inefficiency (\\citealt{Falcke1995}). Extrinsic differences generally invoke the rapid deceleration of initially relativistic jets by collisions in a dense surrounding BLR or ISM (e.g. \\citealt{1984AaA...141...85N}). With so many possible causes, and so few observational constraints, the question remains open why Seyfert galaxies have such low radio luminosities.  One approach to investigate whether the radio quietness of Seyferts is due to intrinsic or extrinsic causes is to measure speeds of radio components as close to the nucleus as possible, hopefully before interactions with the ISM affect them much. How close that needs to be is unclear, but signs of jet-NLR interaction have been seen on scales of tens of pc (\\citealt{1986MNRAS.222..189W}; \\citealt{1994AJ....107.1227W}; \\citealt{1995ApJ...454..106B}; \\citealt{1997ApJ...487..560C}; \\citealt{1998ApJ...502..199F}), so sub-pc scales are desirable. Proper motions of Seyfert galaxy radio components have now been measured with VLBI in a few cases (e.g. \\object{Mrk~231} and \\object{Mrk~348} by \\citealt{1999ApJ...517L..81U}, \\object{NGC~1068} by \\citealt{2000evn..proc....7R}, and \\object{NGC~4151} by \\citealt{1998ApJ...496..196U}, \\citealt{2002AAS...201.4814U}), and the motions were found to be $\\le 0.25~c$. In \\object{III~Zw~2}, $v<0.04~c$ was seen between two barely-resolved components over eight months, followed by a sudden displacement at $v\\ge1.25~c$ over seven months (\\citealt{2000AaA...357L..45B}).  The observations presented here were made for three reasons: 1) to measure proper motions of radio components on pc scales.  2) to look for extended, low surface-brightness radio emission on the scale of the NLR.  In the event that slow proper motions were measured, the next question would be whether components were ejected slowly or whether they were ejcted at high, perhaps relativistic velocities and then slowed by interaction with the NLR gas.  Were such braking to occur, one would expect shocks in the jet and possibly jet disruption. The in-situ particle acceleration associated with shocks could lead then to extended regions of radio emission.  3) to measure spectral indices, which provide a useful diagnostic of emission mechanisms, source compactness, and possible foreground free-free absorption. Most spectral index measurements of Seyfert galaxies to date have used the VLA, ATCA or WSRT, whose relatively large beams can average together many components in the nucleus. Measurements with parsec-scale resolution in Seyfert galaxies enable one to separate a possible compact, flat-spectrum component from surrounding steep-spectrum, optically-thin synchrotron emission, but such measurements require VLBI with matched beams at two frequencies and are still relatively few. In contrast, powerful radio sources have been well studied with VLBI and flat-spectrum core components are often seen at the base of a conically-expanding jet, thus indicating the location of the nucleus.  Do Seyfert nuclei also show a single, compact, synchrotron self-absorbed core component at the base of a collimated jet?\\\\  To meet these goals, we observed using Global VLBI, the European VLBI Network (EVN), the Multi-Element Radio Linked Interferometer Network (MERLIN) and the Very Long Baseline Array (VLBA) the Seyfert galaxies \\object{NGC~5506} and \\object{NGC~7674} at multiple epochs for proper motion measurements of nuclear radio sources, and \\object{Mrk~1210} and \\object{NGC~2110} for the first epoch.  To enlarge the sample size for deriving statistics on the spectral and geometric properties of the pc scale structure of Seyferts, we made a complete literature review of VLBI observations of Seyfert galaxies. We assume $H_0 = 75~{\\rm km\\,s^{-1}\\,Mpc^{-1}}$ and $q_0$ = 0.5 throughout this paper, and calculate spectral indices using $S\\propto\\nu^{\\alpha}$. ",
        "conclusions": "We have observed the Seyfert galaxies \\object{NGC~5506}, \\object{NGC~7674}, \\object{NGC~2110} and \\object{Mrk~1210} at 18~cm and 6~cm with the EVN and MERLIN.  \\object{NGC~5506} and \\object{Mrk~1210} were additionally observed with the VLBA.  These observations provided high angular resolution and good surface brightness sensitivity, spectral index measurements, and up to four epochs of observations for proper motion measurements.  Our conclusions are as follows.\\\\  1. The structure seen on the tens-of-parsec-scale in \\object{NGC~7674} shows two bright, steep-spectrum components that might be the working surfaces where double-sided jets ejected from an (unseen) AGN core located between them impact the interstellar medium. Low-surface-brightness emission was seen extending from the hot spots in the direction of VLA-observed emission by \\cite{Momjian2003}.  The diffuse emission could be from material that has flowed through the hotpots and was then either deflected or rose buoyantly in the pressure gradient in the ISM.  Alternatively, the bend after the hot spot might be caused by accretion disk precession, which might move the hot spots sideways along a wall of ISM leaving behind a trail of fading extended plasma.  These structures are reminiscent of those seen in compact symmetric objects, so we suggest that \\object{NGC~7674} is a low-luminosity member of the CSO family.  \\object{NGC~5506} shows two bright radio components separated by 3.5\\,pc with low-surface-brightness emission on larger scales extending towards the VLA-scale emission. The brightest compact component (B0) has a flat radio spectrum ($\\alpha=0.06$) and could represent the nucleus.  In \\object{NGC~2110} and \\object{Mrk~1210}, the overall structure is a triple radio source. For \\object{NGC~2110}, the flat spectrum ($\\alpha=0.12$) and the symmetric structure of the VLA-observed jets and the components at the ends of these jets (also seen in our MERLIN image) show conclusively that the central source represents the nucleus. In \\object{Mrk~1210}, both the SE and the NW components have steep radio spectra and here the location of the nucleus is unclear.  If jets in Seyferts are initially relativistic and are then disrupted by interaction with the ISM, we might expect significant faint, diffuse radio emission. Such faint emission (brighter than $2~{\\rm mJy\\,beam^{-1}}$) was seen in only \\object{NGC~5506}.\\\\  2. A proper motion of $(0.92\\pm0.07)~c$ was measured in \\object{NGC~7674}, and a $3~\\sigma$ upper limit of $0.50~c$ was found for the motion of B1 in \\object{NGC~5506} with respect to B0.  From the high speed measured in NGC 7674 between components 282~pc apart we derived a dynamical age of 1000~yr.  This speed and age is similar to those seen in CSOs.  That the components have high speed after propagating hundreds of parsecs shows that the ISM in this object has not stopped the jet very effectively.  The speed measured in \\object{NGC~7674} is faster than those measured in other Seyferts by \\cite{1998ApJ...496..196U}, \\cite{1999ApJ...517L..81U} and \\cite{2000evn..proc....7R}, except for IIIZw2, though is slow compared to the speeds at equivalent scales in powerful radio galaxies.  This suggests, but does not prove, that the jets are launched with an intrinsically lower speed than in radio galaxies.  Further support comes from the preference of Seyfert nuclei for early-type spirals, which have relatively low-density ISM even close to the nucleus, and hence it seems unlikely that the jets are significantly slowed.\\\\  3. We mentioned in the introduction the many possible reasons for low jet power in Seyferts that have been suggested in the literature.  Our observations of proper motions favour an intrinsic over an extrinsic cause since the ISM has not stopped the component in \\object{NGC~7674}, and no emission from disrupted jets was seen in three of four Seyferts. Any extrinsic deceleration must happen on scales less than 3.5~pc.  (However, the bends in the radio sources at distances of 3.5 pc and 282~pc in \\object{NGC~5506} and \\object{NGC~7674} show that the putative jets may nevertheless feel an influence from the NLR.)  Of the intrinsic causes mentioned in the introduction, we cannot constrain the role of black-hole spin or magnetic field strength in the accretion disc.  The expanding plasmon model of \\cite{1985MNRAS.214..463P} predicts that the radio components expand into an ISM of density $1~{\\rm cm^{-3}}$ at speeds initially of $\\ge 10^4~{\\rm km\\,s^{-1}}$ and would reach a size of 30~pc after 2000~yr by which time the expansion speed would have been reduced to $5000~{\\rm km\\,s^{-1}}$.  At this speed and at the distance of \\object{NGC~5506} over the 6.25~yr span of our observations, the component diameters would increase by 0.3~mas, or, in \\object{NGC~7674}, over 14.58~yr by 0.1 mas. An increase by this amount is, unfortunately, much smaller than could be measured with our data.\\\\  4. Spectral indices: We asked in the introduction whether Seyfert nuclei show a single compact synchrotron self-absorbed core component at base of a collimated jet, like scaled-down radio galaxies.  We found flat- or absorbed-spectrum components in \\object{NGC~5506} and \\object{NGC~2110} with brightness temperatures of $\\ge 3\\times10^7~{\\rm K}$.  The components in \\object{NGC~5506} have no clear axis of symmetry and the flat spectrum could be caused by either free-free absorption or synchrotron self-absorption.  The morphology of \\object{NGC~2110} shows collimated jets on VLA scales and so is the only object in our sample with a flat-spectrum VLBI core and large-scale jets.  Thus, the classical conically-expanding Blandford-K\\\"onigl jet structure (\\citealt{Blandford1979}, see also \\citealt{Falcke1995}) is sometimes but not often seen in our Seyfert galaxies.\\\\  5. Statistics of selected radio properties: To improve the reliability of any conclusions we enlarged our sample by drawing Seyferts from the literature that were observed with VLBI and found that flat- or absorbed- spectrum cores were present in 13 of 16 cases.  The most strongly absorbed components tend to be in objects that also show other evidence for absorption, such as water maser emission or large X-ray absorption columns.  In those cases, according to the Seyfert unification schemes, the flat-spectrum components would be viewed through an edge-on obscuring disk or torus, the ionized inner edge of which could be causing free-free absorption.  If so, the situation in those objects is different from the Blandford-K\\\"onigl jet model. However the case for free-free absorption is unambiguous in only one galaxy (NGC~3079) for which the spectral index is $>+2.5$.  The other 12 Seyferts have $-0.3<\\alpha<+2.5$ and so spectral turnovers could be produced by either synchrotron self-absorption or free-free absorption.  Where extended structures are present in the enlarged sample (14 Seyferts), the position angle differences between pc-scale and kpc-scale radio emission were found to be uniformly distributed between $0^{\\circ}$ and $90^{\\circ}$.  Such bends could be due to changes in the jet ejection axis or due to pressure gradients in the ISM. No correlation was found between the axis of pc-scale radio structure and the rotation axis of the host galaxy."
    },
    "0402/astro-ph0402374_arXiv.txt": {
        "abstract": "The source of the very bright Gamma-Ray Burst GRB 030329 is close enough to us for there to be a hope to measure or significantly constrain its putative superluminal motion. Such a phenomenon is expected in the ``Cannonball'' (CB) model of GRBs. Recent  precise data on the optical and radio afterglow of this GRB ---which demonstrated its very complex structure--- allow us to pin down the CB-model's prediction for the afterglow-source position as a function of time. It has been stated that (the unpublished part of) the new radio data ``unequivocably disprove'' the CB model. We show how greatly exaggerated that obituary announcement was, and how precise a refined analysis of the data would have to be, to be still of interest. ",
        "introduction": "The currently best-studied theories of Gamma-Ray Bursts (GRBs) and their afterglows (AGs) are the {\\it Fireball} models (see, e.g., Zhang \\& Meszaros 2003 for a recent review) and the {\\it Cannonball} (CB) model (see, e.g., Dar \\& De R\\'ujula 2003a; Dado, Dar \\& De R\\'ujula, 2002a; 2003a and references therein). The first set of models is often considered to be {\\it the standard model} of GRBs. In spite of their similarly-sounding names, these two models are (or were initially) completely different in their basic hypothesis, in their description of the data, and in their predictions. In this note we concentrate on a CB-model prediction which is not (to date) a standard-model one, the apparently superluminal motion of the source of GRBs and their afterglows: the ``cannonballs\". In quite exceptional cases ---relatively close-by GRBs with sufficiently bright (radio) AGs--- it may be possible to observe this superluminal motion of CBs relative to ``fixed stars''. Prior to GRB 030329, the case in which a possible superluminal motion (Dar \\& De R\\'ujula 2000) came closest to being observable was that of GRB 980425, which could ``almost'' be classified (Dar \\& De R\\'ujula 2003a) as an X-ray flash (XRF). For XRFs, the observation of a superluminal motion may be simpler than for GRBs, for the source's apparent displacement in the sky is proportional to the (small) observer's viewing angle, and we interpret XRFs as jetted GRBs viewed at larger angles (Dar \\& De R\\'ujula 2003a; Dado, Dar \\& De R\\'ujula 2003f). In a different model, this interpretation of XRFs has also been advocated by Yamazaki, Yonetoku \\& Nakamura (2003).  The CB-model is a very explicit elaboration of the original proposal by Shaviv and Dar (1995): that the $\\gamma$-rays of a GRB would be generated by inverse Compton Scattering (ICS) of stellar light by the electron constituents of transient narrow jets, emitted in stellar processes leading to gravitational collapse.   In the CB model, long-duration GRBs, as well as XRFs (Dar \\& De R\\'ujula 2003a; Dado, Dar \\& De R\\'ujula 2003f)  are produced in the explosions of {\\it ordinary} core-collapse SNe, akin to SN1998bw\\footnote{The X-ray and radio signals of the SN1998bw/GRB980425 pair are, in the CB model, attributed and well fit to the CB's AG, depriving the SN of these ``peculiar'' emissions. The observed large velocity of the SN's ejecta is attributed to their being observed exceptionally close to the jet axis. Neither this SN, nor its GRB, were exceptional (Dar \\& Plaga 1999, Dar \\& De R\\'ujula 2000, Dado et al.~2002a, 2003a).} (Dar \\& De R\\'ujula 2000), the first SN to be observed in ``association'' with a GRB (GRB 980425, Galama et al.~1998). Two opposite jets of CBs are emitted in the process, travelling with initially large Lorentz factors: $\\gamma_0\\sim 10^3$. The CBs initially expand (in their rest system) at a velocity comparable to, or smaller than, the speed of sound in a relativistic plasma ($c/\\sqrt{3}$), so that the jet opening angle (subtended by a CB's radius as observed from its emission point\\footnote{We are neglecting the initial CB's radius, presumably comparable or not much bigger than that of the collapsed core of the parent star, and thus entirely negligible by the time the GRB is emitted.}) is $\\alpha_j<\\! 1/(\\gamma_0\\,\\sqrt{3})$. The CBs' highly relativistic motion collimates their emitted radiation ---the GRB and its AG--- within a forward beam of characteristic opening angle $1/\\gamma$. An observer sees the ``Doppler-favoured'' jet, travelling at a small angle $\\theta={\\cal{O}}(1/\\gamma)$ relative to the line of sight. Typically $\\theta>\\alpha_j$, so that the jet's opening angle can be neglected and the observer's angle is the {\\it only} relevant one.  In the CB model, during the AG, the jet opening angle {\\it diminishes} with time, while the beaming angle $1/\\gamma(t)$ increases, both evolutions being due to the CB's interaction with the matter of the interstellar medium (Dado, Dar \\& De R\\'ujula, 2002a). As a consequence, the AG's source becomes increasingly ``pointlike'', and its motion in the sky can in principle be followed. Since a CB's motion is relativistic for days or months of (observer's) time, its apparent displacement in the sky is {\\it superluminal} (Courdec 1939, Rees 1967), as we argued in Dar \\& De R\\'ujula (2000).  The closest-by GRB observed so far ---GRB 980425, at a redshift $z=0.0085$--- came very close to having an observable superluminal motion (see Dado, Dar \\& De R\\'ujula 2003a for a detailed discussion). In a GCN note (Dar \\& De R\\'ujula 2003b) we argued that this motion may be observable in the next-closest GRB (030329, at $z=0.1685$).  In this note we sharpen the predictions for this putative motion, given the current availability of precise optical data at early times and times later than the first $\\sim 6$ days (Lipkin et al.~2003 and references therein), as well as sparse X-ray data (Marshall \\& Swank 2003;  Marshall, Markwardt \\& Swank 2003; Tiengo et al.~2003) and abundant radio data (Sheth et al.~2003; Berger et al.~2003; Pooley 2003; Kuno et al.~2004).   A very relevant new information (Lipkin et al.~2003 and references therein) on the AG of GRB 030329 is the abundance of multiple deviations of the optical light curves relative to a smoothly declining behaviour.  We shall refer to these deviations as {\\it features}. In Dado et al.~(2003c) we attributed the most obvious optical-AG feature ---a ``shoulder\" starting at $t\\sim$ 1 day (after burst)--- to a transition between a first to a second dominant CB, a choice supported by the fact that the $\\gamma$-ray light-curve of this GRB has a very clear two-pulse structure (Vanderspek et al.~2003; http:// space.mit.edu/ Hete/ Bursts/ GRB030329/; see also Vanderspeck et al.~2004), as shown in Fig.~(\\ref{GRB}). With the emergence of a handful of other similarly-significant features in the first week ---fast ups and downs of the optical fluences, $F_\\nu(t)$, by some 20 to 40\\%--- a more elaborate interpretation is required. In the AG model $F_\\nu(t)$ is a direct and {\\it quasi-local} tracer of the density of the ISM through which a CB travels, a fact to which we have attributed previous similar observations, e.g.~the ``humps\" in the optical AGs of GRBs 000301c and 970508 (Dado et al.~2002a).  It is therefore necessary to investigate the effect of local ISM density-inhomogeneities on the expected superluminal motion. This is what we do in this note for the case of GRB 030329. We conclude that the sources' motion results in a displacement of $\\sim 0.3$ (0.6) mas from day 3 to day 30 (100) after burst. Detecting such a motion may not be out of the question.   For the benefit of readers not familiar with the CB model, we present in an appendix a brief overview of the model and its current confrontation with data. We also offer there some commentary on the evolution of the standard model. ",
        "conclusions": "The data on GRB 030329 are now sufficiently complete to allow for a detailed prediction of the motion of the source of its AG ---allegedly superluminal in the CB model. The $\\gamma$-ray light-curve and the optical AG require the presence of two CBs, one of which dominates the AG at late times. The parameters needed to predict the motion of the two CBs in the sky are determined by the optical data, so that the individual motion of each CB is predicted. The main result of this paper is the prediction of the sky-motion of the fastest-moving CB, which dominates the late AG, and is shown in the lower panel of Fig.~(\\ref{tororo}).  The main caveat concerning a putative superluminal signature (Dar \\& De R\\'ujula 2000, 2003b) concerns the input to the estimate of its magnitude. Indeed, we have shown in this paper that the predictions are very sensitive to the details of the density profile of the ISM. But we are encouraged by the fact that in the only case in which the superluminal jets of CBs made by a core-collapse SN could be seen, they were seen. Indeed, observations of SN 1987A (Ninenson and Papaliolios, 1999) showed two sources, moving in opposite directions along the SN's axis at (real) velocities compatible with the speed of light and at an apparent superluminal velocity for the approaching source. Mercifully, the jet of that SN was not pointing in our direction (Dar, Laor \\& Shaviv 1998; Dar \\& De R\\'ujula 2001b).  Regarding the search for a superluminal motion, we learned by reading the e-version of NYT 030529 (the New York Times of that date, in GRB's parlance) that, according to Dale Frail {\\it ``[Our observations] are sufficient to rule out predictions of the cannonball model\"}. We have shown that, indeed, the observations of complicated features in the optical AG of GRB 030329 imply that our earlier results (Dar \\& De R\\'ujula 2003b) ---which ignored the presence of these features--- constituted an overestimate of the predicted superluminal displacement. In this sense, Frail was right in stating that the observations ruled out the {\\it predictions}, as opposed to the model itself.  In a setting more scientific than the NYT, Bloom et al.~(2003) state: {\\it ``Owing to the proximity and bright radio emission, high-resolution ($\\sim1$ pc) {\\it Very Long Baseline Array} imaging of the compact afterglow was used by Frail (2003) to} {\\bf unequivocally disprove} {\\it the cannonball model for the origin of GRBs.\"} The emphasis is ours. We have seen that these news of the death of the CB model may have been premature. Even though Mark Twain eventually died for sure, the CB model ---though probably not immortal--- is still in an excellent shape. Yet, trying to disprove the best available model(s), or even the proof of a difficult theorem, is the acceptable standard attitude in many a realm of the exact sciences. In this sense, the apparently strong motivation of the quoted observers to disprove the CB model is ---in our opinion, and regarding this particular model--- the healthiest of all possible attitudes.     \\noindent {\\bf Acknowledgements.} One of us (A.~De R.) is indebted to the Physics Department and Space Research Institute of Technion for its hospitality. This research was supported in part by the Helen Asher Space Research Fund for research at the Technion."
    },
    "0402/astro-ph0402004_arXiv.txt": {
        "abstract": "We investigate what the lensing information contained in high resolution, low noise CMB temperature maps can teach us about cluster mass profiles.  We create lensing fields and Sunyaev-Zel'dovich effect maps from N-body simulations and apply them to primary CMB anisotropies modeled as a Gaussian random field.  We examine the success of several techniques of cluster mass reconstruction using CMB lensing information, and make an estimate of the observational requirements necessary to achieve a satisfactory result. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0402/astro-ph0402010_arXiv.txt": {
        "abstract": "{ We present here new spectroscopic observations  of  Mrk 1040 and LEDA 212995 (Mrk 1040 companion) obtained with the Isaac Newton Telescope (INT). The intensity ratios and widths for the narrow emission lines found in LEDA 212995 are typical of H II regions.  The red-shift (0.0169$\\pm$0.00015) of the object derived from these emission lines is  very close to the red-shift of Mrk 1040 (z=0.01665). The weak narrow and broad  absorption lines were detected in the H$\\alpha$ wavelength band of LEDA 212995 spectra. These absorptions indicate that the companion might be at least partly obscured by Mrk 1040. Using this and previous observations we discuss the possible physical relationship between these two galaxies. ",
        "introduction": "LEDA 212995\\footnote{In NED database this object is noted also  as UGC 01935 NOTES01 and RX J0228.2+3118, in literature  is often noted as 'Mrk 1040 companion'} is a small galaxy with dimensions $0.20' \\times 0.10'$ and magnitude $>$19  visually located near the Sy 1 galaxy Mrk 1040 (z=0.01665, Huchra et al. 1999). The galactic extinction for this object was given by Burstein \\& Heiles (1982) and Schlegel et al. (1998). The other conversion factors for this galaxy were given by Cardelli et al. (1989). \\begin{figure*} \\includegraphics[width=16cm]{fig1a.ps} \\includegraphics[width=16cm]{fig1b.ps} \\caption{ The comparison of the Mrk 1040 and LEDA 212995 spectra in the H$\\beta$  (top) and in the H$\\alpha$ (bottom) wavelength range. } \\end{figure*}   \\begin{figure} \\includegraphics[width=8.5cm]{fig2a.ps} \\caption{Decomposition of H$\\alpha$  line region of LEDA 212995. The 'absorption deep' can be fitted with three broad absorption (at the bottom) which correspond to the H$\\alpha$+[NII] lines redshifted $z\\approx0.0160$. The narrow absorption (solid line) are H$\\alpha$ and [NII]$\\lambda$6583.6 \\AA\\ with the red-shift of Mrk 1040 ($z\\approx0.01665$). } \\end{figure}  \\begin{figure} \\includegraphics[width=8.5cm]{fig2b.ps} \\caption{Decomposition of  the H$\\beta$ line region of LEDA 212995.} \\end{figure}  \\begin{figure} \\includegraphics[width=8.5cm]{fig3.ps} \\caption{Decomposition of Mrk 1040 H$\\alpha$ line. The dots represent observation, and the solid line is the best fit. The Gaussian components are shown at the bottom. The dashed lines at the bottom represent the narrow  lines.} \\end{figure}  \\begin{figure} \\includegraphics[width=8.5cm]{fig4.ps} \\caption{The same as in Fig. 4, but for H$\\beta$ line. The broad dashed line at the bottom represents Fe II template.} \\end{figure}   Previous spectroscopical observations  of this object were reported by Ward \\& Wilson (1978), Dahari (1985), Veilleux \\& Osterbrock (1987), Amram et al. (1992) and Keel (1996). Ward \\& Wilson (1987) found that the galaxy emits narrow and 'sharp' emission lines with red-shift 5070$\\pm$60 km/s (0.01691$\\pm$0.0002) which is close to the Mrk 1040 red-shift (4910$\\pm$60 km/s, i.e. 0.01638$\\pm$0.0002  according to the authors). Dahari (1985) in Table 1, summarized the red-shifts of companions of Seyfert galaxies where the red-shift of 0.0163 for this object was given. Also, the spectroscopical observation of this object with CCD transmission-grism spectrograf at the Cassegrain focus of the Shane 3m telescope was analyzed  by Veilleux \\& Osterbrock (1987). Using the emission line ratio they classified the galaxy as Narrow Emission Line Galaxy (NLRG). Veilleux \\& Osterbrock (1987) found that systemic red-shift of the object is 0.0160  with the uncertainties from 0.0002 to 0.0005. Assuming this value of the red-shift, the companion should be closer to the observer. Moreover, Afanas'ev \\& Fridman (1993) pointed out that an analysis of the (B-R) color distribution in the galactic disk and the presence of a distinct dust lane in the disk show that 'the northeast side of the galaxy is farther away, and the companion, which is bluer than the disk of Mrk1040, is closer to the observer'. The spectroscopical observations of  Mrk1040 and LEDA 212995 were given in Amram et al. (1992), where the asymmetry in velocity field of companion is found and assumed that this asymmetry is due to interaction of these two galaxies. They found that red-shift of LEDA 212995 is 5100$\\pm$ 10 km/s that is in agreement with measurements by Ward \\& Wilson (1978). Keel (1996) presented imaging and optical spectroscopy of a sample paired Seyfert galaxies and found that LEDA 212995 is about 18 arc seconds to the north along the Mrk 1040  minor axis.  In this paper  the  spectroscopic characteristics (red-shift, emission lines) of the object were not given.  According to the red-shifts estimates of different authors,  it is possible to conclude that  galaxies, Mrk 1040 and its companion, are physically related which can be of interest in what respect to the origin of the nuclear activity in  Mrk 1040 (see e.g. Corbin 2000).   We present here new observations of Mrk 1040 and its companion obtained with the Isaac Newton Telescope (INT) in the H$\\beta$ and H$\\alpha$ spectral regions with the aim of characterize the line emission in Mrk 1040 companion and discuss the physical relationship between these galaxies. ",
        "conclusions": "We present our analysis of LEDA 212995 spectra in H$\\beta$ and H$\\alpha$ wavelength regions. The spectra of this object and Mrk 1040 were observed simultaneously (both object were in the slit) with INT. From this analysis we found that:  i) the intensity ratio of narrow emission lines and their widths ($\\approx$ 50 km/s) in the spectra  of LEDA 212995 are  typical for H II regions;  ii) the red-shift inferred from the narrow emission  lines is $z=0.0169\\pm 0.0015$ that is in agreement with Ward \\& Wilson (1978) and Amram et al. (1992).  This red-shift imply that LEDA 212995, as was earlier commented, is really a neighboring galaxy to the Sy 1 galaxy Mrk 1040.  iii) The 'absorption deep' seen in H$\\alpha$ wavelength region indicates that LEDA 212995 is at least partly obscured by  Mrk 1040, it means that  Mrk 1040 may be a foreground galaxy of LEDA 212995, or in another case LEDA 212995 might be very close to Mrk 1040 nuclei. To ascertain this question, further spectroscopic observations  of the Mrk 1040 companion and north-east part of the Mrk 1040 stellar disk are needed."
    },
    "0402/astro-ph0402360_arXiv.txt": {
        "abstract": "Intermediate degree modes of the solar oscillations have previously been used to determine the solar helium abundance to a high degree of precision. However, we cannot expect to observe such modes in other stars. In this work we investigate whether low degree modes that should be available from space-based asteroseismology missions can be used to determine the helium abundance, $Y$,  in stellar envelopes with sufficient precision. We find that the oscillatory signal in the frequencies caused by the depression in $\\Gamma_1$ in the second helium ionisation zone can be used  to determine the envelope helium abundance of low mass main sequence stars. For frequency errors of 1 part in $10^4$, we expect errors $\\sigma_Y$ in the estimated helium abundance to range from $0.03$ for 0.8M$_\\odot$ stars to  $0.01$ for 1.2M$_\\odot$ stars.  The task is more complicated in evolved stars, such as subgiants, but is still feasible if the relative errors in the frequencies are less than $10^{-4}$. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0402/astro-ph0402156_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:1}  Due to the atmospheric opacity, the far--IR domain has been the last window used in Astronomy. For the first time, the \\textit{Infrared Space Observatory} (ISO) has opened this spectral frequency range through \\textit{molecular spectroscopy}. The sensitivity of the instrumentation on board this platform  has no comparison with the few previous space missions or airborne observations carried before the launch of ISO. Almost all the operative range of ISO in the far--IR was not explored before. The far--IR spectrum of the most significant galactic sources was unknown. In particular, the main radiation emitters, the molecules, were unidentified. Far--IR observations are specially suitable to the study of the \\textit{warm gas} in molecular clouds. Among these sources, Sgr~B2 in the Galactic Center, is a  paradigmatic object for our knowledge of the \\textit{chemical complexity} of the Galaxy.  The molecular species and the atomic fine structure lines that can be detected in the far--IR domain are essential for the study of the physical and chemical conditions of the interstellar medium. The bulk of these species can only be observed from space platforms. As a example, the water vapor abundance can determine if stars will be formed during the gravitational collapse of a molecular cloud. Another example are the non--polar  carbon chains. These species can be the \\textit{``skeletons''} from which the large carbon molecules responsible of a great part of the  IR emission in the Galaxy can be formed. Due to the lack of permanent electric dipole, these species do not have rotational spectrum to be observed from radio telescopes.  \\clearpage  In this contribution we present the main results of our detailed study of the \\textit{Long--wavelength spectrometer} (LWS)  spectrum of Sgr~B2(M) between 43~$\\mu$m (7.0~THz) and 197~$\\mu$m (1.5~THz). Both with the grating ($\\lambda/\\Delta\\lambda$$\\sim$200) and with the  Fabry--Perot (FP; $\\lambda/\\Delta\\lambda$$\\sim$10000). ",
        "conclusions": ""
    },
    "0402/astro-ph0402406_arXiv.txt": {
        "abstract": "We present the optical identification of mid-IR and radio sources detected in the European Large Area ISO Survey (\\ELAIS) areas N1 and N2. Using the r' band optical data from the Wide Field Survey we apply a likelihood ratio method to search for the counterparts of the 1056 and 691 sources detected at 15$\\mu$m and 1.4\\,GHz respectively, down to flux limits of $S_{15}=0.5$\\,mJy and $S_{1.4\\,{\\rm GHz}}=0.135$\\,mJy.  We find that $\\sim$92\\% of the 15$\\mu$m \\ELAIS sources have an optical counterpart down to the magnitude limit of the optical data, r'=24.  All mid-IR sources with fluxes $S_{15}\\geq3$\\,mJy have an optical counterpart.  The magnitude distribution of the sources shows a well defined peak at relatively bright magnitudes r'$\\sim$18. About 20\\% of the identified sources show a point-like morphology; its magnitude distribution has a peak at fainter magnitudes than those of galaxies.  The mid-IR-to-optical and radio-to-optical flux diagrams are presented and discussed in terms of actual galaxy models. Objects with mid-IR-to-optical fluxes larger than 1000 are found that can only be explained as highly obscured star forming galaxies or AGNs. Blank fields being 8\\% of the 15$\\mu$m sample have even larger ratios suggesting that they may be associated with higher redshift and higher obscured objects. ",
        "introduction": "The Infrared Space Observatory \\citep[ISO, ][]{1996A&A...315L..27K} was the second infrared space mission, providing a great improvement in sensitivity over the IRAS mission.  The European Large-Area ISO survey \\citep[\\ELAIS, ][]{2000MNRAS.316..749O} was the largest Open Time programme on ISO.  This project surveyed 12 square degrees, divided into four main fields, three in the north (N1, N2, N3) and one in the south (S1). The main survey bands used were 6.7, 15, 90 and 170\\micron. The ISOCAM camera \\citep{1996A&A...315L..32C} was used for the shorter wavelengths while the ISOPHOT \\citep{1996A&A...315L..64L} camera was used for the longer ones.      Optical imaging is essential to study the properties of the sources detected. Due to the large errors ellipses of the mid-IR detections, typically several seconds of arc, it is necessary to carry out a detailed process of identification. Since more than one optical source can be inside those ellipses, a method which provides the likelihood of each counterpart to be the true association is needed. The identification does not only provides us the optical properties of the mid-IR sources but also allows us to further proceed with followup observations of interesting sources.  This paper presents the optical identification of the mid-IR and radio sources in the N1 and N2 areas. Section~\\ref{sec:wfs} presents a summary of the optical observations carried out in these areas as well as the reduction steps and products. Section~\\ref{sec:elaiscat} describes the mid-IR and radio catalogues used.  Section~\\ref{sec:id} discusses the actual procedure to determine the optical counterparts of the sources, while sections~\\ref{sec:properties} and~\\ref{sec:properties2} describe the optical properties of the sources. ",
        "conclusions": "The association of sources detected at 15$\\mu$m in the \\ELAIS N1 and N2 areas with optical objects is presented. A 92\\% of the sample presents an optical identification to r'=24. The magnitude distribution presents a maximum at r'=18 and a tail which extends to fainter magnitudes. The distribution of point-like objects presents a maximum one magnitude fainter.  The mid-IR to optical flux ratios, $S_{15}/S_{r'}$, of the bright optical sources are in the range 1 to $10^2$ and can be explained using simple models of cirrus, starbursts, AGN and Arp220 spectral energy distributions. The tail of faint objects, show larger $S_{15}/S_{r'}$, from $10^2$ to $10^3$ and can only be explained assuming large luminosities and obscurations. Point like objects show bluer g'-r' colour while higher mid-IR to optical flux ratio than galaxies. The remaining 8\\% of objects not identified are faint in the mid-IR , with a $15\\mu$m flux lower than 3\\,mJy. However, their mid-IR-to-optical flux ratio is larger than $10^3$ favouring the interpretation that they are associated with starbursts or AGNs at high redshifts and highly obscured. The identification of radio sources in the same areas is also presented. Their magnitude distribution shows an increase on the number of sources towards faint magnitudes. The number of unidentified objects is 44\\%."
    },
    "0402/astro-ph0402540_arXiv.txt": {
        "abstract": "The atmospheres of weakly-magnetized neutron stars expand hydrostatically and rotate differentially during thermonuclear X-ray bursts.  Differential rotation is probably related to the frequency drifts of millisecond burst oscillations exhibited by about a dozen nuclear-powered X-ray pulsars.  Here, we analyze the linear stability of this differential rotation with respect to local, axisymmetric, multi-diffusive MHD perturbations, at various heights in the neutron star atmosphere.  Unstable magneto-rotational modes are identified from within to well above the burning layers.  Properties of the fastest growing modes depend sensitively on the local magnetic field geometry. Linear estimates suggest that momentum transport due to magneto-rotational instabilities can affect atmospheric rotation profiles on time-scales relevant to burst oscillations. This transport would likely strengthen the coherence of burst oscillations and contribute to their drifts. ",
        "introduction": "Since their discovery in the 1970s (see Lewin, van Paradijs \\& Taam 1995 for a review), thermonuclear X-ray bursts have been observed from a large number of low-mass X-ray binaries containing weakly-magnetized accreting neutron stars. The basic physical mechanism responsible for X-ray bursts, namely mass accumulation on the neutron star surface at rates such that a thermonuclear He or mixed H/He flash is eventually triggered, is well understood (see, e.g., Bildsten 1998 and references therein). The phenomenology of X-ray bursts is rich, however, and theoretical models face some difficulties when detailed comparisons with the observations are attempted (e.g. Bildsten 2000; Galloway et al. 2003).   This situation has become more serious, and more exciting, in the past decade, with the discovery of transient millisecond ($300$--$600$~Hz) oscillations during some X-ray bursts, the so-called burst oscillations (Strohmayer et al. 1996). While the presence of oscillations during burst rise may be understood as resulting from the rotational modulation of a nuclear burning spot spreading around the stellar surface (Spitkovsky, Levin \\& Ushomirsky 2002), it is unclear why oscillations should be visible past the burst peak, when the burning fuel has presumably been entirely ignited. Also intriguing is the nature of frequency drifts exhibited by most burst oscillations (generally spin-ups, by a few Hz at most) and the origin of the diversity they show in terms of drift time-scales and amplitudes (see Strohmayer \\& Bildsten 2003 for a review).  Observations of the burst oscillation phenomenon have been very rewarding. In particular, it has recently been shown that the burst oscillation asymptotic frequency (after complete spin-up) and the independently-known neutron star spin frequency are closely related in two millisecond accreting pulsars (Chakrabarty et al. 2003; Strohmayer et al. 2003). Still, our understanding of the physical mechanisms involved in burst oscillations and associated frequency drifts remains limited.  Following Strohmayer et al. (1997), Cumming \\& Bildsten (2000) proposed that frequency drifts originate from the rotational evolution of heated burning layers, as they cool, hydrostatically contract and nearly recover their original rotation, assuming their specific angular momentum is conserved at all times. These same authors later realized, however, that the largest observed frequency drifts are in excess of the amplitudes expected in their burning layer scenario, by as much as a factor 2-3. Apparently, only layers higher up in the atmosphere would have the right amount of rotational offset to explain the largest drift amplitudes (Cumming et al. 2002). More recently, Chakrabarty et al. (2003) have also argued that the very rapid drift (spin-up) that they have observed during the rise phase of one of the bursts of SAX J1808-3658 effectively rules out the scenario of Cumming \\& Bildsten (2000). There is currently no solid alternative explanation for the origin of burst oscillations and associated frequency drifts, even though both Rossby wave (Heyl 2003; but see Lee 2003) and zonal flow (Spitkovsky et al. 2002) scenarios have been proposed.    In their study of neutron star atmospheres, Cumming \\& Bildsten (2000) have analyzed a number of mechanisms which could make the vertical differential rotation established via burst-induced hydrostatic expansion unstable. They have not identified any clear destabilizing process (which would presumably lead to angular momentum transport between atmospheric shells) and this was the main motivation for their assumption of conserved specific angular momentum.  The purpose of the present study is to reconsider the issue of stability of differential rotation in the atmospheres of weakly-magnetized, rapidly-rotating neutron stars.   As noted by Cumming \\& Bildsten (2000), the strong vertical thermal stratification of a neutron star atmosphere stabilizes it with respect to adiabatic perturbations, even in the presence of strong vertical differential rotation. Although not considered by these authors in their analysis, one indeed verifies that the vertical differential rotation expected in the present context is stable according to both the hydrodynamical Solberg-H\\o iland adiabatic criteria (e.g. Tassoul 1978) and their generalization for weakly-magnetized fluids (Balbus 1995).    This situation is reminiscent of the solar radiative interior, where thermal stratification also plays a strongly stabilizing role. Goldreich \\& Schubert (1967) and Fricke (1968) have shown, however, that rotational instabilities can still exist, provided perturbed fluid elements are allowed to exchange heat with their environment much faster than they exchange momentum. The stabilizing role of the atmospheric thermal stratification is then effectively neutralized when a perturbed fluid element reaches thermal equilibrium with its environment while its original momentum remains largely unchanged. In the solar interior, radiative heat diffusion is indeed orders of magnitude faster than viscous diffusion of momentum, and this was the initial motivation for the double-diffusive analysis of Goldreich \\& Schubert and Fricke (hereafter GSF altogether). The basic double-diffusive mechanism invoked in these instabilities is not too different from that of ``salt-finger'' instabilities operating in the Earth's oceans, except that the destabilizing salinity stratification is replaced by a destabilizing angular momentum stratification in the stellar context.     As we shall see below, conditions in the atmospheres of weakly-magnetized neutron stars during bursts are also strongly multi-diffusive. It is then possible that rotational instabilities will develop and affect the differentially-rotating atmospheric layers in such a way as to influence the observational properties of burst oscillations. To address this question properly, it is necessary to account for the presence of magnetic fields in neutron star atmospheres. Recently, Menou, Balbus \\& Spruit (2004) have generalized the work of GSF to weakly-magnetized fluids. They have shown that the combination of weak magnetic fields and a multi-diffusive situation may have contributed to establishing the state of near solid-body rotation inferred from helioseismology for the solar interior.    In this paper, we present a specific application of the stability analysis of Menou et al. (2004) to weakly-magnetized, differentially-rotating neutron star atmospheres and we discuss possible consequences for the phenomenology of burst oscillations and associated frequency drifts. We describe our method of solution in \\S2. In \\S3, we present the results of our stability analysis at various heights in the neutron star atmosphere. The relevance of these results to burst oscillations is discussed in \\S4. ",
        "conclusions": "The atmospheres of weakly-magnetized neutron stars are subject to vertical differential rotation during thermonuclear X-ray bursts. Conditions in these atmospheres are also strongly multi-diffusive and thus favor the development of magneto-rotational instabilities. We have presented a linear stability analysis of the expected differential rotation, with respect to local axisymmetric MHD perturbations, in three representative atmospheric layers. Unstable diffusive magneto-rotational modes were identified in much of the differentially-rotating atmosphere. This indicates that momentum transport between the various atmospheric layers could occur.   Our main motivation to perform a stability analysis is the possibility that the transport resulting from instabilities may have observational consequences for millisecond burst oscillations and associated frequency drifts seen so far in a dozen accreting neutron star systems (nuclear-powered X-ray pulsars; e.g. Strohmayer \\& Bildsten 2003). The differential rotation established during thermonuclear X-ray bursts is transient, lasting perhaps several tens of seconds (a typical burst duration). Therefore, it is not sufficient to identify magneto-rotational instabilities in this context. For these instabilities to have any effect, they must act to modify atmospheric rotation profiles on short enough timescales. We have attempted to address this issue by using linear estimates of the efficiency of momentum transport and have concluded that, indeed, these instabilities could be efficient enough to be of relevance to the bust oscillation phenomenon.    There are reasons to believe that diffusive magneto-rotational instabilities will, in a stellar context, tend to reduce the level of existing differential rotation and perhaps make the system approach a state of solid body rotation (see discussion in Menou et al. 2004; see also Korycansky 1991). If correct, this conjecture would have interesting consequences for the coherence of burst oscillations. Indeed, it was noted by Cumming \\& Bildsten (2000) that the strong coherence of some burst oscillations implies, in the absence of any momentum transport between atmospheric shells, that the feature at the origin of the oscillations must be confined to a surprisingly thin layer (much thinner than a local scale height in their burning layer scenario). This feature would otherwise be smeared out by differential rotation and the oscillations would lose their coherence.   If, however, magneto-rotational instabilities act to reduce the amount of differential rotation present during bursts, as we have suggested, they would effectively promote the coherence of burst oscillations. Constraints on the thickness of the layer from which the modulated signal originates would then be relaxed. Magneto-rotational instabilities would also contribute to the observed frequency drifts of burst oscillations if they act efficiently enough to modify the rotational evolution of atmospheric layers. This possibility is consistent with the results of Cumming et al. (2002), who concluded that conservation of specific angular momentum cannot account for the largest observed frequency drifts.   As attractive as it is, this scenario suffers from a number of significant uncertainties. For instance, we have cautioned about the risks of estimating efficiencies of momentum transport from linear mode properties only.  Also, the feature which is at the origin of burst oscillations must be non-axisymmetric. The assumption of axisymmetry made in our analysis could therefore be a limitation. We have illustrated how the properties of fastest growing magneto-rotational modes strongly depend on the local magnetic field geometry. This may be viewed as a model uncertainty because the global field topology on a given neutron star is a priori unknown. On the other hand, this sensitivity to magnetic field geometry might in principle be related to the link recently established between the phases of persistent and burst oscillations, because it suggests a special role for magnetic polar caps in the burst oscillation phenomenon (Chakrabarty et al. 2003; Strohmayer et al. 2003). One then wonders whether the diversity in drift timescales could be attributed to different field geometries in different systems, perhaps caused by varying degrees of field burying under continuous accretion (Cumming, Zweibel \\& Bildsten 2001).   Another possible limitation of our work comes from neglecting the stabilizing role of composition gradients. In the vicinity of burning layers, substantial composition gradients may be expected in ashes, following the main episode of nuclear burning. The stratification due to composition gradients is important because, contrary to thermal stratification, it is not sensitive to a fast rate of heat transfer. Although it is difficult to assess the importance of stabilization due to composition gradients without a detailed study, the specific example described by Goldreich \\& Schubert (1967), for a purely hydrodynamical situation, suggests that this effect should be included when discussing neutron star atmosphere during bursts. Another motivation to elucidate the role of stabilizing composition gradients is the possibility that non-axisymmetric versions of diffusive magneto-rotational modes could contribute to the friction term entering the burning front models of Spitkovsky et al. (2002).    Rather than focusing on burning layers only, we have performed stability analyses at various heights in the differentially-rotating neutron star atmosphere. One motivation for this is the stabilizing effect of composition gradients that we have just discussed. Even if such stabilization occurs in the vicinity of burning layers, magneto-rotational instabilities may still be able to operate higher-up in the atmosphere. A second motivation is the indication, from the work of Cumming et al. (2002; see, e.g., their figure~4), that layers above the burning regions are involved in the burst oscillation phenomenon, since they are the only ones having the right amount of rotational offset to explain the largest observed frequency drifts.    Our analysis for high atmospheric layers shows that magneto-rotational instabilities cease gradually to exist as one approaches the stellar photosphere, because the weak magnetic tension they require becomes increasingly difficult to satisfy. Still, according to our middle-atmospheric models, there is a region above the burning layers in which magneto-rotational instabilities can operate and possibly lead to a rather efficient transport of momentum. While burst oscillations have traditionally been associated with burning layers, it may be worth giving more attention to these middle-atmospheric layers. Not only do they have the right amount of rotational offset to explain large frequency drifts, but these layers are also more advantageous for the visibility of oscillations, with a short heat diffusion time compared to the time they take to revolve around the star (see Cumming \\& Bildsten 2000 for a discussion of the visibility of burst oscillations).   We note that Cumming \\& Bildsten (2000) and Cumming et al. (2002) have invoked the mechanism of magnetic field wind-up described by Spruit (1999) as a possible source of recoupling between differentially-rotating layers. In regions with weak enough stratification from composition gradients for magneto-rotational instabilities to operate, this process may not be relevant. Indeed, in the case of differentially-rotating accretion disks, it has been argued analytically and demonstrated numerically that the linear growth of the azimuthal field component caused by differential rotation is not important because of the exponential growth of magneto-rotational modes (Balbus \\& Hawley 1991; 1998). By analogy, one may expect the same linear growth to be irrelevant in the stellar context, given the existence of exponentially-growing diffusive magneto-rotational modes.   Beyond simple qualitative points, it is difficult to deepen an interpretation of burst oscillations and associated frequency drifts in terms of a rotational evolution which is partly driven by magneto-rotational instabilities.  A time-dependent model for the coupled thermal, chemical and rotational evolution of a neutron star atmosphere during burst seems to be required to make more definite predictions. Such time-dependent models, ignoring the rotational evolution, already exist and exhibit complex behaviors. For example, transient convection early in the burst leads to some mixing of momentum and elements, while inverted composition gradients (which would presumably promote instabilities) are sometimes obtained after the main nuclear burning phase (see, e.g., Woosley \\& Weaver 1984; Woosley et al. 2003). In order to include the rotational evolution, these models would require reliable estimates of the efficiency of momentum transport (and of the resulting chemical mixing), which could presumably be obtained from fully turbulent numerical simulations of the non-linear development of diffusive magneto-rotational instabilities.  Until such advanced tools are developed, it is likely that observations will remain our best guide for understanding burst oscillations and associated frequency drifts."
    },
    "0402/astro-ph0402589_arXiv.txt": {
        "abstract": "We discuss the origin of HE0107-5240 which, with a metallicity of $ \\feoh =-5.3$, is the most iron-poor star yet observed. Its discovery has an important bearing on the question of the observability of first generation stars in our Universe. In common with other stars of very small metallicity ($-4 \\lesssim \\feoh \\lesssim - 2.5$), HE0107-5240 shows a peculiar abundance pattern, including large enhancements of C, N, and O, and a more modest enhancement of Na. The observed abundance pattern can be explained by nucleosynthesis and mass transfer in a first generation binary star, which, after birth, accretes matter from a primordial cloud mixed with the ejectum of a supernova. We elaborate the binary scenario on the basis of our current understanding of the evolution and nucleosynthesis of extremely metal-poor, low-mass model stars and discuss the possibility of discriminating this scenario from others.  In our picture, iron-peak elements arise in surface layers of the component stars by accretion of gas from the polluted primordial cloud, pollution occurring after the birth of the binary. To explain the observed C, N, O, and Na enhancements as well as the $\\nucm{12}{C}/\\nucm{13}{C}$ ratio, we suppose that the currently observed star, once the secondary in a binary, accreted matter from a chemically evolved companion, which is now a white dwarf. To estimate the abundances in the matter transferred in the binary, we rely on the results of computations of model stars constructed with up-to-date input physics. Nucleosynthesis in a helium flash driven convective zone into which hydrogen has been injected is followed, allowing us to to explain the origin in the primary of the observed O and Na enrichments and to discuss the abundances of s-process elements.  From the observed abundances, we conclude that \\mmps\\ has evolved from a wide binary (of initial separation $\\sim 20$ AU) with a primary of initial mass in the range $1.2 \\sim 3 \\msun$. On the assumption that the system now consists of a white dwarf and a red giant, the present binary separation and period are estimated at $\\simeq 34$ AU and a period of $\\simeq 150$ years, respectively.  We also conclude that the abundance distribution of heavy s-process elements may hold the key to a satisfactory understanding of the origin of HE0107-5240. An enhancement of $[{\\rm Pb} / {\\rm Fe}] \\simeq 1 \\sim 2$ should be observed if \\mmps\\ is a second generation star, formed from gas already polluted with iron-group elements. If the enhancement of main-line s-process elements is not detected, \\mmps\\ may be a first generation secondary in a binary system with a primary of mass less than 2.5 $\\msun$, born from gas of primordial composition, produced in the Big Bang, and subsequently subjected to surface pollution by accretion of gas from the parent cloud metal-enriched by mixing with the ejectum of a supernova. ",
        "introduction": "The discovery of the exceedingly metal-poor red giant \\mmps\\ \\citep{chr02} has great importance for our understanding of early star formation in the Galaxy. It has the smallest metallicity ($\\feoh = -5.3$) of any star yet observed and it exhibits several abundance ratios which differ markedly from abundance ratios in solar system material. In particular, it shows large enhancements of carbon ($[\\nucm{}{C}/\\nucm{}{Fe}] = 4.0$), nitrogen ($[\\nucm{}{N}/\\nucm{}{Fe}] = 2.3$), and oxygen ($\\mbox{[O/Fe]}=2.4^{+0.2}_{-0.4}$) \\citep{bes04}, as well as a mild enhancement of sodium ($[\\nucm{}{Na}/\\nucm{}{Fe}] = 0.8$); currently, only an upper bound exists for the important s-process element Ba ($[\\nucm{}{Ba}/\\nucm{}{Fe}] < 0.8$) \\citep{chr02}. We use the traditional term ``metal poor'' in referring to \\mmps\\ even though, thanks to the large abundances of CNO elements relative to Fe, it is certainly not ``heavy element poor.''  Recent papers which discuss the origin of \\mmps\\ use arguments involving supernova nucleosynthesis in a primordial cloud \\citep{ume03,lim03}, external pollution after birth \\citep{shi03}, and the mass distribution predicted by theories of star formation \\citep{sal03, omu03}. \\citet{shi03} predict a metallicity distribution function for first generation (Pop.~III) stars currently burning hydrogen and conclude that \\mmps \\ is a first generation object with a surface affected by accreting interstellar matter polluted with heavy elements. \\citet{ume03} adjust parameters in a first generation supernova model in such a way as to produce a C/Fe ratio in the supernova ejectum that agrees with the ratio observed in \\mmps\\ and argue that \\mmps\\ is a second generation object formed from the primordial cloud after it has been mixed with the ejectum of the supernova. A similar scenario is presented by \\citet{lim03} who argue that the \\mmps\\ abundances can be produced by a combination of two types of supernova ejecta: a normal ejectum consisting of $\\sim 0.06 \\msun$ of iron and an abnormal ejectum consisting only of products of partial helium burning. Though all extant scenarios address important aspects of the problem, further discussion is warranted of the physics of star formation and of the chemical composition expected in a primordial cloud into which matter ejected by a supernova has been mixed. More importantly, the modifications of surface abundances which \\mmps\\ has suffered during its long life should be elucidated. In particular, the possibility of accretion from an evolved first generation companion which has mixed to its surface products of internal nucleosynthesis should be explored. In this paper, we describe results of such an exploration.  A major characteristic of models of extremely metal poor stars is that, although the p-p chain reactions are the dominant source of the stellar luminosity and the main driver of evolution during most of the core hydrogen-burning phase, the CNO cycles play an increasingly important role as evolution progresses beyond the main sequence phase. This is because, as temperatures increase, carbon is produced by the highly temperature-sensitive triple-$\\alpha$ reaction and because, at high temperatures, only a small abundance of CNO elements is needed for CNO cycle reactions to become the dominant driver of evolution. This characteristic behavior of the evolution of zero metallicity stars has been explored by many authors, from the early works of \\citet{wag74}, \\citet{dan82} and \\citet{gue83} through more recent works restricted to the low and intermediate mass stars of concern here \\citep{wei00,mar01,chi01,sch01,sch02,sie02}. In particular, \\citet{fuj90}, \\citet{hol90} and \\citet{fuj00} have shown that low and intermediate mass stars follow evolutionary trajectories significantly different from those followed by population I and population II stars, becoming carbon stars at a much earlier stage.  Since the early 1990's, the nature of the evolution of metal poor stars has been illuminated considerably by the observations. Studies of the abundance patterns found in stellar spectra have revealed that carbon stars are not rare among metal poor stars and that the relative frequency of such stars increases with decreasing metallicity, amounting to more than $\\sim 20 \\%$ for stars with $\\feoh \\lesssim -2.5$ \\citep{ros98}. It has been known for a long time that most population I carbon stars which are not evolved beyond the first red giant branch are in binaries with a white dwarf companion \\citep{mcc80};  this is also true of population II CH stars \\citep{mcc84}. Thus, the frequency of carbon stars among low metallicity stars may suggest a high rate of binary star formation in the gas clouds out of which such stars were formed. A current lack of evidence for radial velocity variations does not necessarily rule out the possibility that the observed star is now in a very wide binary or that the initial binary was disrupted at some point after the mass-transfer episode occurred. Since an initial binary must be wide enough to accommodate an AGB star which does not overflow its Roche lobe, and since, if anything, the binary becomes wider as the AGB star loses mass \\citep[see, e.g., the discussion in][]{ibe00}, it is probable that the binary separation does not decrease. Furthermore, given the fact that wide binaries are relatively easily disrupted in close stellar encounters \\citep{heg75} , it may happen that, in $10^{10}$ yr of traveling through the Galaxy, \\mmps\\ has lost a one-time companion. The radial velocity of \\mmps\\ has recently been determined to be $44.5 \\pm 0.5$ km s$^{-1}$, based on observations over a 373 day time span \\citep{bes04}. Although significant radial velocity variations have not been detected, observations over a much longer baseline would be well worth undertaking.  Extremely metal-poor stars show a large dispersion in the abundance ratios among metals. The fact that, for metallicities less than $\\feoh \\simeq -2.5$, this dispersion increases with decreasing metallicity may be interpreted to mean that these stars were born from interstellar gas which was polluted by a small number of supernovae, so that stellar abundances reflect intrinsic variations in the yields of those individual supernovae that contribute to the pollution of the gas.  The abundances of many neutron-rich elements vary considerably from star to star even among those stars which have the abundance ratio $[\\nucm{}{Ba}/ \\nucm{}{Eu}]$ very close to the solar r-process value \\citep{mcw98,aok02,hon04}. This fact has been used to argue that the Ba found in the spectra of stars of low metallicity has been created by an r-process mechanism in supernovae and that the production rate varies considerably from one supernova to another \\citep{mcw98}. However, there are also large variations from one star to another in the abundances of those neutron-rich elements which are unquestionably made by the s-process \\citep[see, e.g.,] [and references therein]{aok01,rya03}. It is known that thermally pulsing AGB stars are the most likely sites for the formation of heavy s-process elements. Since most extremely metal-poor stars are not AGB stars, it seems plausible to invoke the erstwhile presence of an initially more massive binary companion which has produced the s-process elements, dredged these elements to the surface, and then transferred enriched matter to the less evolved component, a scenario suggested by \\citet{fuj00} and explored by \\cite{gor01}, \\cite{aok01}, and \\citet{iwa04}. These authors point out that s-process nucleosynthesis in very metal-poor stars may differ significantly from that which is thought to operate in younger populations.  The purpose of the present paper is to review the possible scenarios for explaining the observed properties of \\mmps\\ and to investigate the possibility of discriminating among these scenarios on the basis of our current understanding of nucleosynthesis in extremely metal-poor stars --- an understanding which has been achieved by comparing results of theoretical evolutionary calculations with the observations. In the following, we define stars of metallicity less than $\\feoh \\lesssim -2.5$ as ``extremely metal poor'' or EMP stars and examine the implications of assuming that \\mmps\\ was born as a single star or as a low-mass component in a binary system. We use the results of new evolutionary models of Pop.~III stars of mass in the range $0.8 - 4 \\msun$ to elaborate the binary scenario more fully. Stellar model construction employs up-to-date input physics \\citep{sud03} which differs in significant ways from the input physics employed by \\citet{fuj90} and by \\citet{fuj00}, and the results, correspondingly, differ quantitatively in significant ways. In addition, we rely upon new results concerning neutron-capture nucleosynthesis in a helium-flash driven convection zone that is a consequence of the ingestion of protons into this convection zone (Aikawa, Nishimura, Suda, Fujimoto \\& Iben 2004, in preparation).  In discussing the origin of \\mmps, it is important to consider separately (a) the source of iron group elements, (b) the source of light elements such as the $\\alpha$-rich elements C and O and the secondary elements N and Na, and (c) the source of s-process elements. The two possibilities for explaining the metals are: (1) accretion of metal-rich gas by a first generation (Pop.~III) star; and (2) formation as a second generation star out of matter in which metals are present in consequence of mixing with the ejectum of a first generation supernova. In the binary scenario, the enhancements of CNO elements and of s-process elements are assumed to be due to accretion from an AGB companion which has experienced one or more dredge-up episodes that bring to the surface the result of nucleosynthesis and mixing in its interior.  There is not yet a consensus as to the conditions which are necessary for the formation of low metallicity, low mass stars. The lack of a sufficient abundance of metals to provide a straightforward and effective cooling mechanism for metallicities $\\feoh \\lesssim -4$ \\citep[] [see also Omukai 2000]{yos80} is the primary hurdle facing low mass star formation in the early Universe. In the absence of metals, the hydrogen molecule can produce cooling, but estimates of the minimum Jean's mass vary from $\\sim 1 \\msun$ \\citep{sab77} and $\\sim 0.1 \\msun$ \\citep{pal83} to $\\gtrsim 60 \\msun$ \\citep{yon72}. Several authors argue that low mass stars can form from a metal-free primordial gas cloud \\citep[e.g.,][see also Uehara et al.~1996]{ree76}, whereas others suggest that only massive stars can form from such a cloud \\citep[e.g.,][]{omu98,bro99}. To complicate the picture even further, two dimensional hydrodynamical simulations \\citep{nak01,nak02} suggest a bimodal initial mass function for Pop.~III stars. Such complexity may arise from the fact that, in primordial clouds of mass $\\sim 10^6 \\msun$ which first collapse \\citep[e.g.,][]{teg97}, the scarcity of free electrons \\citep[mole density $\\sim 10^{-4} $;][]{gal98} limits the formation of ${\\rm H}_2$ molecules which can act as cooling agents. On the other hand, if the cloud temperature is raised above $10^4$K and gas is re-ionized, the ionization fraction remains quite high during the subsequent radiative cooling phase and increases the abundances of ${\\rm H}_2$ and HD sufficiently to allow the birth of low mass stars even in the complete absence of metals (Shapiro \\& Kang 1987; see Uehara \\& Inutsuka 2000 for a recent computation including HD molecules). This may be the case for primordial clouds as massive as $ \\gtrsim 10^8 \\msun$ for which virial temperatures are larger than $10^4$ K, although the collapse is delayed somewhat compared with the case of less massive clouds \\citep[see e.g.,][]{nis99}. The same situation also prevails even for less massive primordial clouds when the interstellar gas is swept up and heated by the supernova shock produced by a first generation massive star \\citep{mac04}. It is also argued that low mass star formation is possible in clouds of the primordial composition if matter is irradiated by sufficiently strong FUV radiation from first generation massive stars \\citep{hai96b,omu03}. If the observed enrichment of CNO elements can be explained as acquired after birth, the very existence of \\mmps\\ indicates that some of these or other formation mechanisms of low mass stars have to work even in the gas clouds completely devoid of metals.  Observationally, there has been long, continuous interest in search for stars of lower metallicity and/or completely devoid of metals since the first discovery of stars of $1/10 \\sim 1/100$ solar metallicity in the 1950's \\citep{cha51}. At the beginning of the 1970's, a survey up to a limiting magnitude of $B \\simeq 11.5$ found no star with $\\feoh < -3$ \\citep{bon70}, and encouraged the idea that no star of metallicity smaller than this had been formed in our Universe. In the early 1980's, however, as by-products of the study of high-velocity stars and blue subdwarfs, the dwarf G64-12 with $\\feoh = -3.5$ \\citep{car81} and the giant CD-38$^{\\circ}$245 with $\\feoh = -4.5$ \\citep{bes84} were found, at magnitudes, respectively, of $B = 11.8$ and $B = 12.0$, both below the limiting magnitude of the survey by \\citet{bon70}. In the early 1990's, the HK survey \\citep{bee92}, with a limiting magnitude of $B = 15.5$, uncovered over 100 stars with $\\feoh < -3$, and, yet, it found no star with $\\feoh < -4.0$. As an aside, we note that the metallicity of CD-38$^{\\circ}$245 has been revised upwards to $\\feoh = - 3.92$ \\citep{rya96}. Some time ago, a high velocity carbon dwarf G77-61, at $B = 15.60$ ($V = 13.90$), was assigned a metallicity of $\\feoh = -5.6$ \\citep{gas88}; however, the low surface temperature of the star and the presence of crowded molecular lines makes the abundance analysis very difficult, and further work is required to confirm the early estimate of metallicity. Finally, the Hamburg/ESO survey, with a limiting magnitude of $B = 17.5$ \\citep{chr99}, discovered \\mmps\\ with metallicity $\\feoh=-5.3$ at magnitude $B = 15.89$. Along with the increase in the limiting magnitude, therefore, we have been able to detect the stars of smaller metallicity. Provided that the sample of stars of metallicity as small as, and still smaller than, that of \\mmps\\ can be considerably augmented, the discovery of \\mmps\\ may be the beginning of a new epoch in the study of the early history of the Universe using low mass stellar survivors as a tool.  The paper is organized as follows: In \\S 2, formation scenarios are discussed. In \\S 3, the characteristics of evolution and nucleosynthesis in low and intermediate mass metal-free and metal-poor stars are reviewed and compared with the results of recent observations of extremely metal-poor stars. In \\S 4, the origin of \\mmps\\ is discussed and the binary scenario is elaborated. Conclusions and further discussion are provided in \\S 5. ",
        "conclusions": "During the past decade, there has been considerable observational and theoretical progress in understanding the properties of extremely metal-poor (EMP) stars in the Galactic halo. Thanks to the HK survey \\citep{bee92}, the number of known EMP stars ($\\feoh \\lesssim -2.5$) has reached more than 100, and the spectroscopic characteristics of these stars have been revealed through detailed studies using large telescopes. By analyzing existing theoretical and observational evidence and by utilizing the results of new computations, we have focused in this paper on the peculiarities that distinguish EMP stars from Pop.~I and II stars and have presented a general theoretical framework for understanding the evolution of EMP stars and for understanding the nucleosynthesis which has taken place in their interiors. Our framework relies heavily on a binary scenario in which the primary has produced, and transferred to the secondary, many of the isotopes which are present at the surface of an observed EMP star.  A distinct feature that characterizes the evolution of EMP stars of low and intermediate mass is the helium-flash driven deep mixing (He-FDDM) phenomenon, which is triggered when the outer edge of a convective zone driven by a first helium-shell flash (which occurs near the beginning of the AGB phase in intermediate mass models and at the tip of the RGB in low mass models) extends into hydrogen-rich material. The first discussions of the He-FDDM mechanism \\citep{fuj90,hol90,fuj95,fuj00} drew attention to the surface enhancement of carbon and nitrogen made possible by this mechanism. In the present paper, we have investigated another consequence of the He-FDDM event, namely, the nucleosynthesis of s-process elements which occurs when neutrons are released by the reaction $\\nucm{13}{C} (\\alpha, n) \\nucm{16}{O}$ in the helium- and carbon-rich convective zone. We demonstrate that products of this nucleosynthesis can lead to surface enrichments of O, Ne, Na, and Mg and, if iron-group seed nuclei are present in the convective zone, to surface enrichments of heavy s-process elements. In addition, stars which evolve to the TPAGB phase after having experienced a He-FDDM event can also develop a large $\\nucm{12}{C}/\\nucm{13}{C}$ ratio at the surface in consequence of third dredge-up events.  Our binary scenario gives a reasonable account of the observed properties of EMP stars, such as the very high frequency of carbon-rich stars and s-process nucleosynthesis products which differ from those made by stars of younger populations. The scenario enables us in principle to identify modifications in surface abundances which, after birth in an unpolluted primordial cloud, first generation EMP stars may have experienced due to accretion of primordial matter polluted by the ejecta of first generation supernovae; this ability is essential if we wish to use low mass EMP star survivors as tools to probe the early Universe.  We have constructed a specific binary scenario to account for the observed abundance characteristics of \\mmps. The initial binary system has the properties: separation $\\sim 18$ AU, orbital period $\\sim 45$ years, a secondary of mass $\\sim 0.8 \\msun$, and a primary of mass in the range $1.2 \\lesssim M / \\msun \\lesssim 3$. The primary follows an evolutionary path of the Case II$^\\prime$ variety in the classification scheme defined by \\citet{fuj00}. In consequence of experiencing a He-FDDM episode, the primary develops surface enhancements of C and N and, in consequence of experiencing third dredge-up events during the subsequent TPAGB phase, develops a large overabundance of C relative to N as well as enhancements of O, Ne, and Na, formed during the He-FDDM and subsequent thermal pulses. After ejecting its hydrogen-rich envelope in a superwind, the primary evolves into a white dwarf. The secondary accretes heavy-element enriched matter from the wind emitted by the primary and, when it evolves into a red giant, it mixes this enriched matter into a deep convective envelope, establishing surface abundance peculiarities similar to those observed for \\mmps.  Because the iron abundance is so small ($\\feoh=-5.3$), the presence of iron-group elements in \\mmps\\ can be attributed to either (1) accretion after birth of gas in a parent primordial cloud which has been polluted with material ejected by one or more first generation supernovae or (2) birth out of already polluted matter in the parent cloud.  We have shown that further light on the source of iron-group elements can be shed by comparing an observed distribution of heavy s-process elements with theoretical expectations. An abundance ratio $[{\\rm Pb}/ {\\rm Fe}] \\simeq 1 \\sim 2$, coupled with a large ${\\rm Pb} / {\\rm Ba}$ ratio would be evidence that \\mmps\\ is a second generation star, or a Pop.~III star in a binary with a primary of initial mass in the range $2.5 \\lesssim M / \\msun \\lesssim 3.0$. A lack of heavy s-process element enrichment would indicate that \\mmps\\ is really a Pop.~III star in a binary with a primary component of initial mass in the range $1.2 \\msun \\lesssim M \\lesssim 2.5 \\msun$. Unfortunately, current observations provide only upper limits on the abundances of light and main-line s-process elements, so no conclusion can as yet be drawn. A ratio $[{\\rm Pb} / {\\rm Fe}] \\simeq 1 \\sim 2$ translates into a ratio $[{\\rm Pb} / {\\rm H}] \\simeq -4 \\sim -3$ for \\mmps. This is difficult to detect with present facilities, and we will probably have to wait for future observations with a next generation large telescope and the elaboration of model atmosphere including, for example, 3D hydrodynamic simulations to establish definitively whether or not HE0107-5240 is a Pop.~III star.  With regard to the current binary status of \\mmps, variations in the radial velocity of the size predicted by our scenario cannot be excluded by extant spectroscopic observations which, to date, cover only 52 days \\citep{chr04} and 373 days \\citep{bes04}. Because of wind mass loss, the initial binary system may have been considerably widened. Adopting Jean's theorem which predicts the constancy of the product, semi-major axis times the total mass, we would expect the current binary to have the characteristics $a \\simeq 34$ AU and $P_{\\rm orb} \\simeq 150$ years, giving an orbital velocity of $\\sim 7 \\hbox{ km s}^{-1}$. Confirmation of such a small velocity demands long term observations at high dispersion.  As far as alternative single star interpretations of \\mmps\\ are concerned, any viable scenario must begin with formation out of gas with a singularly unusual abundance distribution (including the currently observed carbon and oxygen enhancements) of the sort not expected by mixing products of normal supernovae with primordial matter. As an example, a scenario proposed by \\citet{ume03} supposes that the gas out of which \\mmps\\ was formed acquired a carbon abundance of $[{\\rm C}/{\\rm H}] = -1.3$ after the mixing of an unusually small amount of primordial matter with a supernova ejectum in which the mass of carbon is $\\simeq 0.2 \\msun$. Given that a typical type II supernova ejects a mass of the order of $0.1\\sim 1 \\msun$ in the form of Fe, a most peculiar process of star formation is required to account at the same time for a ratio of $\\feoh \\simeq -3$ characteristic of EMP stars. It is to be noted that the mass of carbon in the supernova ejectum cannot be much larger than the mass of iron in the ejectum since, in order to obtain the observed ${\\rm C} /{\\rm O}$ ratio, the carbon can only have come from the shell of partial helium burning existing above the region where iron-group elements are formed. In addition, since, in contrast with nitrogen, which can be produced in the interior and brought to the surface during the first dredge-up phase on the RGB, sodium cannot be formed during the evolution of a low mass star, and a non-canonical mechanism for sodium enrichment has to be invoked. Given this hurdle for the (any) single star scenario, it seems reasonable to accept the observed Na enhancement as evidence that some of the matter in the convective envelope of \\mmps\\ was formed in the interior of a primary companion during the AGB phase.  Although there is a controversy as to whether or not a single star of mass as small as that of \\mmps\\ can be formed in a primordial cloud, a condensation as massive as the $2 \\sim 3.8 \\msun$ predicted by our binary scenario is not excluded by current theory \\citep{nak01}, which suggests a bi-modal star-formation mass function for first generation stars. The initial condensation could be formed as a first generation object in the first collapsed, primordial cloud of total (dark and baryonic) mass $\\sim 10^6 \\msun$. Another possible site is a primordial cloud of mass $\\gtrsim 10^8 \\msun$, for which virial temperatures are higher than $10^4$ K. In such a cloud, \\mmps\\ could have been born as a first generation star, but the formation epoch would be delayed in comparison with the formation epoch of first generation stars in a lower mass collapsed cloud.  If \\mmps\\ is a second generation star, it also must have been formed in a cloud of total mass $ \\gtrsim 10^8 \\msun$. In a primordial cloud of total mass $ \\sim 10^6 \\msun$ and baryonic mass $\\sim 10^5 \\msun$, contamination of primordial gas with the $0.1 \\sim 1 \\msun$ of iron ejected by a typical supernova would produce an iron abundance $\\feoh = -3 \\sim -2$. This metallicity is appropriate for most EMP stars known to date, but it is hundreds of times larger than the metallicity of HE0107-5240. In order to achieve a metallicity appropriate for \\mmps, the contaminating supernova must eject an abnormally small mass of iron. But, the supernova explosion would be so weak that the remnant would dissolve into the interstellar gas instead of compressing it to collapse conditions, and star formation would not be triggered \\citep{mac04}.  As described in the introduction, the history of the search for extremely metal-poor stars has taught us that, as the limiting magnitude of a survey is increased, stars of lower metallicity are detected, suggesting that, for EMP stars, there may be a correlation between typical metallicity and apparent luminosity, and hence, a relationship between typical metallicity and spatial distribution. In the current framework of structure formation, galaxies have been formed by the merging of lower-mass building blocks. If the parent cloud of \\mmps\\ is of mass $ \\gtrsim 10^8 \\msun$, it differs from parent clouds of mass $\\simeq 10^6 \\msun$ out of which many second generation stars with $\\feoh \\simeq -4 \\sim -2.5$ may have been formed. Accordingly, in addition to commenting on initial abundances in primordial matter and characteristics of nucleosynthesis in supernova explosions in the early Universe, EMP stars may also reveal the masses and locations in our Galaxy of their parent clouds.  The discovery of \\mmps\\ has also demonstrated the importance of probing for Pop.~III stars by sorting according to carbon-star characteristics rather than according to the weakness and/or absence of the Ca II K line. As pointed out by FII00, it makes sense to search for carbon stars since low-mass stars of $\\feoh \\lesssim -4.5$ spend their final nuclear burning lives as luminous carbon stars on the horizontal and asymptotic giant branches. Assuming that it has the luminosity of a typical red giant, \\mmps\\ is at a distance of $\\sim 10$ kpc. Because the luminosity of a typical intermediate mass TPAGB star is larger by about a factor of ten than that of a typical low mass RGB star, TPAGB stars are detectable at much larger limiting magnitudes than are RGB stars. The down side is that the lifetime of a typical TPAGB star is over ten times smaller than that of a typical RGB star. Nevertheless, with CCD cameras, larger telescopes, and patience, we may be able to find Pop.~III carbon stars further out in the Galactic halo and perhaps even in intergalactic space.  Recently, \\citet{mar02} have reported a number of faint high latitude carbon stars from the data obtained by the Sloan Digital Sky Survey (SDSS). Among these stars there may be EMP carbon stars of the sort that we suggest looking for, even though \\citet{mar02} argue that nearby dwarf carbon stars outnumber giants in their sample. Follow-up spectroscopic observations may decide the issue. Since their selection is based on comparisons with known types of carbon stars in the five color system of SDSS \\citep{kri98}, there is the possibility that Pop.~III carbon stars may be missed. Because of different atmospheric properties related to higher surface temperatures and different chemistry, the color properties of EMP carbon stars may differ significantly from those of more familiar carbon stars of Pop.~I, or even from those of EMP carbon stars with $\\feoh \\simeq -3$. In any case, it is important to continue the search for carbon stars of low metallicities using every stratagem that can be devised, including methods that make use of spectral lines for which carbon molecules are responsible. From extremely metal-poor carbon stars found at extreme distances, or even from the absence of such stars, we may learn important information about the early Universe."
    },
    "0402/astro-ph0402666_arXiv.txt": {
        "abstract": "A new method is presented for modeling the transformation between two polarimetric pulse profiles in the Fourier domain.  In practice, one is a well-determined standard with high signal-to-noise ratio and the other is an observation that is to be fitted to the standard.  From this fit, both the longitudinal shift and the polarimetric transformation between the two profiles are determined.  Arrival time estimates derived from the best-fit longitudinal shift are shown to exhibit greater precision than those derived from the total intensity profile alone.  In addition, the polarimetric transformation obtained through this method may be used to completely calibrate the instrumental response in observations of other sources. ",
        "introduction": "\\label{sec:modeling}  In radio astronomy, a pulsar's mean polarimetric pulse profile is measured by averaging the observed Stokes parameters as a function of pulse longitude.  By integrating many well-calibrated pulse profiles, a standard profile with high signal-to-noise ratio (SNR) may be formed and used as a template to which individual observations are fit.  For example, in Appendix A of Taylor (1992), a method is presented for modeling the relationship between standard and observed total intensity profiles in the Fourier domain.  In the current treatment, the scalar equation that relates two total intensity profiles is replaced by an analogous matrix equation, which is expressed using the Jones calculus.  The polarization of the electromagnetic field is described by the coherency matrix, ${\\mbf\\rho}=(S_0\\,\\pauli{0}+\\mbf{S\\cdot\\sigma})/2$, where $S_0$ is the total intensity, $\\mbf{S} = (S_1,S_2,S_3)$ is the Stokes polarization vector, $\\pauli{0}$ is the $2\\times2$ identity matrix, and $\\mbf{\\sigma} = (\\pauli{1},\\pauli{2},\\pauli{3})$ are the Pauli spin matrices (Britton 2000).  Under a linear transformation of the electric field vector as represented by the Jones matrix, ${\\bf J}$, the coherency matrix is subjected to the congruence transformation, ${\\mbf{\\rho}^\\prime}={\\bf{J}}\\mbf{\\rho}{\\bf{J}}^\\dagger$ (Hamaker 2000).  Let the coherency matrices, $\\mbf{\\rho}^\\prime (\\phi_n)$, represent the observed polarization as a function of discrete pulse longitude, $\\phi_n$, where $0\\le n< N$ and $N$ is the number of pulse longitude intervals.  Each observed polarimetric profile is related to the standard, $\\mbf{\\rho}_0(\\phi_n)$, by the matrix expression, \\begin{equation} \\label{eqn:model} \\mbf{\\rho}^\\prime (\\phi_n) = \\mbf{\\rho}_{\\mathrm DC} + \\mbf{\\rho}_{\\mathrm N}(\\phi_n) + {\\bf J} \\mbf{\\rho}_0 (\\phi_n - \\varphi) {\\bf J}^\\dagger, \\end{equation} where $\\mbf{\\rho}_{\\mathrm DC}$ is the DC offset between the two profiles, $\\mbf{\\rho}_{\\mathrm N}$ represents the system noise, ${\\bf J}$ is the polarimetric transformation and $\\varphi$ is the longitudinal shift between the two profiles.  The Jones matrix, ${\\bf J}$, is analogous to the gain factor, $b$, in equation (1) of Taylor (1992). However, as ${\\bf J}$ has seven non-degenerate degrees of freedom, the matrix formulation introduces six additional free parameters.  The discrete Fourier transform (DFT) of \\eqn{model} yields \\begin{equation} \\mbf{\\rho}^\\prime (\\nu_m) = \\mbf{\\rho}_{\\mathrm N}(\\nu_m) + {\\bf J} \\mbf{\\rho}_0 (\\nu_m) {\\bf J}^\\dagger \\exp(-i2\\pi m \\varphi), \\label{eqn:fourier_rho} \\end{equation} where $\\nu_m$ is the discrete pulse frequency.  Given the measured Stokes parameters, $S_k^\\prime(\\phi_n)$, and their DFTs, $S_k^\\prime(\\nu_m)$, the best-fit model parameters will minimize the objective merit function, \\begin{equation} \\chi^2 = \\sum_{m=1}^{N/2} \\sum_{k=0}^3 { |S_k^\\prime(\\nu_m) - \\trace[\\pauli{k}\\;\\mbf{\\rho}^\\prime(\\nu_m)]|^2 \\over \\varsigma_k^2 }, \\label{eqn:merit} \\end{equation} where $\\varsigma_k$ is calculated from the noise power and $\\trace$ is the matrix trace.  As in van Straten (2004), the partial derivatives of \\eqn{merit} are computed with respect to both $\\varphi$ and the seven parameters that determine ${\\bf J}$.  The Levenberg-Marquardt method is then applied to find the parameters that minimize $\\chi^2$. ",
        "conclusions": "When compared with the scalar equation used to model the relationship between total intensity profiles, the matrix equation presented in Section~1 quadruples the number of observational constraints while introducing only six additional free parameters.  By completely utilizing all of the information available in mean polarimetric pulse profiles, arrival time estimates may be obtained with greater precision than those derived from the total intensity profile alone. In addition, the modeling method may be used to uniquely determine the polarimetric response of the observatory instrumentation using only a short observation of a well-known source."
    },
    "0402/astro-ph0402385_arXiv.txt": {
        "abstract": "The apparently repeating microlensing event OGLE-2003-BLG-095 is analyzed. Data were obtained from the OGLE Internet archive and exist in the public domain. The source is relatively bright, with an unmagnified (but possibly blended) $I$-band magnitude of 15.58, and the signal-to-noise ratio of the data is excellent. The light curve shows two distinct, smooth peaks characteristic of a double microlensing event. It can be modeled as either (1) microlensing by a binary lens or (2) microlensing of a binary source, with the latter model providing a statistically superior fit. However due to apparent low-amplitude variability of the source, the interpretation is somewhat ambiguous. OGLE-2003-BLG-095 is only the second possible case in the literature for microlensing of a well-resolved binary source. ",
        "introduction": "Gravitational microlensing surveys were originally suggested by \\citet{pacz86} as a means of detecting dark matter in the Galaxy in the form of massive compact objects, commonly abbreviated as MACHOs following \\citet{grie91}. Various efforts were undertaken to search for such objects, including the Optical Gravitational Lensing Experiment (OGLE; \\citealt{udal92}), the MACHO project (e.g., \\citealt{alco93}), and several others. Currently ongoing projects, most notably OGLE and MOA \\citep{bond01}, are detecting hundreds and dozens of new microlensing events each year, respectively. The new events found by these groups are made public in real time to maximize scientific gain by enabling follow-up of interesting objects by the astronomical community.  In addition to the possibility of detecting dark matter MACHOs, microlensing surveys also offer the interesting opportunity to study the populations of stellar lenses (which constitute at least a significant fraction of all events) and sources. Since most stars are located in multiple systems, many microlensing events show deviations from the standard Paczy{\\'n}ski (point lens, point source) light curve. For example, some events show the effects of caustic crossings \\citep{mao91,goul92}; this provides clear evidence of lens binarity and can be used to constrain the combination of the distances to the source and lens, the source-lens relative motion, and the physical extent of the source, as well as revealing the mass ratio of the lens system. A small additional fraction of microlensing events should show clear signatures of source binarity in multiple peaks and color shifts; this fraction was estimated to be 2--5\\% by \\citet{grie92}. To date no convincing example of such an event has been reported in the literature; the best case is MACHO-96-BLG-4, but the light curve of this event can be equally well fit by a binary source or binary lens model, with the binary lens providing the most natural interpretation \\citep{alco00}. Several authors \\citep{domi98,han98,dist00} have proposed explanations for the lack of such clear binary source microlensing events. Essentially, a variety of different effects including blending, unequal luminosity of binary components, and simple coincidence conspire to make most binary source microlensing light curves resemble single source, single lens events. A handful of other events, such as MACHO-96-LMC-2 \\cite{alco01}, show signatures of binary orbital motion of the source; in this particular case, it is not clear whether both components of the binary contribute significantly to the total source flux. In any case the light curves of such events do not match the expectation of having multiple distinct, well-resolved peaks as described above.  This work analyzes the light curve of \\evt, which shows two distinct microlensing peaks. The light curve can be fit by a binary lens or a binary source model; however, the binary source model is statistically preferred. The light curve is presented in \\S\\ref{sec:data}, \\S\\ref{sec:model} contains a description of the models and fitting procedure, and the results are analyzed in \\S\\ref{sec:disc}. ",
        "conclusions": "\\label{sec:conc}  The light curve of \\evt\\ shows two well-separated smooth peaks that can be modeled as either microlensing by a binary lens or microlensing of a binary source. Both models are physically plausible and can possibly be tested by future observations of the source and/or lens. The binary lens model provides a more familiar explanation of the event since such phenomena have been observed many times in the past, while the binary source model is preferred on statistical grounds. Given the apparent low-amplitude variability of the source (or at least one component of it) however, statistical distinctions of the magnitude that separate the two models are somewhat suspect, especially considering the improbably large microlens parallaxes indicated by the best fit models. Despite these complications, \\evt\\ remains a plausible candidate for microlensing of a well-resolved binary source -- one of only two such observed to date."
    },
    "0402/astro-ph0402450_arXiv.txt": {
        "abstract": "\\shell is a well-defined neutral hydrogen shell discovered in the VLA Galactic Plane Survey (VGPS). Only the blueshifted side of the shell was detected. The expansion velocity and systemic velocity were determined through the systematic behavior of the \\HI emission with velocity.  The center of the shell is at ($l$,$b$,$v$)=($23\\fdg05$,$-0\\fdg77$,$+117\\ \\kms$). The angular radius of the shell is $6\\farcm8$, or 15 pc at a distance of 7.8 kpc. The \\HI mass divided by the volume of the half-shell implies an average density $n_H = 11\\ \\pm 4\\ \\rm cm^{-3}$ for the medium in which the shell expanded.  The estimated age of \\shell is 1 Myr, with an upper limit of 2 Myr. The modest expansion energy of $2 \\times 10^{48}\\ \\rm erg$ can be provided by the stellar wind of a single O4 to O8 star over the age of the shell. The $3\\ \\sigma$ upper limit to the 1.4 GHz continuum flux density ($S_{1.4} < 248\\ \\rm mJy$) is used to derive an upper limit to the Lyman continuum luminosity generated inside the shell. This upper limit implies a maximum of one O9 star (O8 to O9.5 taking into account the error in the distance) inside the \\HI shell, unless most of the incident ionizing flux leaks through the \\HI shell. To allow this, the shell should be fragmented on scales smaller than the beam (2.3 pc). If the stellar wind bubble is not adiabatic, or the bubble has burst (as suggested by the \\HI channel maps), agreement between the energy and ionization requirements is even less likely.  The limit set by the non-detection in the continuum provides a significant challenge for the interpretation of \\shell as a stellar wind bubble.  A similar analysis may be applicable to other Galactic \\HI shells that have not been detected in the continuum. ",
        "introduction": "The VLA Galactic Plane Survey (VGPS) is part of an international effort to map atomic hydrogen and other tracers of the Galactic interstellar medium with a resolution of $1'$. Previously, large parts of the Galactic plane in the northern sky were covered by the Canadian Galactic Plane Survey (CGPS) \\citep{taylor2003}, and in the southern sky by the Southern Galactic Plane Survey (SGPS) \\citep{mcclure2001}. The VGPS covers the first Galactic quadrant in the vicinity of the celestial equator, for which the Very Large Array is the most suitable instrument. The VGPS survey area extends from Galactic longitude $18\\arcdeg$ to $67\\arcdeg$. The latitude coverage varies from $|b| < 1\\arcdeg$ at the low longitudes to $|b| < 2\\arcdeg$ at the high longitudes.  An outline of the survey area was shown by \\citet{taylor2002}.  One important objective of these high resolution \\HI surveys is to study the effect of stellar wind and supernova explosions on the interstellar medium. There is a rich literature on this subject, and we limit the discussion to some examples that relate to the subject of this paper.  The effects of stellar wind and supernovae may be manifested on scales of hundreds of parsecs for super bubbles, chimneys and worms, e.g.  \\citet{heiles1979}, \\citet{heiles1984}, \\citet{normandeau1996}, \\citet{mcclure2000}, \\citet{english2000}, \\citet{stil2001}, \\citet{uyaniker2002}, \\citet{mcclure2002}, to $\\sim 10\\ \\rm pc$ for winds of single stars. Smaller bubbles originating from a single star may be found around Wolf-Rayet stars \\citep[for a discussion of radio observations]{cappa2002} and some other early type stars, e.g. \\citet{higgs1994}, \\citet{normandeau2000}, \\citet{carral2002}.  An interesting question in this respect is what fraction of the stellar wind and supernova ejecta produced in the disk breaks out of the Galactic disk and flows into the Galactic halo. Whether or not a breakout occurs depends on the scale of the bubble and the scale height of \\HI in the disk. In this context, small bubbles represent events in which matter and energy ejected by massive stars are retained in the disk. As such, smaller bubbles provide a different perspective on the Galactic energy budget, as well as a probe of conditions that relate to the release of enriched matter and energy into the disk and the halo. Parsec scale \\HI bubbles have become accessible for systematic study through the recent high-resolution \\HI surveys.  The stellar winds of OB stars are driven by the ultraviolet continuum. Therefore, a strong stellar wind and a high ionizing flux are correlated.  A stellar wind bubble can be completely or partly ionized due to the Lyman continuum flux of the central star. It is not clear a priori whether the ionization of a shell can be detected in the radio continuum images of the VGPS. Confusion with unrelated emission may inhibit detection for larger shells. Well-documented examples of smaller shells that would be detectable in a survey such as the VGPS exist in the literature, e.g.  \\citet{higgs1994}, \\citet{cappa1999}. In this paper we present the small \\HI shell \\shell discovered in the VGPS, and the implications of its non-detection in the continuum. ",
        "conclusions": "We report the discovery of the small \\HI shell \\shell in the VLA Galactic Plane Survey (VGPS). The physical parameters of \\shell are well constrained because its velocity places it close to the tangent point.  At $(l,b,v) =$ ($23\\fdg05$,$-0\\fdg77$,$+117\\ \\rm \\kms$), the distance of the shell is $7.8\\ \\pm\\ 2\\ \\rm kpc$. The expansion kinetic energy is found to be $2 \\times 10^{48}\\ \\rm erg$ for a shell mass of $2.5\\times 10^3\\ \\rm M_{\\odot}$ and expansion velocity $9\\ \\pm\\ 1\\ \\kms$. The age of the shell is $1\\ \\rm Myr$ with a strong upper limit of $2\\ \\rm Myr$. The average density of the medium in which the shell expanded derived from the mass and the volume of the shell is $n_H = 11\\ \\pm 4\\ \\rm cm^{-3}$. This relatively high density and the one-sided morphology of the \\HI shell are suggestive of a density gradient or a cloud in the vicinity of the shell.  The \\HI shell has no counterpart in the VGPS 1.4 GHz continuum image, with an upper limit to the 1.4 GHz flux density $S_{1.4} < 248\\ \\rm mJy$.  The interpretation of \\shell as a stellar wind bubble is tested by combining the energy requirements of the shell with a limit to the number of OB stars inside the shell derived from the upper limit to the 1.4 GHz continuum emission.  If \\shell is an adiabatic stellar wind bubble \\citep{weaver1977}, the expansion energy requires the equivalent of the stellar wind of a single O4 to O8 star.  However, the $3\\sigma$ upper limit to the 1.4 GHz continuum flux density excludes more than one O9 star (O8 to O9.5 given the uncertainty in the distance), unless most of the incident ionizing flux leaks through the \\HI shell. For this to be the case, the \\HI shell should be highly fragmented on scales smaller than the beam (2.3 pc). The energy requirements for the shell are significantly larger if the bubble is not adiabatic, or if the bubble has burst, as suggested by the morphology in the \\HI channel maps. In this case, the discrepancy between the energy requirements of the shell and the maximum number of OB stars allowed inside the shell is even larger.  We conclude that the interpretation of the low-energy \\HI shell \\shell as a stellar wind bubble is questionable in view of the absence of continuum emission associated with the shell. A similar argument may be applicable to other Galactic \\HI shells that have not been detected in the continuum."
    },
    "0402/astro-ph0402499_arXiv.txt": {
        "abstract": "We have used the Wide Field and Planetary Camera~2 on board the Hubble Space Telescope to obtain $V$ and $I$ images of seven nearby galaxies. For each, we have measured a distance using the tip of the red giant branch (TRGB) method. By comparing the TRGB distances to published Cepheid distances, we investigate the metallicity dependence of the Cepheid period-luminosity relation. Our sample is supplemented by 10 additional galaxies for which both TRGB and Cepheid distances are available in the literature, thus providing a uniform coverage in Cepheid abundances between 1/20 and  2 (O/H)$_\\odot$. We find that the difference between Cepheid and TRGB distances decreases monotonically with increasing Cepheid abundance, consistent with a mean metallicity dependence of the Cepheid distance moduli of ${{\\delta{(m - M)}}/{\\delta[O/H]}} =  {-0.24 \\pm 0.05}$ mag~dex$^{-1}$. ",
        "introduction": "In the past decade, the uncertainty in the value of the Hubble constant based on the local distance scale ladder has decreased from roughly a factor of two to $\\pm$10--15\\% (e.g., Mould, Kennicutt,\\& Freedman 2000).  This breakthrough was made possible by the determination of HST-based Cepheid distances to 25 nearby galaxies, carefully selected to provide an accurate calibration for a variety of secondary distance indicators (Kennicutt et al.\\ 1995, Saha 1997).  The dominant source of systematic errors in the distance scale as a whole and in H$_0$ in particular resides in the calibration of the Cepheid period luminosity relation, most notably its zero point (which is tied to the distance to the Large Magellanic Cloud), and possible dependence (both in zero point and slope) on the metallicity of the variable stars. Reducing these extant  errors is imperative. For instance, the recent WMAP analysis of fluctuations in the cosmic  microwave background has produced a value of the Hubble constant $H_0 = 71 \\pm 4$ km~s$^{-1}$~Mpc$^{-1}$ (Bennett et al.\\ 2003). While in perfect agreement with the local value derived by the HST Key Project on the Extragalactic Distance Scale (72 $\\pm$ 8 km~s$^{-1}$~Mpc$^{-1}$; Freedman et al.\\ 2001, hereafter F01), tighter constraint on the latter would allow a more meaningful comparison of these two estimates.  The uncertainty in the metallicity dependence of the Cepheid period-luminosity (PL) relation is particularly troublesome. The galaxies used by the Key Project span a range of gas-phase metal abundances of roughly a factor of 50 ($-1.5 \\le [O/H] \\le 0.3$; Ferrarese et al.\\ 2000a), wide enough that a systematic change of 0.5 mag in the Cepheid distance moduli per factor 10 increase in abundance would by itself introduce a systematic error of approximately 10\\% in the Key Project value of H$_0$ (Kennicutt et al.\\ 1998, hereafter K98; Mould et al.\\ 2000; F01). Not only there are significant metallicity offsets between Cepheids in the Key Project  galaxies and those in the Large Magellanic Cloud (LMC), on which the PL calibration itself rests,  such offsets are different for  the mean samples used to calibrate secondary distance indicators (e.g., SNe~Ia, fundamental plane). This can potentially lead to systematic offsets in the values of $H_0$ derived from individual calibrators. A lack of constraints on the metallicity dependence of the Cepheid PL relation also hampers the interpretation of fully external tests of the zero point of the Cepheid distance scale (e.g., Hernstein et al.\\ 1999).   The magnitude of the metallicity dependence of the Cepheid PL relation is, unfortunately, poorly constrained, both theoretically or observationally.  Following K98, we describe such dependece in terms of the parameter $\\gamma$:  \\begin{equation} \\gamma = {\\delta {{(m - M)}_0}} / {\\delta{\\log Z}}, \\end{equation}  where $\\delta{{(m-M)}_0} = (m-M)_{\\mbox{0,Z}} - (m-M){\\mbox{0,LMC}}$, the difference between the distance modulus obtained with and without the metallicity correction, and $\\delta{\\log Z} = (\\log Z)_{\\mbox{LMC}} - (\\log Z)_{\\mbox{galaxy}}$. Note that $\\gamma$ reflects the net effect on {\\it distance determination}. Metallicity can affect both the luminosity of a Cepheid, as well as the color boundaries of the instability strip. Since a mean period-color relation is used to deduce reddening and extinction, the second of these effects can be the dominating influence. Thus $\\gamma$ really depends on the specific passbands chosen. In this paper we are mainly concerned with Cepheid distances based on $V$ and $I$ observations ($\\gamma$) which covers essentially all of the HST measurements.  Theoretical models of Cepheids predict metallicity dependences ranging from near zero (Saio \\& Gautschy 1998; Alibert et al.\\ 1999) to significant dependences (in either direction!) of up to $\\pm$0.3 mag~dex$^{-1}$ (Chiosi, Wood, \\& Capitanio 1993; Bono et al.\\ 1999; Sandage, Bell, \\& Tripicco 1999; Caputo et al.\\ 2000; Fiorentino et al.\\ 2002). Recent empirical determinations of $\\gamma$ have yielded an even larger range of values. The HST Key Project attempted to constrain $\\gamma$ in two ways, by comparing measured PL relations for two Cepheid fields in M101 differing in [O/H] by 0.7 dex, and by investigating a possible systematic difference between Cepheid and tip of the red giant branch (TRGB) distances for a sample of 10 galaxies (K98).  These two tests yielded marginal (1.5 $\\sigma$) detections of a metallicity dependence, with $\\gamma = -0.24 \\pm 0.16$ and $-0.12 \\pm 0.08$ mag~dex$^{-1}$ respectively. This led to a provisional correction of $-0.20$ mag~dex$^{-1}$ to the final Cepheid distances published by the Key Project team (F01).  However, value of $\\gamma$ between 0 and  $-$0.9 mag dex$^{-1}$ are supported by independent studies. By comparing Cepheid, TRGB, and RR Lyrae distances to the Magellanic Clouds and IC~1613, Udalski et al.\\ (2001) detected no significant metallicity dependence.  A null result was also derived by Ciardullo et al.\\ (2002) from a comparison of Cepheid and planetary nebula luminosity function (PNLF) distances to nearby galaxies. On the other hand, a comparison of Cepheid PL relations in the LMC and SMC by Sasselov et al.\\ (1997) produced $\\gamma = -0.4{^{+0.1}_{-0.2}}$.  A similar analysis, but applied to the Key Project galaxies, yielded $\\gamma = -0.4 \\pm 0.2$  (Kochanek 1997).  An even larger dependence was reported by Gould (1994), based on a re-analysis of M31 Cepheid observations of Freedman \\& Madore (1990).  The reason for the discrepancy between the various studies is the small number of galaxies used, the limited range of metal abundances spanned, and/or lack of quantifiable systematic errors.  Most recently, Tammann, Sandage \\& Reindl (2003) examined 321 Cepheid variables in the Galaxy with good $B$, $V$, and $I$ photometry by Berdnikov et al. (2000), and compared them to more than 1000 Cepheids in the LMC and SMC (Udalski et al. 1999b,c). They found that the Cepheid variables followed different period-color relations; LMC Cepheids were bluer than the Galactic ones, for example. They suggested that the observed differences in three galaxies for Cepheids with $\\log P > 1.0$ were due to metallicity differences. Kanbur et al. (2003) then studied Cepheids from the HST Key Project, and also from the Sandage-Tammann-Saha sample. They measured the distances to these galaxies using several different PL relations calibrated using Galactic and LMC Cepheids, including the new relation by Tammann et al. (2003). Kanbur et al. (2003) found that the Tammann et al. Galactic calibration yielded the same distances as the Udalski et al. (1999) LMC calibration, if the latter were corrected for a metallicity effect of $\\gamma = -0.2$ from Freedman et al. (2001). This suggested that the Galactic and LMC Cepheids did indeed follow different PL relations, and constrained the metallicity dependence to be $\\gamma \\sim -0.2$ mag dex$^{-1}$.  The goal of this paper is to correct all of these shortcomings and perform a more robust test of the metallicity dependence of Cepheid PL relation. We will follow the same technique used by K98, which was based on a comparison of Cepheid and TRGB distances for galaxies spanning a wide range in Cepheid metallicity. Compared to K98, our study benefits  from an increased sample size and a wider and more uniform range in Cepheid  metallicity. To  the 10 galaxies analyzed by K98, we add seven new measurements, covering a range in Cepheid abundances of 0.05 $-$ 2 in $Z/Z_\\odot$. A comparison of  TRGB and Cepheid distances provides an especially powerful test for a metallicity dependence of the Cepheid PL relation.The method is transparent and robust: over the metallicity range spanned by our galaxies, the TRGB magnitude in the $I$-band is insensitive to the metal abundance and age of the stellar population (Da Costa \\& Armandroff 1990; Lee, Freedman \\& Madore 1993; Salaris \\& Cassisi 1997; Sakai 1999). Furthermore, the metallicities of the halo fields targeted by TRGB observations do not correlate with those of the disk Cepheids, therefore even a small metallicity dependence of the TRGB would not introduce any systematic biases in our results. Finally, we also make the implicit assumption that the Oxygen fraction with respect to the total metal content is constant for all galaxies used in the application presented in this paper.   The paper is organized as follows: in \\S2, we discuss the observations and reduction of the HST/WFPC2 TRGB data obtained as part of this program for six galaxies (plus one downloaded from the public HST archive). \\S3 deals with the TRGB distances, including those that had been published prior to this paper. Cepheid distances, all of which have been previously published, are discussed in \\S4. The Cepheid abundances, which are derived from those of nearby \\hii\\ regions, are presented in \\S5. Results and discussion are presented in \\S6 and \\S7 respectively. ",
        "conclusions": "The results of this analysis provide the strongest evidence to date for a non-negligible dependence of Cepheid distances on metal abundance. Our best estimate of the magnitude of this dependence is $\\gamma = -0.24 \\pm 0.05$ mag~dex$^{-1}$, when referenced to the Zaritsky et al.\\,(1994) HII region metallicity scale (we consider the effects of adopting a different metal abundance scale below).  This result is consistent with the dependence measured from a direct comparison of  metal-rich and metal-poor Cepheid fields in M101 ($\\gamma = -0.24 \\pm 0.16$; K98). In the remainder of this section, we explore the consequences of such a $Z$-dependence on the calibration of the distance scale as a whole and H$_0$.  The consequences of a Cepheid metallicity dependence of roughly this magnitude on the calibration of several extragalactic standard candles was explored in detail in the final series of papers from the HST H$_0$ Key Project (Sakai et al.\\,2000, Ferrarese et al.\\,2000b; Gibson et al.\\,2000; Kelson et al.\\,2000; Mould et al.\\,2000; F01). We have summarized these results in Table~6, which shows the net effect of a Cepheid metallicity dependence of 0.20 mag~dex$^{-1}$ on the zeropoint calibrations of the secondary distance indicators used by the Key Project team.  These are expressed in terms of the luminosity zeropoints and on the mean net change in the derived distances for the Key Project samples.  A Cepheid $Z$-dependence in the direction measured here causes {\\it all} of the secondary distance scales to be systematically underestimated (thus leading to an over-estimate of H$_0$).  This is because the PL relation is calibrated with a relatively metal-poor galaxy, the LMC.  The magnitude of the effect is slightly different for the different secondary distance indicators, but for $\\gamma = -0.20$ mag~dex$^{-1}$ it is significant but small, lowering the net value of H$_0$ by 3.5\\%, or about 2.5 km\\,s$^{-1}$\\,Mpc$^{-1}$ for H$_0$ = 72 km\\,s$^{-1}$\\,Mpc$^{-1}$ (F01). This correction already has been incorporated into the value given above.  As mentioned in \\S\\,6 the absolute slope of the Cepheid $Z$-dependence is also sensitive to the metallicity scale adopted.  As an illustration of this point Figure~\\ref{figure:znew} shows the same Cepheid vs TRGB comparison as Figure~\\ref{figure:metaldep}, but with the metal abundances adjusted to agree with the electron temperature based HII region abundances in Kennicutt et al.\\,(2003).  As discussed earlier this has the effect of preferentially reducing the metallicities of the most metal-rich Cepheid fields, and the result is a somewhat ($\\sim$25\\%) steeper $Z$-dependences, with an average $\\gamma = -0.31 \\pm 0.09$ mag~dex$^{-1}$. Note however that adopting this different abundance scale {\\it would have an identical effect on the distance scale} as given in Table~6, because the effect of the steeper $Z$-dependence would be canceled by a correspondingly narrower abundance range in the calibrating galaxies; in other words as long as the metallicity corrections are applied using a consistent abundance scale, the precise calibration of the metallicity scale is not important.  Of course the absolute slope of the dependece is important for understanding the physical origins of the period-luminosity dependence of the Cepheid variable stars.  In \\S3, it was suggested that because the study presented in this paper is a {\\it differential} test, it would not matter which TRGB calibration is used.  We test this assumption by examining the metallicity dependence using two independent calibrations. The first one is that by Lee et al. (1993) which was used throughout this paper. The second calibration is that by Salaris \\& Cassisi (1998), which is based on the stellar evolution models.  The dominant different between the two calibration is that the theoretical model by Salari \\& Cassisi predicts the TRGB magnitude $\\sim 0.1$ mag brighter than the empirical calibration of Lee et al. The authors suggest that the difference arises from the fact that the globular cluster samples used in the empirical calibration may be missing the brightest RGB stars due to the small number statistics, and thus systematically dimming the TRGB magnitude. In Figure~\\ref{figure:zpcomp},  we show two correlations, one using the Lee et al. calibration, and the other based on Salaris \\& Cassisi 1998. As expected, the zero points of the two correlations differ by $\\sim$ 0.1 mag. For the MF91, multi-wavelength sample, using the theoretical calibration, we obtain $\\gamma = -0.26 \\pm 0.08$, which agrees well with the fit using the empirical, Lee et al. calibration, $\\gamma = -0.23 \\pm 0.08$. In summary, we emphasize again that the results shown in this paper are based on {\\it differential} tests, and as indicated by our simple comparison, the value of $\\gamma$ should not be affected by the use of another TRGB calibration.  Finally, our measurment of the metallicity dependence cannot distinguish between a variation in the zero point or in the slope. As discussed in Section~4, the slope of the Cepheid PL relation is not always well determined; the MF91 and U99 calibrations in fact yield I-band slopes that are significantly different from each other. Thus, there is a need to check if the slope is the cause of the metallicity dependence of the Cepheid variables. A detailed study is beyond the scope of this paper; here a simple exercise is carried out to examine the effect of the slope, by calculating the mean period of Cepheid variable sample for each galaxy.  When the mean periods are plotted against the metallicities, we find that there are two groupings: one around $12+log(O/H) \\sim 8.7$ and mean $P \\sim 1.4$, and the other one at $12+ \\log (O/H) \\sim 7.7$ and mean $P \\sim 1.1$.  That is, one group at high Z corresponds to the longer mean period, and the second one at low Z corresponds to the shorter mean period. The second low-Z, shorter period group consists of four galaxies. The mean Cepheid periods were also calculated for all the galaxies used as the calibrators for the Tully-Fisher relation. This is especially important to check if slope slope is responsible for the metallicity dependence of the Cepheids, affecting the calibration of the secondary distance indicators and finally the value of H$_0$. The Tully-Fisher calibrators all lie in the high-Z, long mean period group. If the metallicity affects the slope of the Cepheid PL relation, then we might need to exclude those galaxies whose mean period is significantly different from others. Thus, excluding four galaxies that have low mean periods, the metallicity dependence, $\\gamma$, was re-calculated. For the multiple-wavelength, MF91 calibration sample, $\\gamma = -0.25 \\pm 0.09$, which agrees well with the value estimated using all galaxies ($\\gamma=-0.24$). Therefore, to first order, the slope of the Cepheid PL relation does not appear to affect the metallicity dependence."
    },
    "0402/astro-ph0402516_arXiv.txt": {
        "abstract": "Galaxies are lighthouses that sit atop peaks in the density field. There is good observational evidence that these lighthouses do not provide a uniform description of the distribution of dark matter. ",
        "introduction": "In the beginning there was dark matter and gas but there were no stars. Today some of the gas that was once dispersed has transformed into stars, and we have light.  In places the stars are young, hot, and bright while elsewhere the stars are old, cool and faint.  In places the stars have been scattered by violent collisions.  In places the gas never coalesced and stars never formed.  We have light, but light in selective places.  It has been pointed out that the abrupt cutoff of the luminosity function at the bright end and the flat slope at the faint end compared with the Press-Schechter mass function anticipated by the hierarchical clustering scenario suggests that there is a relative deficiency of light at high and low mass extremes compared with the situation at intermediate masses (Yang et al. 2003, van den Bosch et al. 2003).  Semi-analytic models that follow the transformation of gas into stars lead to similar expectations (Blanton et al. 1999, Somerville et al. 2001, Ostriker et al. 2003).  These same results are found directly in observations. A full discussion of the observational situation is provided by Tully (2004).  A condensed version is presented here. ",
        "conclusions": ""
    },
    "0402/astro-ph0402270_arXiv.txt": {
        "abstract": "{We report the discovery of one unique cataclysmic variable drawn from the Hamburg Quasar Survey, HS\\,2331+3905. Follow-up observations obtained over three years unveiled a very unusual picture. The large amplitude 3.5\\,h radial velocity variations obtained from our optical spectroscopy is not the orbital period of the system, as one would normally expect. Instead, extensive CCD photometry strongly suggests that HS\\,2331+3905 is a short orbital period cataclysmic variable with $P_{\\rm orb}= 81.09$\\,min, containing a cold white dwarf which appears to exhibit ZZ\\,Ceti pulsations.}         \\addkeyword{Stars: binaries:close} \\addkeyword{Stars: Cataclysmic Variables} \\addkeyword{Stars: individual: HS\\,2331+3905}   \\begin{document} ",
        "introduction": "HS\\,2331+3905 (HS\\,2331 thereafter) was selected as a cataclysmic variable (CV) candidate on the basis of its spectral characteristics in the Hamburg Quasar Survey (HQS; Hagen et al.~1995; G\\\"ansicke et al. 2002). The identification spectrum of HS\\,2331 contains broad double-peaked Balmer emission lines, clear signs of the presence of an accretion disc, flanked by extremely broad absorptions throughs, indicating that this CV contains a relatively cold white dwarf. The red part of the spectrum does not contain any spectral features that could be ascribed to the emission of the secondary. Here we report follow-up (ground and space) photometry and spectroscopy of HS\\,2331, obtained over a three year period after its identification.  \\begin{figure*}[!t] \\hfill \\includegraphics[width=5.6cm, angle=270]{saraujo_fig1.ps}% \\hfill \\includegraphics[width=5.6cm, angle = 270]{saraujo_fig2.ps} \\hfill \\caption{\\textit{Left panel:} Samples of the light curves of HS\\,2331 obtained from differential CCD photometry. The names and numbers indicate the observatory and time resolution used in each of the observations respectively. \\textit{Right panel:} Combination of $FUV$, optical spectra and 2MASS colours of HS2331 (dark line and filled circles) plotted with the best three-component model fit (grey line and open circles). See text for details.} \\label{f-fig1} \\end{figure*} ",
        "conclusions": "We have discovered a short orbital period system, HS\\,2331, as part of the HQS quest for new CVs. The orbital period of HS\\,2331, $P_{\\rm orb}=81.09$\\,min, was primarily defined by the detection of coherent eclipses. From our three years of photometric data, HS\\,2331 appears to be a permanent superhumper with $P_{\\rm SH}=83.38$\\,min. The light curves of HS\\,2331 display double-humps with a period that is exactly half the orbital period (evident from a direct inspection of the light curves in the left panel of Fig.\\,\\ref{f-fig1}), suggesting that we are seeing some sort of symmetric structure, such as e.g. two bright spots.  In addition, HS\\,2331, exhibits the photometric behaviour typical of ZZ\\,Ceti pulsators, showing multifrequency variability in the range $\\sim 60$\\,s to $\\sim300$\\,s. The white dwarf temperature derived from our fit, 11\\,000\\,K, is well within the instability strip for ZZ\\,Ceti pulsators.  In order to disentangle the multiperiodic signature of the likely white dwarf pulsator in HS\\,2331, we need to organize a multi-site observing campaign to obtain long, continuous stretches of high time resolution photometry.  All in all, the pieces of the jigsaw seems to come together, and we are beginning to understand the nature of HS\\,2331. There are nevertheless, several points that we still need to address. The fact that the dominant radial velocity variability does not correspond to the orbital period of the system is particularly disconcerting. At present, we have no explanation for this phenomenon, nor for the physical significance of the 3.5\\,hr radial velocity period. We are not aware of any other system suffering from this problem, but the reason for this may be that periods determined from radial velocity studies are usually adopted unquestioned as reflecting the corresponding orbital periods."
    },
    "0402/astro-ph0402046_arXiv.txt": {
        "abstract": "It was predicted more than 40 years ago that the cores of the coolest white dwarf stars should eventually crystallize. This effect is one of the largest sources of uncertainty in white dwarf cooling models, which are now routinely used to estimate the ages of stellar populations in both the Galactic disk and the halo. We are attempting to minimize this source of uncertainty by calibrating the models, using observations of pulsating white dwarfs. In a typical mass white dwarf model, crystallization does not begin until the surface temperature reaches 6000-8000 K.  In more massive white dwarf models the effect begins at higher surface temperatures, where pulsations are observed in the ZZ Ceti (DAV) stars. We use the observed pulsation periods of \\object{BPM~37093}, the most massive DAV white dwarf presently known, to probe the interior and determine the size of the crystallized core empirically. Our initial exploration of the models strongly suggests the presence of a solid core containing about 90\\% of the stellar mass, which is consistent with our theoretical expectations. ",
        "introduction": "More than four decades have passed since \\cite{abr60}, \\cite{kir60} and \\cite{sal61} predicted that the cores of white dwarf stars should crystallize as they cool down over time. There has never been a direct empirical test of this theory. The discovery of pulsations in the massive hydrogen-atmosphere (DA) white dwarf \\object{BPM~37093} \\citep{kan92} provided the first opportunity to search for the observational signature of crystallization in an individual star. Theoretical calculations by \\cite{win97} and \\cite{mw99} suggested that the core of this star might be up to 90\\% crystallized, depending on its mass and internal composition.  In addition to providing the first test of the theory of crystallization in a dense stellar plasma, knowing whether and to what degree this star is crystallized has implications for a more fundamental question. Recent {\\it Hubble Space Telescope} observations of the faintest white dwarfs in the globular cluster M4 \\citep{han02} have led to a resurgence of interest in using these stars to provide independent constraints on the ages of stellar populations. The largest potential sources of error in this method arise from uncertainties about the composition and structure of white dwarf interiors. Fortunately, all of the major uncertainties can be minimized through detailed observation and modeling of pulsating white dwarfs, providing a crucial method to probe the stellar interiors and calibrate the cooling models.  The crystallization process leads to one of the largest sources of uncertainty in the ages of cool white dwarfs \\citep{seg94}. When a typical mass white dwarf star \\citep[$0.6~M_\\odot$,][]{ngs99} cools down to $T_{\\rm eff}\\sim 6000$-8000 K (depending on the core composition), the high-density core will undergo a phase transition from liquid to solid. An associated latent heat of crystallization will be released, providing a new source of thermal energy that introduces a delay in the gradual cooling of the star \\cite[for a recent review, see][]{hl03}. In mixed C/O cores, phase separation of the ions during crystallization can provide an additional source of energy, delaying the cooling even further.  Early attempts to determine the crystallized mass fraction in \\object{BPM~37093} were plagued by difficulties with uniqueness \\citep{mw99}. Recent improvements in our ability to match the observed pulsation periods in white dwarf stars with theoretical models have been driven by the development of an optimization method based on a parallel genetic algorithm \\citep{mc03}. This method allows the objective global exploration of the defining parameters, which is essential to ensure a unique solution. In this Letter, we present the initial application of this method to fit the observed pulsation periods of \\object{BPM~37093} with asteroseismological models. ",
        "conclusions": ""
    },
    "0402/astro-ph0402336_arXiv.txt": {
        "abstract": "{ We present an analysis of the complex associated system of the high-redshift QSO \\object{HE~2347-4342}. Absorption features of \\ion{H}{i}, \\ion{C}{iii}, \\ion{C}{iv}, \\ion{N}{v}, and \\ion{O}{vi} with up to 16 components occur in the optical spectral range located up to $1500\\,\\mathrm{km\\,s}^{-1}$ redwards from the emission line. Apparently, \\ion{C}{iv} and \\ion{N}{v} show the line locking effect. A quantivative analysis of the line distribution comparing simulated spectra with randomly distributed doublets reveals, however, no statistical evidence for its physical reality. Using photoionization calculations to emulate the observed ion column densities we constrain the quasar's spectral energy distribution. Absorbers in the velocity range of $200 - 600\\,\\mathrm{km\\,s}^{-1}$ can be modelled successfully with a spectral index of $\\alpha \\sim -3$ at energies higher than $3 - 4\\,\\mathrm{Ryd}$, which is an energy distribution similar to the QSO continuum suggested by Mathews \\& Ferland (\\cite{mathewsferland}). The analysis of a group of high velocity absorbers ($v > 1300\\,\\mathrm{km\\,s}^{-1}$) leads to a harder energy distribution. The large amount of helium ($\\log N_{\\mathrm{\\ion{He}{ii}}} > 16.3$) associated with these absorbers implies that they are responsible for the observed absence of the proximity effect (Reimers et al. \\cite{reimers97}). Clouds located more distant from the quasar may be shielded from the high energy part of the quasar continuum due to optically thick absorption shortward of $228\\,\\mathrm{\\AA}$ by the high velocity absorbers. A group of absorbers with $900 < v < 1200\\,\\mathrm{km\\,s}^{-1}$, in particular a cloud at $1033\\,\\mathrm{km\\,s}^{-1}$, which has the most reliable column density measurements, can be modelled neither with  photoionzation nor under the assumption of collisionally ionized gas. Possible explanations are a multiphase medium with a mixture of photo and collisionally ionized gas and/or gas in non-equilibrium. ",
        "introduction": "Absorption complexes with $z_{\\mathrm{abs}} \\sim z_{\\mathrm{em}}$ are apparently associated with the QSO. These so-called associated absorption systems are characterized by highly ionized absorption lines and are generally defined by the criterion that the velocity relative to the quasar is less than $5000\\,\\mathrm{km\\,s}^{-1}$ (Weymann et al. 1979, Foltz et al. 1986, 1988). Obviously, the absorbing material belongs to the inner region of active galaxies. Studies of associated systems have potentially important implications for galaxy formation and evolution. The gas dynamics and velocity fields of the flows in the central region are still a matter of debate. The metalicities of these systems are solar to several times solar (Wampler et al. 1993). High metalicities would agree with predictions of galactic chemical evolution (Hamann \\& Ferland 1993).  In contrast to normal intergalactic absorption line systems, which are formed in intervening gas clouds at distances corresponding to their cosmological redshifts, associated systems typically show strong high ionization lines from \\ion{O}{vi} and in particular \\ion{N}{v} which are rarely detected in intergalactic absorption line systems. Obviously, the origin of the high ionization is the hard, nonthermal EUV radiation of the parent QSO. Therefore, column density ratios of \\ion{C}{iii}, \\ion{C}{iv}, \\ion{Si}{iv}, \\ion{O}{vi}, and \\ion{N}{v} combined with photoionization calculation should -- at least in principle -- allow to estimate the EUV energy distribution of the QSO which is otherwise largely unobservable.  The high-redshift quasar HE~2347-4342 ($z_{\\mathrm{em}} = 2.885, V = 16.1$) was discovered in the course of the Hamburg/ESO Survey (HES; Reimers \\& Wisotzki \\cite{reimerswisotzki}). A strong associated system with absorption components up to $1500\\,\\mathrm{km\\,s}^{-1}$ redwards from the QSO emission line redshift is observed in highly ionized metal lines. In this paper we present a detailed analysis of the absorption characteristics and the inferences on the spectral energy distribution of the QSO.  In absorbing gas clouds close to the QSO, radiation pressure due to absorption by strong resonance lines and therefore line locking can play a role, and this appears to happen in the complex associated system of HE~2347-4342. Line locking occurs if the velocity separation between two absorbing clouds is equal to the velocity separation of a doublet splitting (Scargle \\cite{scargle}; Braun \\& Milgrom \\cite{braunmilgrom}). Aside from the incidence in stars this effect has been observed mainly in BAL--QSOs (Foltz et al. \\cite{foltzetal}; Vilkoviskij \\& Irwin \\cite{vilkoviskij}). However, it also occurs in other AGNs with associated systems (e.g. Srianand \\cite{srianand}; Srianand et al. \\cite{srianandetal}). With our analysis of the \\ion{C}{iv} and \\ion{N}{v} complex associated system of HE~2347-4342 we shall perform a statistical test in order to clarify whether the apparent line locking has a physical basis or the line coincidences are within the expectation in a random distribution of a rich associated system with up to 16 components.  The strong associated system of HE~2347-4342 also seems to be responsible for the observed absence of a \\ion{He}{ii} proximity effect in this QSO (Reimers et al. \\cite{reimers97}) which would imply that at least one of the components of the associated system is optically thick to radiation shortward of $228\\,\\mathrm{\\AA}$. If the corresponding clouds were closer to the QSO than other components, the radiation shortward of $228\\,\\mathrm{\\AA}$ would be invisible for the more distant absorbers. This could be observable and our analysis attempts to reveal whether some subcomponents ``see'' a softer ionizing radiation field than others. ",
        "conclusions": "We have analyzed the complex associated system of HE~2347-4342 showing absorption features of \\ion{H}{i}, \\ion{C}{iii}, \\ion{C}{iv}, \\ion{N}{v}, and \\ion{O}{vi} in the optical. The state of line locking apparently visible in \\ion{C}{iv} and \\ion{N}{v} was investigated. Comparing simulated spectra with randomly distributed absorption doublets we found no statistical evidence for the presence of this effect. We point out that apparently present line locking has to be confirmed quantitatively.  We have used column densities of visible ions and upper limits of non-detected ions redwards from the QSO emission line redshift to model its SED adopting a simple model matching the observed low energy part. Applying our method to the data leads to widely consistent results. Five absorbers within the velocity range $200 < v < 600\\, \\mathrm{km\\,s}^{-1}$ can be modelled with a confidence of at least 90\\,\\% with a spectral slope $\\alpha \\sim -3$ at the energy $\\sim 3 \\dots4$\\,Ryd resulting in a SED similar to a Mathews--Ferland continuum. Two other absorbers at 850 and $1259\\, \\mathrm{km\\,s}^{-1}$ satisfy similar models with a steeper slope (99\\,\\% confidence, absorber 6) or a higher turnover energy (47\\,\\% confidence, absorber 10), respectively. Thus, the velocities of the absorbers and their distance to the quasar seem to show no correlation.  We failed to find parameter combinations to describe the absorption features in the range $900 < v < 1200\\, \\mathrm{km\\,s}^{-1}$. We conclude that these absorbers are either not purely photoionized or the effective ionizing continuum is not compatible with our model assumptions. Computations with the MF continuum indicate a higher metalicity than solar with reduced silicon abundance. Differential abundance effects are not considered within our simple model, since additional parameters would cause a higher degree of freedom and increase the uncertainties.  Collisional ionization at a single gas temperature could not be confirmed as dominating process for these absorbers. We found that the absorber at $1033\\,\\mathrm{km\\,s}^{-1}$ may be described with a mixture of temperatures of $100\\,000 < T < 300\\,000\\,\\mathrm{K}$. However, the presented data are not sufficient to constrain multiphase gas models.  A further complication is the possible variability of the UV continua of luminous quasars in strength and shape. It has recently been discovered that the UV continuum of the similar luminous QSO HS~1700+6416 ($z = 2.73$) has varied by a factor of 3 to 4 at $1200\\,\\mathrm{\\AA}$, while in the optical the amplitude is only $0.1^{\\mathrm{m}}$ (Reimers et al., in preparation). Consequently, the shape of the observed UV continuum (Fig. \\ref{sed}) might be different from the one responsible for photoionization of the associated system. However, new observations made in June 2003 with the UVES spectrograph with the same resolution but an improved signal-to-noise ratio reveal no variability of the observed components of \\ion{O}{vi} and \\ion{N}{v}. The time interval between the observations analyzed in this paper and the new data is about 2.5 yr, corresponding to about 8 months in the QSO rest frame.  Modelling the two high velocity absorbers ($1300 < v < 1600\\,\\mathrm{km\\,s}^{-1}$) we found best results with flat slopes at high turnover energies. However, an energy distribution like this is in conflict with the non-detection of HE~2347-4342 in the ROSAT all sky survey as well as the non-detection of \\ion{Mg}{viii} in the FOS portion of the spectrum. Assuming the high velocity absorbers are the closest to the quasar they are expected to be illuminated by a hard UV continuum. We confirm a spectral energy distribution harder than the MF continuum. Since the two high velocity absorbers lead to the hardest continuum and large \\ion{He}{ii} column densities ($\\log N_{\\mathrm{\\ion{He}{ii}}} > 16.3$), they may shield the other absorbers from the QSO radiation. If they are optically thick in the \\ion{He}{II} continuum, lower velocity absorbers are exposed to a filtered, softer radiation field, as observed.  The results could be improved if more ions were included in the analysis. Therefore, column densities of ions observed in the UV portion of the spectrum like \\ion{O}{v} have to be estimated more accurately. In the available FOS and GHRS spectra the resolution and $S/N$ do not allow to determine column densities for individual absorbers. Especially, the models for the high velocity clouds which suffer from saturated hydrogen and uncertain \\ion{O}{vi} absorption features could be improved, if future UV observations allow to resolve further lines like \\ion{O}{v} 630\\,\\AA, \\ion{O}{iv} 608\\,$\\mathrm{\\AA}$ etc.  In order to constrain the high energy portion of the SED more precisely, several ionizations stages of neon (\\ion{Ne}{v} -- \\ion{Ne}{viii}) could be used, which are expected to be visible in the UV spectral range. If these lines are observed with a sufficiently high resolution and signal-to-noise ratio, e.g. by the Cosmic Origin Spectrograph (COS) on HST, the energy range covered by the ionization potentials of the observed ions will be extended by a factor of 2. Thus, the high energy portion of the SED could be modelled more confidently."
    },
    "0402/astro-ph0402100_arXiv.txt": {
        "abstract": "{We present a near-infrared J and K photometric catalog containing more than 73,000 stars in the central region of the giant Globular Cluster $\\omega$~Centauri. This is the largest IR data-set ever published for this cluster and has been used to completely characterize  the morphology and the properties of the Red Giant Branch (RGB). In particular, we concentrated our attention on (i) the anomalous RGB (RGB-a), recently discovered in this cluster and (ii) the RGB of the dominant metal poor population (RGB-MP) in both the infrared $(K,J-K)$ and optical-infrared $(K,V-K)$ color magnitude diagrams. The full set of morphological parameters and photometric indices has been measured and compared with the empirical relations by Ferraro et al. (2000). We find that the detailed photometric properties of the RGB-a are in full agreement with the recent spectroscopic metallicity estimates, that place it at the metal-rich extreme of the stellar population mix in $\\omega$~Centauri. ",
        "introduction": "The origin and star formation history in $\\omega$~Centauri, the most luminous and massive globular cluster in our Galaxy, is one of the most intriguing problems of modern stellar astrophysics. $\\omega$~Centauri is the only known Galactic globular cluster which shows clear variations in the metal content of its giants. This evidence has been firmly estabilished in the past by extensive low (Norris et al. 1996, Suntzeff \\& Kraft 1996) and high resolution (Norris \\& Da Costa 1995, Smith et al. 2000) spectroscopic surveys. More recently, the scenario has become more complicated due to the discovery of an additional, metal-rich population with its own distinct RGB (hereafter RGB-a) that contains approximately 5$\\%$ of the red giants in $\\omega$~Cen (Lee et al. 1999, Pancino et al. 2000, 2002). In spite of the huge observational effort carried out so far, the global picture of the cluster formation and evolution is far from being completely understood.  In this framework, we have started a long--term project devoted to the detailed study of the properties of the different stellar populations in this cluster (see the overview of the project by Ferraro, Pancino \\& Bellazzini, 2001). Within this project a number of results have been published, in particular on the identification of the anomalous RGB-a, and on the definition of its chemical and kinematic properties (see Pancino et al. 2000, 2002, 2003; Bellazzini et al. 2001; Ferraro et al. 2002; Origlia et al. 2002). Most of the actual observational knowledge comes from the optical (photometric or spectroscopic) study of RGB stars. Only a few sparse literature of near infrared observations existed up to now. Two pioneering studies by Glass \\& Feast (1973, 1977) and Persson et al. (1980) present J, H and K magnitudes for a few tens of bright giants. More recently, the NICMOS camera on board of HST has been used by Pulone et al. (1998) to obtain extremely deep photometry of a tiny area ($20\"\\times20\"$) 7 arcminutes away from the cluster center. The 2MASS survey has instead covered a very wide area ($3^{\\circ}\\times2^{\\circ}$) around $\\omega$~Cen, which however does not include the central region of the cluster. This paper presents a large J and K photometric catalog of more than 73,000 stars in an area covering $\\sim 13'\\times13'$ around the center of $\\omega$~Cen. By combining the IR-data set with wide field optical photometry (Pancino et al. 2000, 2003), we measured the complete set of morphological parameters defined by Ferraro et al. (2000, hereafter F00), which fully characterize the photometric properties of the RGB of different populations in $\\omega$~Centauri. ",
        "conclusions": "We presented an extensive near IR, J and K catalog of stars in the giant globular cluster $\\omega$~Cen. More than 73,000 stars have been measured allowing an accurate photometric characterization of the RGB. In particular, the colors at different magnitude levels, the magnitude at different colors, the RGB slope and the RGB bump position have been measured, scaled to the absolute plane and compared to similar features measured in clusters with different metallicity by F00. The agreement with the F00 relations is quite good, and the photometric properties of the newly discovered anomalous RGB (RGB-a) consistently reflect the high metal content of this sub-population found by previous spectroscopic work."
    },
    "0402/hep-ph0402059_arXiv.txt": {
        "abstract": "{ \\normalsize A pseudo Nambu-Goldstone boson as curvaton avoids the $\\eta$-problem of inflation which plagues most curvaton candidates. We point out that a concrete realization of the curvaton mechanism with a pseudo Nambu-Goldstone boson can be found in the supersymmetric Peccei-Quinn mechanism resolving the strong CP problem. In the flaton models of Peccei-Quinn symmetry breaking, the angular degree of freedom associated with the QCD axion can naturally be a flat direction during inflation and provides successful curvature perturbations. In this scheme, the preferred values of the axion scale and the Hubble parameter during inflation turn out to be about $10^{10}$ GeV and $10^{12}$ GeV, respectively. Moreover, it is found that a significant isocurvature component, (anti)correlated to the overall curvature perturbation, can be generated, which is a smoking-gun for the curvaton scenario. Finally, non-Gaussianity in the perturbation spectrum at potentially observable level is also possible. }    \\thispagestyle{empty}  \\end{titlepage} ",
        "introduction": "The primordial curvature perturbation is caused presumably by some scalar field, which acquires its perturbation during inflation. For a long time it was generally agreed that the curvature-generating field would be the inflaton \\cite{treview}. Then it was suggested  instead  that this field is some other `curvaton' field \\cite{LW,moroi} (see also \\cite{also}), which generates the curvature perturbation only when its  density becomes a significant fraction of the total. The curvaton proposal has received enormous attention because it opens up new possibilities both for model-building and for observation.\\footnote {More recently still it has been suggested that the curvaton acts by causing inhomogeneous reheating \\cite{dzlev} or through a preheating mechanism \\cite{steve}.}  In all cases,  the field responsible for the curvature perturbation must be  light during inflation, in the sense that its effective mass $\\meff$ is much less than the Hubble parameter $H$. (To be precise, we need during inflation $\\meffs \\lsim 0.1 H^2$, so that the spectral tilt $1-n\\simeq \\frac23 \\meffs/H^2$ satisfies the observational constraint.) If the responsible field is a curvaton, it should preferably also remain light after inflation, until $H$ falls below the true mass \\cite{CD}. These requirements are in mild conflict with the generic expectation from supergravity, that the mass-squared of each scalar field  will be at least of order $H^2$ \\cite{DRT}. This is the famous $\\eta$-problem of inflation. Ways have been proposed to keep the responsible  field sufficiently light \\cite{treview}, the most straightforward of them being to make it a pseudo Nambu-Goldstone boson (PNGB) (for the curvaton see \\cite{LW,pngb}).  For economy, and also to facilitate contact with observation, one would like a candidate for the responsible field to be one that is present in an already-existing model, designed for some purpose other than the generation of  the curvature perturbation. Several such candidates have been proposed \\cite{pngb,andmore,pqcurv,kkt,McD}, but in general they do not come with a satisfactory mechanism for keeping the curvaton sufficiently light. {\\em The purpose of this paper is to suggest a natural curvaton candidate, which is present in flaton models of Peccei-Quinn symmetry breaking, and which is a PNGB.}  The Peccei-Quinn (PQ) symmetry  provides a nice solution to the strong CP problem \\cite{pq}. It  is  an anomalous global symmetry, $U(1)_{PQ}$, spontaneously broken at an intermediate scale $f_{PQ}$ with preferred value around $10^{12}$ GeV providing enough dark matter (axion) of the universe. Its PNGB is the axion, which must be extremely light to satisfy the observational constraints \\cite{INV}. In the context of supersymmetry (SUSY), two complex fields (or more) are needed to implement a global symmetry due to holomorphicity of the superpotential (unless the symmetry in question is R--symmetry). Hence, in SUSY, to spontaneously break the  PQ symmetry one needs at least two complex fields, corresponding to two radial and two real angular degrees of freedom. One of the latter is the axion, and the other is our curvaton candidate.\\footnote {By `degree of freedom' we mean as usual a normal mode of the coupled oscillations of the fields, which after quantization corresponds to a particle species.} We can define a symmetry (explicitly broken of course) acting only on the combination of phases which corresponds to the curvaton, and then the curvaton is the  PNGB of that symmetry.  To understand how our model works, recall first that two  fundamentally different paradigms exist for the PQ symmetry breaking fields. In one of them,  the potentials involve only renormalizable terms in the superpotential plus small soft SUSY breaking terms, leading to a normal Mexican-hat potential whose height and width have the same  scale $f_{PQ}$. According to this paradigm, PQ symmetry breaking persists in the limit of unbroken SUSY so that the axion has a well-defined supersymmetric scalar partner called the saxion (as well as a fermionic partner called the axino). The saxion is the only light degree of freedom apart from the axion, its mass coming from SUSY breaking and being of order $\\TeV$ for gravity-mediated SUSY breaking. The other two degrees of freedom are heavy with mass of order $f_{PQ}$.  According to the other  paradigm, the PQ symmetry breaking fields are instead flaton fields \\cite{flaton}, so-called because they are symmetry-breaking fields, which correspond to flat directions of the potential (i.e., directions in which the quartic term is negligible). Their  potential contains only soft SUSY breaking terms and non-renormalizable terms, which means that it is very flat. In the limit of unbroken SUSY there would be no spontaneous breaking of the PQ symmetry. A nice feature of this paradigm is that the intermediate axion scale $f_{PQ}$ is not put by hand but is generated through a geometric mean of the SUSY breaking scale and the Planck scale. The other degrees of freedom accompanying the axion  get their mass only from SUSY breaking, making them of order $\\TeV$ for gravity-mediated SUSY breaking. In this paper, we point out that this kind of model contains a natural curvaton candidate as the angular degree of freedom. The radial degrees of freedom are less suitable for this purpose, because in the early Universe they presumably acquire the mass of order $H$ which is generic \\cite{DRT} for scalar fields in a supergravity theory. In contrast, a mass of order $H$ is unlikely to be generated for the angular degree of freedom, because it would correspond to a generalized $A$-term which is forbidden if the fields responsible for the energy density are charged under certain symmetries \\cite{DRT}. We will analyze how the model parameters like the axion scale, the Hubble parameter and the curvaton decay temperature are constrained to produce successful curvature perturbations. ",
        "conclusions": "Within the curvaton paradigm, we have proposed a new candidate for the field which causes the curvature perturbation. We believe that it is one of the most attractive  candidates yet   proposed for that field. Being a field that is part of the flaton realization of PQ symmetry-breaking, and it can easily be kept light during and after inflation. In a typical part of the universe, the observed magnitude $\\sim 10^{-5}$ of the curvature perturbation comes  from a  modest and quite reasonable hierarchy between the inflationary Hubble parameter and the vacuum expectation value  of the Peccei-Quinn field. The predicted values of the axion scale and the Hubble parameters are $10^{10}$ and $10^{12}$ GeV, respectively, corresponding to $n=1$.  In this case, the entropy dumping due to the curvaton decay is negligible so that the conventional cosmological predictions concerning  dark matter components and  baryogenesis remain unchanged.  The model can generate two kinds of isocurvature pertubations. One is an axion isocurvature perturbation, uncorrelated with the curvature perturbation.  The other is an isocurvature perturbation in some other kind of dark matter, or in the baryonic matter, which is produced either  before or at curvaton decay. As described in \\cite{iso1,iso2}, such a perturbation is generic to curvaton models, and is fully correlated or anti-correlated with the curvature perturbation.\\footnote{Correlation or anti-correlation of the CDM \\{baryon\\} isocurvature perturbation depends on whether CDM creation  \\{baryogenesis\\} takes place (just) before curvaton decay or due to the curvaton decay itself. Note, that there is no residual isocurvature perturbation if CDM creation \\{baryogenesis\\} occurs after curvaton decay.} The detection of any such isocurvature perturbation would be evidence in support of the curvaton model.  Finally, we have shown that, in this case, non-Gaussianity at at an observable level (with \\mbox{$f_{\\rm NL}\\sim {\\cal O}(10)$}) is possible.    Our  curvaton candidate is  a PNGB, corresponding to an angular degree of freedom associated with the QCD axion. As a consequence, it can avoid the usual mass of order $H$ during (and after) inflation, because supersymmetry breaking affects the mass only through $A$ terms, which can be controlled  by appropriate symmetries.  Our model of PQ symmetry breaking is similar to  one already extensively investigated \\cite{comelli,hangbae}, except that we take the PQ symmetry to be spontaneously broken throughout the history of the Universe. In contrast, the investigations of \\cite{comelli,hangbae} assume that the symmetry is initially unbroken, leading to thermal inflation. These possibilities correspond respectively to radial masses-squared of order $\\mp H^2$, and the two signs should be deemed equally likely in the absence so far of a string-theoretic prediction. The unbroken paradigm has a very different cosmology from the one that we are adopting, producing in particular copious saxion- or axino-like particles which may  decay into relativistic axions whose energy density is enough to affect nucleosynthesis \\cite{comelli}. Also, the lightest axino-like particle may be the lightest supersymmetric particle \\cite{hangbae}, providing a more natural implementation of the Cold Dark Matter scenario that was originally  proposed \\cite{leszek} in the context of non-flaton models. None of this occurs within our paradigm.  As it invokes the axion, our model, as it stands, is open to the criticism that the axion mass is implausibly small. Indeed, in the kind of models that we have considered where the axion is an angular part of a complex field, a non-renormalizable term  in the potential  with dimension $d$ generically breaks PQ symmetry and  contributes to the axion mass an amount of order $\\sim v^{(d-2)/2}$ in Planck units. To keep the axion mass to the required value of order $10^{-30}\\mpl$, terms up to $d\\sim 12$ must  respect the PQ symmetry. (We take $v\\sim f_{PQ} = 10^{12}\\GeV$ for an estimate.) A widely-discussed possibility to ensure this is to impose a $Z_n$ subgroup of the PQ symmetry \\cite{lukas}.\\footnote {Imposing the full PQ symmetry of course works, but exact continuous global symmetries are widely regarded as incompatible with string theory and even the existence of gravity. A different  possibility, which does not seem to have been  mentioned before, might be to suppose that the PQ fields are actually moduli, making the origin a point of enhanced symmetry. Then the non-renormalizable terms may be suppressed by a factor $(\\TeV/\\mpl)^2 \\sim 10^{-30}$. However this is small enough for all $d\\geq 5$ only if $f_{PQ}$ has an  implausibly small value of order $10^8\\GeV$. In any case, this mechanism cannot be used in the flaton case, which invokes an unsuppressed coefficient for one term.}  Because of these considerations, one may favour a solution of the CP problem in which the QCD axion is identified with a string axion \\cite{bdaxion}, or one without any axion at all.\\footnote {The string axion as a curvaton candidate is discussed in \\cite{pngb}. As the string axion is a PNGB one may be hopeful that it can be kept sufficiently light in the early Universe,  but string phenomenology is not yet sufficiently developed that one can be sure.} In that case the field considered in our proposal becomes ad hoc, introduced solely to explain the curvature perturbation. Still, because of its  simplicity and the ease with which it is made  sufficiently light, we feel that the present curvaton model is extremely attractive in comparison with others. Let us end by mentioning those   other models which have a close connection with ours.  The closely  related  models are considered in Refs.~\\cite{kkt,McD}, where as in our case the effect of the angular part of a flat direction has been considered. The difference from our case is that the flat direction is supposed to have zero VEV, so that the curvaton does not exist as a particle in the vacuum. As a consequence, the phase can only induce the curvature fluctuations in the radial field which leads to different predictions.    In addition, the proposal of \\cite{kkt} does not invoke the mass-squared of order $-H^2$, so that the radial potential is very flat with only a two-loop thermal correction breaking the symmetry. (In both cases, successful curvature perturbations can arise for $n=3$, contrary to our case with $n=1$.)  Because of such features the success of these models depends on computations which are much more tricky than in our case, though that is of course not necessarily an argument against them. However, we note that the computation of \\cite{kkt}, involving the very flat radial potential, works only if inflation lasts for a limited amount of time, because it takes the radial field during inflation to be at the edge of the slow-roll regime, $V'' \\sim H^2$. It seems more reasonable to assume that inflation lasts long enough to allow the quantum fluctuation to randomize the radial field  within  the smaller  region  $V\\lsim H^4$. (The analogous assumption for the angular  field is of course the one that we made.)  The other related model \\cite{pqcurv} also invokes a flaton model of PQ symmetry breaking, but now the curvaton candidate is a radial field and it is not clear how to keep it light in the early Universe. Also, the model works only if inflation is of short enough duration that the curvaton does not enter the randomization regime.  \\medskip  {\\bf Acknowledgments:} EJC was supported by the Korea Research Foundation Grant, KRF-2002-070-C00022.  \\bigskip  \\appendix"
    },
    "0402/astro-ph0402608_arXiv.txt": {
        "abstract": "{ It has recently been suggested by \\protect\\nocite{LumNat03}{Luminet} {et~al.} (2003) that the WMAP data are better matched by a geometry in which the topology is that of a Poincar\\'e dodecahedral model and the curvature is ``slightly'' spherical, rather than by an (effectively) infinite flat model. A general back-to-back matched circles analysis by \\protect\\nocite{CSSK03}{Cornish} {et~al.} (2004) for angular radii in the range $25-90\\ddeg$, using a correlation statistic for signal detection, failed to support this. In this paper, a matched circles analysis specifically designed to detect dodecahedral patterns of matched circles is performed over angular radii in the range $1-40\\ddeg$ on the one-year WMAP data. Signal detection is attempted via a correlation statistic and an rms difference statistic. Extreme value distributions of these statistics are calculated for one orientation of the 36$\\ddeg$ `screw motion' (Clifford translation) when matching circles, for the opposite screw motion, and for a zero (unphysical) rotation. The most correlated circles appear for circle radii of $\\alpha =11\\pm1 \\ddeg$, for the left-handed screw motion, but not for the right-handed one, nor for the zero rotation. The favoured six dodecahedral face centres in galactic coordinates are $(\\lII,\\bII)$ $\\approx (252\\ddeg,+65\\ddeg), (51\\ddeg,+51\\ddeg),$ $(144\\ddeg,+38\\ddeg), (207\\ddeg,+10\\ddeg),$ $(271\\ddeg,+3\\ddeg), (332\\ddeg,+25\\ddeg)$ and their opposites. The six pairs of circles {\\em independently} each favour a circle angular radius of $11\\pm1\\ddeg$. The temperature fluctuations along the matched circles are plotted and are clearly highly correlated. Whether or not these six circle pairs centred on dodecahedral faces match via a $36\\ddeg$ rotation only due to unexpected statistical properties of the WMAP ILC map, or whether they match due to global geometry, it is clear that the WMAP ILC map has some unusual statistical properties which mimic a potentially interesting cosmological signal. ",
        "introduction": "In the past twelve months, several authors have analysed the possibility that the primordial temperature fluctuations in the cosmic microwave background, as measured in the first-year data of WMAP (Wilkinson Microwave Anisotropy Probe) satellite \\nocite{WMAPSpergel}({Spergel} {et~al.} 2003, and accompanying papers), could be better matched by a perturbed Friedmann-Lema\\^{\\i}tre-Robertson-Walker (FLRW) model in which the global geometry is multiply connected rather than simply connected.  The possibility of the Universe being multiply connected is first known to have been suggested by \\nocite{Schw00,Schw98}{Schwarzschild} (1900, 1998). For recent reviews on cosmic topology, see \\nocite{LaLu95}{Lachi\\`eze-Rey} \\& {Luminet} (1995), \\nocite{Lum98}{Luminet} (1998), \\nocite{Stark98}{Starkman} (1998) and \\nocite{LR99}{Luminet} \\& {Roukema} (1999). For workshop proceedings on the subject, see \\nocite{Stark98}{Starkman} (1998) and following articles, and \\nocite{BR99}{Blanl{\\oe}il} \\& {Roukema} (2000). Detection strategies include both two-dimensional methods (based on temperature fluctuations in the surface of last scattering) and three-dimensional methods (based on distributions of gravitationally collapsed objects distributed in three-dimensional comoving space). For a list and discussion of both two-dimensional and three-dimensional methods, see Table~2 of \\nocite{LR99}{Luminet} \\& {Roukema} (1999) and the accompanying discussion.  The WMAP data has motivated many two-dimensional analyses.  Some of the authors mentioning either multiply connected models consistent with the WMAP data, or indirect hints of multiple connectedness, include \\nocite{WMAPSpergel,WMAPTegmarkFor,WMAPChiang,WMAPmultipol}{Spergel} {et~al.} (2003); {Tegmark}, {de Oliveira-Costa}, \\&  {Hamilton} (2003); {Chiang} {et~al.} (2003); {Copi}, {Huterer}, \\& {Starkman} (2003), --- finding low values of low $l$ multipoles or applying a multipole vector analysis.  However, other authors \\nocite{FengZhang03,Cline03,Contaldi03,EfstSph03}(e.g.~ {Feng} \\& {Zhang} 2003; {Cline}, {Crotty}, \\& {Lesgourgues} 2003; {Contaldi} {et~al.} 2003; {Efstathiou} 2003a) have suggested various non-topological explanations for the low multipole WMAP $C_l$ spectrum, such as double inflation or other phenomena from early universe physics, or positive curvature.  Yet others \\nocite{Naselsky03a,Naselsky03b}({Naselsky}, {Doroshkevich}, \\&  {Verkhodanov} 2003, 2004) point out that at least the ``internal linear combination'' (ILC) map of the WMAP data contains non-Poissonian signal, due to foreground residues on large scales. On smaller scales, \\nocite{Giommi03}{Giommi} \\& {Colafrancesco} (2003) find that somewhere between 20-100\\% of the signal at spherical harmonic $l$ values in the range $500 < l < 800$ may be due to blazars.  Attempts to {\\em exclude} classes of global geometry models using the WMAP data include the calculations of \\nocite{WMAPTegmarkAgainst}{de Oliveira-Costa} {et~al.} (2004) and of \\nocite{CSSK03}{Cornish} {et~al.} (2004). The latter performed a general back-to-back matched circles analysis for angular radii in the range $25-90\\ddeg$, using a correlation statistic for signal detection. They failed to find matched circles for a wide class of models, including the torus models, up to a scale of 16.8{\\hGpc}, which \\nocite{CSSK03}{Cornish} {et~al.} 2004 state as 24~Gpc, since they adopt $H_0= 70$\\kms Mpc$^{-1}$\\footnote{The Hubble constant is parametrised as $h\\equiv H_0/100$km~s$^{-1}$~Mpc$^{-1}.$}.  Disagreement also exists on whether or not the low quadrupole is really significant in rejecting the infinite flat `concordance' model \\nocite{EfstNoProb03a,EfstNoProb03b}({Efstathiou} 2003b, 2004).  Possibly one of the strongest claims in favour of a possible detection is that by \\nocite{LumNat03}{Luminet} {et~al.} (2003), who point out that given standard assumptions on the statistics of the fluctuations, the Poincar\\'e dodecahedral model implies a quadrupole and an octupole very close to those calculated from the WMAP data, for the same value of the total density parameter, $\\Omtot \\approx 1.013\\pm 0.02$.  In principle, this is excluded by the \\nocite{CSSK03}{Cornish} {et~al.} (2004) analysis.  The Poincar\\'e dodecahedral model requires positive (spherical) curvature. \\nocite{LumNat03}{Luminet} {et~al.} (2003) favour a total density parameter of $\\Omtot \\approx 1.013\\pm 0.002$ based on the spherical harmonic statistical analyses of the WMAP data, with non-relativistic matter density parameter, $\\Omm=0.28$ and cosmological constant $\\Omega_\\Lambda = \\Omtot - \\Omm$. This implies that points on the surface of last scattering which are multiple topological images of single physical points in space-time should correspond to matched circles \\nocite{Corn96,Corn98b}({Cornish}, {Spergel}, \\& {Starkman} 1996; {Cornish}, {Spergel}, \\&  {Starkman} 1998) which subtend angular radii at the observer of `about $35\\ddeg$' \\nocite{LumNat03}({Luminet} {et~al.} 2003).  Since the angular relations between face centres are identical for Euclidean and spherical dodecahedra, Euclidean calculations of the relative positions of circles are sufficient for testing the Poincar\\'e dodecahedral hypothesis. For example, the Euclidean half-angle of about $31.7\\ddeg$ is valid for angular separations of adjacent face centres in the spherical dodecahedron. Since the adjacent circles expected by \\nocite{LumNat03}{Luminet} {et~al.} (2003) have larger angular radii (subtended at the observer), this implies that they intersect with the face edges.  However, as is shown in \\SSS\\ref{s-notthirtyfive}, this angular radius is extremely sensitive to the value of $\\Omtot$. For example, keeping $\\Omm=0.28$ fixed, it is sufficient to decrease $\\Omtot$ to $\\Omtot=1.009$ to bring the angular radius to nearly zero.  This implies that a Poincar\\'e dodecahedral signal may have been missed by \\nocite{CSSK03}{Cornish} {et~al.} (2004) because they did not explore the part of parameter space for small matched circles.  In this paper, the missing part of parameter space is investigated.  A matched circles analysis specifically designed to detect dodecahedral patterns of matched circles is performed over angular radii in the range $1-40\\ddeg$ on Internal Linear Combination map (ILC) of the one-year WMAP data. The WMAP data are briefly discussed in \\SSS\\ref{s-wmap}. While the ILC is unlikely to be ideal for the studies of $C_l$ statistics, and some authors (cited above) claim correlations with foregrounds, it is hard to see how any signal mimicking matched circles oriented in a dodecahedral pattern could be imposed, either by the construction method of the ILC or by foregrounds.  Signal detection is attempted via a correlation statistic and an rms difference statistic, similarly to \\nocite{Rouk00a,Rouk00b}{Roukema} (2000b, 2000a). Extreme value distributions of these statistics between a right-handed rotation when matching circles, a left-handed rotation, and a zero rotation. A genuine signal should appear for either the right-handed or left-handed rotation, but not both, and should not appear for the zero rotation.  The relation between circle angular radius ($\\alpha$) and local cosmological parameters ($\\Omm, \\Omtot$) and the statistics used are presented in \\SSS\\ref{s-method}.  The tentative detection of dodecahedrally distributed matched circles with $\\alpha \\approx 10\\ddeg$ and analysis of their statistical significance are presented in \\SSS\\ref{s-results}. Further discussions and conclusions are made in \\SSS\\ref{s-conclu}.    For reviews on cosmological topology, see \\nocite{LaLu95}{Lachi\\`eze-Rey} \\& {Luminet} (1995), \\nocite{Lum98}{Luminet} (1998), \\nocite{Stark98}{Starkman} (1998) and \\nocite{LR99}{Luminet} \\& {Roukema} (1999). For workshop proceedings on the subject, see \\nocite{Stark98}{Starkman} (1998) and the following articles, and \\nocite{BR99}{Blanl{\\oe}il} \\& {Roukema} (2000). For a list and discussion of both two-dimensional and three-dimensional methods, see Table~2 of \\nocite{LR99}{Luminet} \\& {Roukema} (1999) and the accompanying discussion. The reader should be reminded that while microwave background data is still the most popular for topology analyses, considerable work in three-dimensional methods has been carried out, including, e.g., \\nocite{LLL96,Rouk96,FagG97,RE97,RB98,Gomero99a,LLU98,FagG99a,ULL99a,FagG99b,Gomero99b,Gomero99c}{Lehoucq}, {Lachi\\`eze-Rey}, \\&  {Luminet} (1996); {Roukema} (1996); {Fagundes} \\& {Gausmann} (1998b); {Roukema} \\& {Edge} (1997); {Roukema} \\& {Blanl{\\oe}il} (1998); {Gomero}, {Reboucas}, \\&  {Teixeira} (2002); {Lehoucq}, {Luminet}, \\& {Uzan} (1999); {Fagundes} \\& {Gausmann} (1998a); {Uzan}, {Lehoucq}, \\& {Luminet} (1999); {Fagundes} \\& {Gausmann} (1999); {Gomero}, {Reboucas}, \\&  {Teixeira} (2000, 2001).  For background on spherical multiply connected spaces, apart from the recent analysis by \\nocite{LumNat03}{Luminet} {et~al.} (2003), see \\nocite{GausSph01,LehSph02,RiazSph03}{Gausmann} {et~al.} (2001); {Lehoucq} {et~al.} (2002); {Riazuelo} {et~al.} (2003) for extremely thorough, in-depth mathematical background directly related to the cosmological context. \\postrefereechanges{ For a general background on geometry and topology, see e.g. \\nocite{Weeks2001}{Weeks} (2001). }  \\falpha  \\fomtot  \\fomm  Comoving coordinates are used when discussing distances (i.e. `proper distances', \\nocite{Wein72}{Weinberg} 1972, equivalent to `conformal time' if $c=1$). ",
        "conclusions": "\\label{s-conclu}  It is clearly premature to claim a highly significant detection of the topology of the Universe based on just one simple analysis of the first year WMAP ILC cosmic microwave background map.  However, the plots are striking and it seems prudent to release them to the scientific community while a companion paper is prepared with formal statistical analyses.  Whether or not the matched circles found are just coincidence or due to global geometry, it is clear that temperature fluctuations around 12 dodecahedrally spaced circles of radius $11\\pm1\\ddeg$ in the WMAP ILC map correlate unusually well in their respective pairs when a phase shift of $36\\ddeg$, corresponding to a left-handed screw motion, is applied.  \\fcmodem \\fcmodep \\fcmodez  If simulations can show that this is a fairly likely occurrence due to Gaussian fluctuations in an infinite flat universe, then this would show that genuine matched circles will be even harder to distinguish from spurious detections than was previously thought.  However, \\nocite{Vielva03}{Vielva} {et~al.} (2003) have found non-Gaussian fluctuations at about $10\\ddeg$ - extremely close to radius of the matched circle radius. While it is not obvious how this non-Gaussianity detection should relate to the radius of matched circles, since correlations are between circles, not along individual circles, this would also complicate simulations, since they would need to be consistent with observational analyses like this one.  \\nocite{CSSK03}{Cornish} {et~al.} (2004) avoided circles of radii smaller than $25\\ddeg$ because of the risk of false positives, and carried out extensive simulations of what signal would be expected from a genuinely multiply connected universe.  However, these simulations risk the problem of cosmic variance and logical circularity. If the Universe really is detectably multiply connected, then it would be fairly reasonable that the density perturbations (eigenmodes) on the largest scales are somewhat affected by the global physics of the Universe, so that modelling them on the basis of Gaussianity and random phases may lead to statistical statements which are incorrect, because they talk about an ensemble of likely universes with different statistical properties to the real Universe.  \\fcircleA \\fcircleB \\fcircleC \\fcircleD \\fcircleE \\fcircleF  The cosmic variance problem is that our actual Universe is just one realisation --- the number of perturbations on large scales are too small for the large number theorem to make statistics of ensembles valid.  An apparently strong heuristic argument against these matches being physical is the need for extreme fine-tuning. If the claimed matched circles are due to topology, then the in-radius of the Universe happens by chance to be about $\\cos(11\\ddeg) \\approx 98\\%$ of the distance to the surface of last scattering. Intuitively, this is difficult to accept.  However, as pointed out by \\nocite{LumNat03}{Luminet} {et~al.} (2003), in the case of a positively curved Universe, especially if $\\Omtot \\approx 1.01$, there is necessarily a fine-tuning, since the curvature scale is only about three times the matter-dominated horizon size.  Moreover, the non-zero cosmological constant, $\\Omega_\\Lambda \\approx 0.7$, is well established observationally (though see \\nocite{BlanchSarkar03,VaucBlanch03}{Blanchard} {et~al.} (2003); {Vauclair} {et~al.} (2003) for a minority viewpoint) and definitely requires fine-tuning of some sort.  Attempts have been made to link a non-zero cosmological constant with detectable cosmic topology, but so far no obvious successes have been found.  Independently of theoretical arguments, the results presented in this paper are testable by several observationally based methods which do not require assumptions on hypothetical statistical ensembles of universes, and each should potentially be able to improve the signal if it is cosmological in origin: \\begin{list}{(\\roman{enumi})}{\\usecounter{enumi}} \\item an attempt to separate out the na\\\"{\\i}ve Sachs-Wolfe effect, the doppler component, and the integrate Sachs-Wolfe effect, and analysis at higher resolution \\item improved removal of foregrounds, either independently of the hypothesis, or using the matched circles hypothesis to predict foregrounds which intervene in one circle but not the other (cf. \\SSS\\ref{s-pointsources}) \\item polarisation data from Planck. \\end{list}  Similarly to the prediction by \\nocite{LumNat03}{Luminet} {et~al.} (2003) that `$\\Omtot \\approx 1.013 > 1$', Fig.~\\ref{f-alpha} shows that for $\\Omm = 0.28\\pm0.02,$ the total density parameter must be $\\Omtot \\approx 1.010 \\pm 0.001$ for these matched circles of radius $11\\pm1\\ddeg$ to be cosmological in origin. A larger uncertainty in $\\Omm$ would correspondingly increase the uncertainty in $\\Omtot$, but the prediction that $\\Omtot > 1$ remains: the fundamental dodecahedron of the Poincar\\'e dodecahedral manifold is that of a manifold of positive curvature."
    },
    "0402/astro-ph0402322_arXiv.txt": {
        "abstract": "Fluctuations in the redshifted 21 centimeter emission from neutral hydrogen probe the epoch of reionization. We examine the observability of this signal and the impact of extragalactic foreground radio sources (both extended and point-like). We use cosmological simulations to predict the angular correlation functions of intensity fluctuations due to unresolved radio galaxies, cluster radio halos and relics and free-free emission from the interstellar and intergalactic medium at the frequencies and angular scales relevant for the proposed 21cm tomography. In accord with previous findings, the brightness temperature fluctuations due to foreground sources are much larger than those from the primary 21cm signal at all scales. In particular, diffuse cluster radio emission, which has been previously neglected, provides the most significant foreground contamination. However, we show that the contribution to the angular fluctuations at scales $\\theta \\approxgt 1$ arcmin is dominated by the spatial clustering of bright foreground sources. This excess can be removed if sources above flux levels $S \\approxgt 0.1$ mJy (out to redshifts of $z \\sim 1$ and $z \\sim 2$ for diffuse and point sources respectively) are detected and removed. Hence, efficient source removal may be sufficient to allow the detection of angular fluctuations in the 21cm emission free of extragalactic foregrounds at $\\theta \\approxgt 1$ arcmin. In addition, the removal of sources above $S=0.1$ mJy also reduces the foreground fluctuations to roughly the same level as the 21cm signal at scales $\\theta \\approxlt 1$ arcmin. This should allow the substraction of the foreground components in frequency space, making it possible to observe in detail the topology and history of reionization. ",
        "introduction": "In the past few years a quantitative study of the high-redshift intergalactic medium (IGM) and its reionization history has finally been made possible by the discovery of quasars at $z>5.8$ (e.g. Fan et al. 2001, 2003). In particular, the detection of a Gunn-Peterson trough (Gunn \\& Peterson 1965) in the Keck (Becker et al. 2001) and VLT (Pentericci et al. 2002) spectra of the Sloan Digital Sky Survey quasar SDSS 1030-0524 at $z=6.28$ and in the Keck spectrum of SDSS 1148+5251 at $z=6.37$ (White et al. 2003), has been interpreted as the signature of the trailing edge of the cosmic reionization epoch. The recent analysis of the first year of data from the Wilkinson Microwave Anisotropy Probe (WMAP) satellite on the temperature and polarization anisotropies of the cosmic microwave background (CMB), infers a mean optical depth to Thomson scattering $\\tau_e \\sim 0.17$, suggesting that the universe was reionized at higher redshift (Kogut et al. 2003; Spergel et al. 2003). Physically, the CMB and the Gunn-Peterson trough probe two different stages of reionization, the former being sensitive to the initial phase, when free electrons appear, the latter to the residual neutral hydrogen in the latest stages of reionization. None of the two methods though is able to constrain the exact ionization level or the details of the reionization history. For this reason, an alternative way to probe the high-redshift IGM is required.  An optimal experiment for probing the various stages of reionization is the proposed 21cm tomography. It has indeed been shown that neutral hydrogen in the intergalactic medium (IGM) and gravitationally collapsed systems should be directly detectable in emission or absorption against the cosmic microwave background radiation (CMB) at frequencies corresponding to the 21cm line (e.g.;Field 1958, 1959; Scott \\& Rees 1990; Kumar et al.1995). In principle it will possible to carry out such an experiment with planned high sensitivity radio telescopes such as the PrimevAl Structure Telescope (PAST) \\footnote{\\tt http://astrophysics.phys.cmu.edu}, the Square Kilometer Array (SKA) \\footnote{\\tt http://www.nfra.nl/skai} and the LOw Frequency ARray (LOFAR) \\footnote{\\tt http://www.astron.nl/lofar}. The 21cm spectral features will display redshift dependent angular structure due to evolving inhomogeneities in the gas density field, hydrogen ionized fraction, and spin temperature. Several different signatures have been investigated in the recent literature: the fluctuations in the 21cm line emission induced both by the inhomogeneities in the gas density and in the ionized hydrogen fraction (Madau, Meiksin \\& Rees 1997; Tozzi et al. 2000; Ciardi \\& Madau 2003, hereafter CM; Furlanetto, Sokasian \\& Hernquist 2004) and by `minihalos' with virial temperatures below $10^4\\,$K (Iliev et al. 2002, 2003); the global feature (`reionization step') in the continuum spectrum of the radio sky that may mark the abrupt overlapping phase of individual intergalactic HII regions (Shaver et al. 1999); and the 21cm narrow lines generated in absorption against very high redshift radio sources by the neutral IGM (Carilli, Gnedin \\& Owen 2002) and by intervening minihalos and protogalactic disks (Furlanetto \\& Loeb 2002).  While the 21cm tomography proposes to map the topology of the reionization process and constrain the nature of the ionizing sources, it remains a challenging project due to foreground contamination from unresolved extragalactic radio sources (Di Matteo et al. 2002), free-free emission from the same halos that reionize the universe (Oh \\& Mack 2003) and the Galactic free-free and synchrotron emission (Shaver et al. 1999). Because the proposed experiments will be carried out in frequency space as well as angle, recent work has discussed the possibility of removing the foreground power spectrum components by comparing maps closely spaced in frequency (Zaldarriaga, Furlanetto \\& Hernquist 2004; Gnedin \\& Shaver 2004). The proposed substraction is nonetheless demanding as the foreground fluctuations overwhelm the primary 21cm signal by up to three orders of magnitude. In particular, it will require knowing in detail the spectral behavior and possible spectral variations of each of the contaminants and their associated angular power spectra.  In this paper we employ computer simulations of a $\\Lambda$CDM universe to evaluate the foreground brightness temperature fluctuations due to extragalactic sources that contaminate the 21cm tomography. In particular, we model the free-free emission from ionizing sources self-consistently with a viable reionization scenario and with the associated 21cm IGM emission (as described in CM). Using the same simulations we adopt simple but physically motivated prescriptions to model the radio galaxy population; the model successfully reproduces both the observed luminosity function (at 1.4GHz) and the two point correlation function of the observed radio sources. Finally, we use a separate simulation (of the same $\\Lambda$CDM universe) that, in addition to dark matter and baryonic gas, follows the evolution of magnetic fields and cosmic rays, to estimate the foreground signals due to extended radio sources such as cluster radio relics and radio halos, which had not been considered in previous work.  We shall show that the detailed information on the spatial and redshift distribution of the contaminant sources is crucial for determining the correct brightness temperature fluctuations due to the unresolved extragalactic foregrounds on the angular scales where the primary 21cm signal is expected to peak. By modeling the foregrounds within the cosmological simulations we are able to study and map the variation of the angular clustering signal as a function of the flux cut above which foreground sources will be detected (and hence will not contribute to the unresolved fluctuating signal). In particular, we will show that the angular power spectra of unresolved foreground contaminants become significantly suppressed at scales $\\theta \\approxgt 1$ arcmin when sources above flux levels of a fraction of mJy are removed.  The structure of the paper is as follows. In Section~2 we describe the numerical simulations adopted to model the physical processes producing the extragalactic backgrounds. In Section~3 we briefly outline the origin of the 21cm emission line from the diffuse IGM used in CM, while in Section~4 we show the total comoving emissivities and spatial correlation functions for the extragalactic foregrounds modeled within the simulations. In Sections~5, and~6 we calculate the contribution to the brightness temperature fluctuations to to the clustering of the different foregrounds and finally, in Section~7 we give our conclusions. ",
        "conclusions": "We have examined the contribution of extragalactic foregrounds to the angular brightness temperature fluctuations at the frequencies relevant for observations of the redshifted 21cm signal. We have used a set of cosmological simulations to model the evolution of radio galaxies, free-free emission from the ISM and IGM and radio halos and relics. We calculated their comoving space density emissivity and spatial correlation functions and projected the simulation volumes to predict their expected sky temperature fluctuations.  Because high resolution observations of the 21cm signal will be carried out in both frequency and angle it is best to summarize the main conclusion from our analysis by showing the intrinsic and foreground signals as a function of both scale and frequency. Fig.~\\ref{cont} shows the angular power spectrum of the 21cm signal (from CM simulations), and that of the radio galaxies and radio halos (the two dominating foreground components for which source removal was studied here) as a function of frequency (where we have taken the spectral index of radio galaxies and radio halos to be $\\alpha = 0.8$ and 1 respectively) and angle (scale $l$). In summary, \\begin{itemize} \\item As previously emphasized by Di Matteo et al. (2002, for the case of radio galaxies) and Oh \\& Mack (2002, for the case of free-free emission from the ISM) we have found that the absolute foreground signal fluctuates more strongly than that of the 21 cm signal at all angular scales and frequencies (left panel Fig.~\\ref{cont}). In addition, here we have shown that the largest fraction of these fluctuations is due to extended cluster radio halos and relics, both of which had previously been neglected in the context of the 21cm tomography experiment (see however Waxman \\& Loeb 2000). Radio emission from the cores of galaxies contributes more strongly than the free-free emission from their ISM. The presence of foreground signal many orders of magnitude above the primary signal needs to be addressed because it may seriously hinder the detection of the 21 cm line emission as the latter could be easily mimicked by small errors/irregularities in the substraction of the foreground spectra.  \\item We have considered source removal above a flux limit $S_c=0.1$~mJy, (close to LOFAR sensitivity) for point sources and radio halos (as an example of an extended source). We have found that, as expected, the amplitude of the foreground power spectra of both radio halo and radio galaxies decreases at all scales. In particular, with the additional spatial information on the foreground sources provided by the simulations, we have shown that the amplitude of the angular foreground fluctuations drops well below that of the 21cm at scales $\\theta \\approxgt 1$ arcmin and becomes comparable to it at scales smaller that $\\theta \\sim 1$ arcmin (right panel Fig.~\\ref{cont}). At angular scales larger than $\\theta \\sim 1$ arcmin and at frequencies $\\nu \\approxlt 150$ MHz (simply corresponding to the redshifts prior to full reionization in the S5 model) the 21cm signal should be detected directly. At scales smaller than $\\theta \\sim 1$ arcmin, substraction of the foregrounds can be feasibly carried out in frequency space, as the overall foreground signal is at most of the same order as the primary 21cm. Comparing maps closely spaced in frequency, (as discussed in detail by Zaldarriaga et al. 2004) after the removal of bright sources should make the detection of the signal feasible at small scales also (where the foreground power spectra are different than the one of the 21 cm; Fig.~\\ref{cont})\\footnote{We note further that the amplitude of the 21cm signal shown in Fig.~\\ref{cont} corresponds to a given bandwidth $\\Delta\\nu=\\nu\\Delta z/(1+z)=1$MHz and is expected to increase with bandwidth resolution. CM show that the signal can increase by a factor up to 4 at small angles and high frequencies if e.g.; $\\Delta\\nu=0.1$MHz can be achieved making foreground substraction less crucial at scales $\\theta \\approxlt 1'$}. \\end{itemize}  From Fig.~\\ref{cont}, we conclude that bright sources contribute to the angular power spectrum at $\\theta \\approxgt 1$ arcmin and, once removed, the power on these scales is suppressed. This effect is due to a scale dependent bias of the spatial correlation functions of sources below a certain flux cut. This result emphasizes the importance, when determining the angular fluctuation and power spectra of foregrounds, of using simulations to project the detailed spatial distribution of foregrounds as a function of flux cut. Analytical estimates of brightness temperature fluctuations (Di Matteo et al. 2001; Oh \\& Mack 2002), which had to assume a certain angular clustering of the radio foregrounds, appear to be insufficient for determining the effect of these contaminants on all scales. This had lead previous authors to generally more pessimistic conclusions for the feasibility of 21cm tomography. We find instead that, due to the effects of bright source removal, 21cm observations carried out at angular scales $\\theta$ of the order of a few arcmin should not be strongly affected by extragalactic foregrounds. However, how source substraction can be carried out from the observed maps, needs detailed work. The maps we have simulated will need to be folded with the instrument response functions as a function of angle and frequency (and other specific details of the instrument design). This is currently being investigated for the case of LOFAR by Valdes et al. (2004).  We note that determining the angular scale at which the foreground signal should drop out is, of course, subject to some uncertainty associated to specifics of source removal. Besides needing to account for the instrument response, which is beyond the scope of this paper, this will of course be a function of the theoretical approach we have adopted to model the foreground source populations down to flux limits currently unobserved. In fact, it is part of the scientific goals of future high sensitivity low frequency radio facilities such as LOFAR, PAST and SKA to sample for the first time faint radio galaxies and diffuse cluster radio emission (radio halos and radio relics, see En\\ss lin \\& R\\\"ottgering 2002, Br\\\"uggen et al. 2004) and hence determine the properties of these population up to high redshifts."
    },
    "0402/astro-ph0402578_arXiv.txt": {
        "abstract": "The Absolute Radiometer for Cosmology, Astrophysics, and Diffuse Emission (ARCADE) is a balloon-borne instrument to measure the temperature of the cosmic microwave background at centimeter wavelengths. ARCADE uses narrow-band cryogenic radiometers to compare the sky to an external full-aperture calibrator. To minimize potential sources of systematic error, ARCADE uses a novel open-aperture design which maintains the antennas and calibrator at temperatures near 3 K at the mouth of an open bucket Dewar, without windows or other warm objects between the antennas and the sky. We discuss the design and performance of the ARCADE instrument from its 2001 and 2003 flights. ",
        "introduction": "The cosmic microwave background is a thermal relic from a hot, dense phase in the early universe. Deviations from a perfect blackbody spectrum carry information on the energetics of the early universe. Measurements across the peak of the spectrum limit deviations from a blackbody to less than 50 parts per million \\citep{fixsen/etal:1996,gush/etal:1990}. Direct observational limits at longer wavelengths, though, are weak: distortions as large as 5\\% could exist at wavelengths of several centimeters or longer without violating existing observations.  Plausible physical processes can generate observable distortions at long wavelengths without violating limits established at shorter wavelengths. The decay of massive particles produced near the Big Bang imparts a chemical potential to the CMB, creating a deficit of photons at long wavelengths \\citep{sunyaev/zeldovich:1970, silk/stebbins:1983, burigana/etal:1995, mcdonald/etal:2001, hansen/haiman:2004}. Reionization of the universe by the first collapsed structures distorts the spectrum through thermal bremsstrahlung by the ionized gas, characterized by a quadratic rise in temperature at long wavelengths \\citep{bartlett/stebbins:1991}. Reionization is expected to produce a cosmological free-free background with amplitude of a few mK at frequency 3 GHz \\citep{haiman/loeb:1997, oh:1999}. Such a signal is well below current observational limits, which only constrain free-free distortions to $\\Delta T < 19 $ mK at 3 GHz \\citep{bersanelli/etal:1994}.  Detecting the signal from reionization requires accuracy of order 1 mK at frequencies below 30 GHz. Coherent receivers using High Electron Mobility Transistor (HEMT) amplifiers can easily achieve this sensitivity, leaving systematic error as the limiting uncertainty in previous CMB measurements below 30 GHz (see \\cite{kogut:1992} for an experimental review). These fall into 3 categories. Below 2 GHz the dominant uncertainty is synchrotron emission within the Galaxy. Ground-based measurements at higher frequencies are limited by atmospheric emission, while balloon-borne experiments have been limited by emission from warm parts of the instrument. The Absolute Radiometer for Cosmology, Astrophysics, and Diffuse Emission (ARCADE) is a fully cryogenic, balloon-borne instrument designed to avoid these sources of systematic error and provide new limits on deviations from a blackbody spectrum at centimeter wavelengths comparable to the limits established at millimeter wavelengths.  ARCADE represents a long-term effort to characterize the CMB spectrum at cm wavelengths in order to constrain the thermal history of the early universe. An engineering flight designed to test the cold open optics launched from Ft Sumner, NM on November 2 2001 UT. A second flight, optimized for CMB observations, launched from Palestine, TX on June 15 2003 UT. Scientific analysis of the 2003 flight is presented by \\citet{fixsen/etal:2004}. We describe the design of the ARCADE instrument and discuss its performance during the 2001 and 2003 flights.  \\begin{figure} \\plotone{arc_inst_schematic.eps} \\caption{ARCADE instrument schematic. Cryogenic radiometers compare the sky to an external blackbody calibrator. The antennas and external calibrator are maintained near 2.7 K at the mouth of an open bucket Dewar; there are no windows or other warm objects between the antenna and the sky. \\label{arcade_schematic} } \\end{figure} ",
        "conclusions": "Both the 2001 and 2003 flights demonstrate the viability of large open-aperture cryogenic optics for CMB measurements. The ARCADE cryogenic design maintains the external calibrator, antennas, and cryogenic radiometers at temperatures near 2.7 K for many hours at 35 km altitude. Cold boiloff gas vented through the aperture plane reduces condensation of atmospheric nitrogen to levels consistent with the desired CMB observations; the primary impact is the cooling required to handle the additional heat load on the aperture. The current instrument is relatively small (two frequency channels at 10 and 30 GHz) and is intended primarily as a pathfinder to verify the cryogenic design under flight conditions. A larger second-generation instrument with 6 channels exending down to 3 GHz is currently under construction and is scheduled to launch in 2005."
    },
    "0402/astro-ph0402052_arXiv.txt": {
        "abstract": "The central region of the Galaxy has been observed at 580, 620 and 1010 MHz with the Giant Metrewave Radio Telescope (GMRT). We detect emission from Sgr~A*, the compact object at the dynamical centre of the Galaxy, and estimate its flux density at 620 MHz to be 0.5$\\pm$0.1 Jy. This is the first detection of Sgr~A* below 1 GHz \\citep{IAU199.ROY,GC2003.ROY}, which along with a possible detection at 330 MHz \\citep{NORD2004} provides its spectrum below 1 GHz.  Comparison of the 620 MHz map with maps made at other frequencies indicates that most parts of the Sgr~A West HII region have optical depth~$\\sim$2.  However, Sgr~A*, which is seen in the same region in projection, shows a slightly inverted spectral index between 1010 MHz and 620 MHz.  This is consistent with its high frequency spectral index, and indicates that Sgr~A* is located in front of the Sgr~A West complex, and rules out any low frequency turnover around 1 GHz, as suggested by \\citet{DAVIES1976}. ",
        "introduction": "Being located two orders of magnitude closer than the nearest large galaxy, the Galactic Centre (GC) can be studied at a much higher spatial resolution and sensitivity than is possible for other galaxies. Because of this advantage, we can identify unique objects like the Radio-arc consisting of linear parallel filaments \\citep{YUSEF-ZADEH1984}, or the 2.6$ \\times 10^6$ M$_{\\odot}$ black hole suggested to be associated with the compact radio source Sgr~A* \\citep{GHEZ1998}. From the high resolution ($\\sim$~arc~seconds) observation by the Very Large Array (VLA) at radio frequencies \\citep{EKERS1983,PEDLAR1989}, the following sources within the central 15$'$ of the Galaxy have been identified. \\\\ (i) At the dynamical centre of the Galaxy is the compact nonthermal radio source known as Sgr-A*. (ii) Around Sgr~A* is the HII region Sgr-A West \\citep{EKERS1983}, whose morphology resembles a face-on spiral galaxy. (iii) Near Sgr-A West is Sgr-A East, which is believed to be a supernova remnant (SNR). (iv) A 7$'$ halo, which has been proposed to be a mixture of thermal and non-thermal emission \\citep{PEDLAR1989}.  Sgr~A* (see Melia \\& Falcke 2001 for a review) was not detected at frequencies below 960 MHz and observations at 408 MHz \\citep{DAVIES1976} and at 330 MHz \\citep{PEDLAR1989} provide upper limits ($\\le$0.1 Jy at 330 MHz) on its flux density. Sgr~A* probably has a low frequency turnover below 1 GHz, but the nature of the turnover has never been examined in detail \\citep{MELIA2001}.  Recently, \\citet{NORD2004} claim to have detected Sgr~A* at 330 MHz.  However, the average brightness of the 7$'$ halo seen towards Sgr~A* at 330 MHz is $\\sim 100$ mJy/Beam (with the beamsize used in their map), which is comparable to the claimed peak intensity of Sgr~A* (Fig.~2 in Nord et al. 2004). The 7$'$ halo could be located in front of the Sgr~A complex \\citep{PEDLAR1989}, and presence of any small scale structure in the halo along Sgr~A* can mimic its claimed detection. Therefore, detection of Sgr~A* at 330 MHz remains provisional.  High resolution radio observations \\citep{ROBERTS1993} show that Sgr~A~West comprises of three major features of ionised gas known as Northern and Southern arm and Western arc, which are embedded in a halo of lower density ionised gas with an extent of about 1.5$'$ \\citep{MEZGER1986, PEDLAR1989}. Along the Northern arm, the gas appears to flow away from us. If this is taken as an indication of gas falling in towards Sgr~A*, then this would imply that the Northern arm is located in front of Sgr~A*. Though \\citet{WHITEOAK1983} have suggested Sgr~A* to be located in front of Sgr A West, \\citet{LISZT1983} detected HI absorption against Sgr~A* at 40$-$60 \\kms, and not against Sgr~A~West. While \\citet{LISZT1983} attributes this discrepancy to the patchiness of the HI screen, the relative location of Sgr~A* with respect to Sgr~A~West need to be established.  Due to free-free absorption at low radio frequencies, HII regions tend to get optically thick and absorption against another continuum object can be used to constrain their relative location.  We have observed the central half a degree region of the Galaxy at 580, 620 and 1010 MHz using the GMRT \\citep{SWARUP1991}. As a cross-check, we have observed Sgr~A* also at 580 MHz. To estimate its spectral index between 1 GHz and 620 MHz and compare with its spectrum obtained at higher radio frequencies, we have further observed this region at 1010 MHz with the GMRT. In this paper, we will mainly discuss Sgr~A* and Sgr~A~West HII region. In Sect.~2, we describe the observations and data analysis.  The results are presented in Sect.~3, and the inferences are discussed in Sect.~4. The conclusions are presented in Sect.~5. ",
        "conclusions": "Observations of the GC region at 1010, 620 and 580 MHz with the GMRT and a comparison with the existing observations made at other radio frequencies have provided us several important details about the region:\\\\ (i) Sgr~A* has been detected at 580 MHz, which is the lowest frequency unambiguous detection of Sgr~A*, and the estimated flux density is consistent with what is expected from its higher radio frequency spectral index and the flux density.  This indicates that there is no low frequency turnover of its emission at freqencies above 0.6 GHz. \\\\ (ii) Our observations at 0.6 GHz breaks the degeneracy between emission measure and electron temperature in the existing data, and allows us to estimate the optical depth of the Sgr~A~West HII region to be about 2.5 at 620 MHz. \\\\ (iii) Though Sgr~A* is located along the same line of sight as Sgr A West, the emission from it undergoes no absorption by this HII region, which indicates that Sgr~A* is located in front of Sgr~A~West."
    },
    "0402/astro-ph0402502_arXiv.txt": {
        "abstract": "Viscosity and magnetic fields drive differentially rotating stars toward uniform rotation, and this process has important consequences in many astrophysical contexts.  For example, merging binary neutron stars can form a ``hypermassive'' remnant, i.e.\\ a differentially rotating star with a mass greater than would be possible for a uniformly rotating star.  The removal of the centrifugal support provided by differential rotation can lead to delayed collapse of the remnant to a black hole, accompanied by a delayed burst of gravitational radiation.  Both magnetic fields and viscosity alter the structure of differentially rotating stars on secular timescales, and tracking this evolution presents a strenuous challenge to numerical hydrodynamic codes.  Here, we present the first evolutions of rapidly rotating stars with shear viscosity in full general relativity.  We self-consistently include viscosity in our relativistic hydrodynamic code by solving the fully relativistic Navier-Stokes equations. We perform these calculations both in axisymmetry and in full 3+1 dimensions. In axisymmetry, the resulting reduction in computational costs allows us to follow secular evolution with high resolution over dozens of rotation periods (thousands of $M$). We find that viscosity operating in a hypermassive star generically leads to the formation of a compact, uniformly rotating core surrounded by a low-density disk.  These uniformly rotating cores are often unstable to gravitational collapse.  We follow the collapse in such cases and determine the mass and the spin of the final black hole and ambient disk. However, viscous braking of differential rotation in hypermassive neutron stars does not always lead to catastrophic collapse, especially when viscous heating is substantial.  The stabilizing influences of viscous heating, which generates enhanced thermal pressure, and centrifugal support prevent collapse in some cases, at least until the star cools.  In all cases studied, the rest mass of the resulting disk is found to be 10-20\\% of the original star, whether surrounding a uniformly rotating core or a rotating black hole. This study represents an important step toward understanding secular effects in relativistic stars and foreshadows more detailed, future simulations, including those involving magnetic fields. ",
        "introduction": "\\label{intro}  The field of numerical relativity has matured to a stage where it is possible to simulate realistic systems of astrophysical interest. In this paper, we examine the global effects of viscosity on differentially rotating, relativistic stars. Viscosity can have significant effects on the stability of neutron stars. For example, it can drive a secular bar instability in rapidly rotating neutron stars, as shown in Newtonian gravitation~\\cite{c69,sl96} and in general relativity~\\cite{sz98}. Viscosity can suppress the $r$-modes~\\cite{lom98,jjl0102} and other gravitational-radiation driven instabilities, including the secular bar modes~\\cite{ll7783}. Viscosity also destroys differential rotation, and this can cause significant changes in the structure and evolution of differentially rotating massive neutron stars.  Differentially rotating neutron stars can support significantly more rest mass than their nonrotating or uniformly rotating counterparts, making ``hypermassive'' neutron stars possible~\\cite{bss00,lbs03}. Such hypermassive neutron stars can form from the coalescence of neutron star binaries~\\cite{rs99,Shibata:1999wm,stu03} or from rotating core collapse. The stabilization arising from differential rotation, although expected to last for many dynamical timescales, will ultimately be destroyed by magnetic braking and/or viscosity~\\cite{bss00,s00}. These processes drive the star to uniform rotation, which cannot support the full mass of the hypermassive remnant. This process can lead to ``delayed'' catastrophic collapse to a black hole, possibly accompanied by some mass loss. Such a delayed collapse might emit a delayed gravitational wave signal detectable by laser interferometers. Moreover, the collapse, together with any residual gas in an ambient accretion disk, could be the origin of a gamma-ray burst (GRB).  Both magnetic fields and viscosity can destroy differential rotation in a rapidly rotating star~\\cite{s00,css03,ls03}. Simple estimates show that the magnetic braking (Alfv\\'en) timescale for a laminar field is much shorter than the timescale of molecular (neutron) viscosity in a typical massive neutron star. Hence magnetic fields are expected to be the principal mechanism driving neutron stars toward rigid rotation. Phase mixing arising from magnetic braking~\\cite{spruit99,ls03}, or other possible magnetohydrodynamic instabilities~\\cite{spruit99,balbus98} might stir up turbulence. Turbulent shear viscosity could then dominate the subsequent evolution.  In this paper, we are primarily interested in identifying the global evolutionary consequences of shear viscosity in a relativistic star, independent of the detailed nature or origin of the viscosity.  To explore the consequences of the loss of differential rotation in equilibrium stars, we study the secular evolution of differentially rotating relativistic stars in the presence of a shear viscosity. Viscosity and magnetic fields have two things in common: (1) they both change the angular velocity profiles of a differentially rotating star, and (2) they both act on {\\em secular} timescales, which can be many rotation periods. The latter inequality poses a severe challenge to numerical simulations using a hydrodynamic code. It is too taxing for a {\\it hydrodynamic} code using an explicit differencing scheme to evolve a star for physical realistic secular timescales. To solve this problem, we artificially amplify the strength of viscosity so that the viscous timescale is short enough for numerical treatment. However, we keep the viscous timescale substantially longer than the dynamical timescale of the stars, so that the evolution of the star remains quasi-stationary. We then check the validity of our results by reducing the viscosity on successive runs and testing that the viscosity-induced physical behavior is unchanged; rather, only the timescale changes and does so inversely with the strength of viscosity. A more detailed discussion of the expected scaling is presented in Section~\\ref{justification}~\\cite{footnote0}.  To study viscous evolution, we need to perform long simulations in full general relativity. Typically, we evolve the stars in axisymmetry. This allows us to follow the secular evolution of the stars with high resolution in a reasonable amount of time. Viscosity can, however, drive nonaxisymmetric instabilities when a star is rapidly rotating. To test for such instabilities, we also perform lower-resolution, three-dimensional (3D) simulations on the most rapidly rotating stars we consider.  For non-hypermassive neutron stars that are slowly and differentially rotating, we find that viscosity simply drives the whole star to rigid rotation. If the non-hypermassive neutron star is rapidly and differentially rotating, however, viscosity drives the inner core to rigid rotation and, at the same time, expels the material in the outer layers. The final system in this case consists of a rigidly-rotating core surrounded by a low-density, ambient disk in quasi-stationary equilibrium.  Our most interesting results concern the fate of hypermassive neutron stars. We numerically evolve four models with different masses and angular momenta. We find that in all cases, viscosity drives the cores to rigid rotation and transports angular momentum outwards into the envelope. As a result, the core contracts in a quasi-stationary manner, and the outer layers expand to form a differentially rotating torus.  Of the four models we have studied, the star with the highest mass collapses to a black hole, with about 20\\% of the rest mass leftover to form a massive accretion disk. On the contrary, the other three stars do not collapse to black holes, but form star + disk systems, similar to the final state of the rapidly rotating non-hypermassive neutron stars described above. As will be discussed in Section~\\ref{rad_cooling}, viscosity generates heat so that the stars do not evolve adiabatically in general. The extra thermal pressure due to viscous heating helps to support the stars. We also consider the limit of rapid cooling, whereby the heat generated by viscosity is immediately removed from the stars. Of the three stars which do not collapse to black holes in the no-cooling limit, we found that the one with the lowest angular momentum undergoes catastrophic collapse in the rapid-cooling limit. About 10\\% of the rest mass is leftover to form an accretion disk in this case. To test the validity of the axisymmetric results, we perform 3D simulations to check for any nonaxisymmetric instabilities. We do not find any unstable nonaxisymmetric modes and the 3D results agree with the axisymmetric results.  Our results suggest that viscous braking of differential rotation in a hypermassive neutron star can, but does not always, lead to catastrophic collapse. When catastrophic collapse does occur, the remnant is a black hole surrounded by a massive accretion disk. This outcome is very different from that of the collapse of an unstable, rigidly-rotating ``supramassive'' neutron star, in which the whole star collapses to a black hole, leaving only a tiny amount of material to form a disk~\\cite{sbs00,s03}. Many models for GRBs require a massive disk around a rotating black hole to supply energy by neutrino processes~\\cite{rosswog}. Our results suggest that viscous forces in a hypermassive star could lead to the formation of a massive disk around such a black hole.  The structure of this paper is as follows. In Section~\\ref{setup}, we derive the relativistic Navier-Stokes equations containing shear viscosity in a 3+1 form suitable for numerical integration, and describe how we evolve them in both axisymmetry and full 3+1 dimensions. We then describe in Section~III several tests that we perform to check our code.  We present the results of our simulations on five selected stars in Section~IV. Finally, we briefly summarize and discuss our conclusions in Section~V. ",
        "conclusions": "We have simulated the evolution of rapidly rotating stars in full general relativity including, for the first time, shear viscosity. Our findings indicate that the braking of differential rotation in hypermassive stars always leads to significant structural changes, and often to delayed gravitational collapse.  The rest mass, angular momentum, and thermal energy all play a role in determining the final state. We performed axisymmetric numerical simulations of five models to study the influence of these parameters. In the presence of shear viscosity, the most hypermassive model which we studied (star~I), collapses to a black hole whether we evolve by ignoring cooling, or by assuming rapid cooling of the thermal energy generated by viscosity.  However, the viscous transport of angular momentum to the outer layers of the star results in mass outflow and the formation of an appreciable disk.  Next, we considered three hypermassive models (stars~II, III, and~IV) with the same rest mass $M_0$, but different values of the spin parameter $J/M^2$. These models have smaller $M_0$ than star~I, and are therefore less prone to collapse.  Star~II, which has $J/M^2 = 0.85$, collapses when evolved in the rapid-cooling limit, leaving behind a disk.  But without cooling, this model evolves to a stable, uniformly rotating core with a differentially rotating massive disk.  The additional thermal pressure support provided by viscous heating prevents collapse in this case.  In contrast, stars~III and~IV, which have $J/M^2 = 1.0$ and 1.1, respectively, do not collapse even in the rapid-cooling limit. This is sensible because these models have a smaller rest mass than star~I, but larger angular momenta than star~II.  Though the cores of stars~III and~IV contract, they are prevented by centrifugal support from reaching the necessary compaction to become dynamically unstable. In both cases, we find low-density disks surrounding uniformly rotating cores.  However, our simulations do not rule out the possibility that slow accretion of the disk material could eventually drive the uniformly rotating cores to collapse. Disk formation also occurs for star~V, which is differentially rotating but non-hypermassive. Since there exist stable, uniformly rotating models with the same rest mass, the braking of differential rotation in this case does not result in collapse.  However, differentially rotating stars can support larger $T/|W|$ than uniformly rotating stars.  In the case of star~V, there does not exist a uniformly rotating star with the same (high) angular momentum and rest mass, so that mass shedding must take place as viscosity drives the star to uniform rotation.  In the final state, we find a rigidly-rotating core surrounded by a low-density, disk.  Since results obtained from our axisymmetric code are physically reliable only for models which are not subject to nonaxisymmetric instabilities, we evolved stars~I and~IV in 3D to check for such instabilities.  Previous studies in Newtonian gravity have found that the secular, viscosity-driven bar instability in uniformly rotating stars should set in when $T/|W| \\gtrsim 0.14$ \\cite{c69,sl96}.  When general relativity is taken into account, the threshold value can be somewhat higher \\cite{sz98}. Thus, of all of our models, stars~I and~IV have the best chances of developing bars since they have the highest $T/|W|$.  We introduced an initial bar-shaped perturbation and ran these cases in the rapid-cooling limit.  We found that, in both cases, the small initial perturbation decays and no bar is formed.  This is somewhat surprising since $T/|W|$ is well above 0.14 in both of these cases. We plan to address this issue in a future report.   For the evolution of each of our  five models, we find that a massive disk or torus forms in the final state. The disk typically carries $\\sim\\!20\\%$ of the rest mass of the initial configuration. Viscosity transports angular momentum from the interior of the star to the more slowly rotating exterior.  The exterior regions then expand to accommodate the additional centrifugal force, forming a low-density disk.  The inner core becomes rigidly rotating and, in some cases, undergoes gravitational collapse. The disk, however, remains differentially rotating since viscosity acts much more slowly in low-density regions. For cases in which black holes are formed, the mass of the disk may be estimated by integrating the rest-mass density for those fluid elements which have specific angular momentum $j$ greater than the value at the ISCO, $j_{\\rm ISCO}$ [see Eq.~(\\ref{eq:DeltaM0})].  The estimates obtained in this way agree reasonably well with the results of our numerical simulations. Particularly good agreement was found for the case of star~I with no cooling, for which we were able to extend the evolution some $55 M$ beyond the first appearance of an apparent horizon. The rest mass and angular momentum of the disk surrounding the rotating black hole could then be calculated directly and agreed well with the estimates.  We expect that excision techniques will continue to be crucial in establishing the final fate of systems involving matter surrounding black holes.   In a recent paper, Shibata \\cite{s03} numerically simulated collapses of marginally stable, supramassive stars. These supramassive models were constructed using polytropic equations of state  with $2/3 < n < 2$ and rotate at the mass-shedding limit with $0.388 \\leq J/M^2 \\leq 0.670$.  Shibata found that the collapse of these stars results in Kerr black holes and that no more than $0.1\\%$ of the initial rest mass remains outside of the hole.  This result is quite different from our finding that disks are usually present following collapse. However, the initial data for the two calculations are quite different, as well as our inclusion of viscosity. The analysis of \\cite{s03} takes uniformly rotating, unstable configurations as initial data and follows their dynamical evolution. Our calculations begin with differentially rotating, stable configurations and follow both their secular (viscous) and dynamical evolution. Viscosity drives our configurations to uniform rotation. We find that massive disks usually form as by-products of the formation of uniformly rotating cores.  This is due primarily to the transport of angular momentum from the inner to the outer layers.  In addition, all of our models have $0.85 \\leq J/M^2 \\leq 1.1$.  (Large angular momentum is required to generate a hypermassive neutron star in equilibrium.) Since this range is higher than that considered in \\cite{s03}, our models more naturally produce disks \\cite{note:disks}.  All of the phenomena observed in our simulations follow from the braking of differential rotation in strongly relativistic stars.  This may be accomplished by viscosity as shown here, but magnetic fields are likely to be more important.  The fate of the hypermassive remnants of binary neutron star mergers may crucially depend on these effects.  The loss of differential rotation support in such a remnant may lead to delayed gravitational collapse.  This collapse could in turn lead to a delayed gravitational wave burst following the quasi-periodic inspiral and merger signal~\\cite{bss00}.  Our results indicate that if the remnant is not sufficiently hypermassive, collapse may not occur, at least not until the star cools by radiating away its thermal energy.  Understanding the evolution of such merger remnants could aid the interpretation of signals observed by ground based gravitational wave detectors, such as LIGO, VIRGO, GEO, and TAMA.  In addition, short-duration GRBs are thought to result from mergers of binary neutron stars or neutron star-black hole systems \\cite{rosswog,npk01}.  In this scenario, the GRB may be powered by accretion from a massive torus or disk surrounding a rotating black hole.  We have demonstrated that such disks are easily produced during the evolution of hypermassive neutron stars.  The braking of differential rotation may also be important in neutron stars formed in core collapse supernovae. Nascent neutron stars are probably characterized by significant differential rotation (see, e.g., \\cite{zm97,rmr98,ll01,l02} and references therein).  Conservation of angular momentum during the collapse is expected to result in a large value of $T/|W|$. However, uniform rotation can only support small values of $T/|W|$ without shedding mass (\\cite{st83}, Chap. 7).  Thus, nascent neutron stars from supernovae probably rotate differentially. If the induced differential rotation is strong enough, hypermassive neutron stars can form. Their subsequent evolution and final fate then depends on the presence of viscosity or magnetic fields. Such considerations may be important for long-duration GRBs in the collapsar model \\cite{mw99}.  In this model, the GRB is powered by accretion onto the central black hole formed through core collapse in a massive star.  Several interesting astrophysical systems undergo secular evolution in strongly gravitating environments.  In this paper, we have shown that it is possible to study secular effects that occur over many dynamical timescales using hydrodynamic computations in full general relativity.  We consider this an important step toward future numerical explorations of secular effects in other contexts. In particular, we plan to incorporate MHD into our evolution code, as magnetic braking probably acts more quickly than viscosity to destroy differential rotation in many systems, like neutron stars or supermassive stars \\cite{smspapers}.  Our results have also raised the following interesting question:  Under what circumstances are differentially rotating, compressible neutron stars with high $T/|W|$ unstable to nonaxisymmetric modes?  We plan to address this issue in a future report."
    },
    "0402/astro-ph0402028_arXiv.txt": {
        "abstract": "{ We ascribe  the interpretation of the twin kilohertz Quasi Periodic Oscillations (kHz QPOs) of X-ray spectra of  Low Mass X-Ray Binaries (LMXBs) to MHD Alfven wave oscillations in the different mass density regions of the accreted matter at the preferred   radius, and the upper kHz QPO frequency coincides with the Keplerian frequency. The proposed model concludes that the kHz QPO frequencies depend inversely on the preferred  radius, and that theoretical relation between the upper frequency ($\\nt$) and the lower frequency ($\\no$) is $\\no \\sim \\nt^{2}$, which is  similar to the measured empirical relation.   The separation between the twin frequencies decreases (increases) with increasing  kHz QPO frequency if the lower kHz QPO frequency is more (less) than $\\sim$ 400 Hz. % ",
        "introduction": "The launch of the X-ray timing satellite, Rossi X-ray Timing Explorer (RXTE), led to the discovery of Quasi Periodic Oscillations (QPOs) of LMXBs in their  X-ray brightness, with frequencies $\\sim 10^{-1} - 10^{3}$ Hz (see van der Klis 2000 for a recent review). Thereafter, much attention has been paid to the kHz QPO mechanism of LMXB; however the proposed models are still far from explaining  all detected data. The Z sources (Atoll sources), which are high (less) luminous  neutron star (NS) LMXBs (Hasinger \\& van der Klis 1989), typically show four distinct types of QPOs (van der Klis 2000). At present, these are the normal branch oscillation (NBO) $\\nn \\simeq 5-20$~Hz, the horizontal branch oscillation (HBO) $\\nh \\simeq 10-70$~Hz, and the kHz QPOs $\\nt(\\no) \\simeq 300-1300$~Hz that typically occur in pairs in more than 20  sources, with upper frequency $\\nt$ and lower frequency $\\no$. In 11  sources, nearly coherent burst oscillations $\\nb \\simeq 270-620$~Hz have also been detected during thermonuclear Type~I X-ray bursts; these are considered to be  the NS spin frequencies $\\ns$ or their first overtone (see, e.g., Strohmayer \\& Bildsten 2003). Moreover, the existence of a third kHz QPO has also been reported in three low-luminosity sources (Jonker et al. 2000). All of these QPOs but the burst oscillations have centroid frequencies that increase with the inferred mass accretion rate \\mdot. Furthermore, the frequencies  $\\nt$ and $\\no$, as well as the frequencies  $\\nt$ and $\\nh$, follow  very similar relations in five Z sources, and the QPO frequencies  of LMXBs and black hole candidates (BHC) have  a tight and systematical correlation over three orders of magnitude in frequency  (Psaltis et al. 1998, 1999;  Belloni et al. 2002).  Various  theoretical models have been proposed to account for the  QPO phenomenon  in X-ray binaries (for a review see, e.g., Psaltis 2000). In the early detection of RXTE, the upper kHz QPO ($\\nt$) was simply considered to originate from the Keplerian orbital frequency at the preferred  radius, and the lower kHz QPO ($\\no$) is  attributed to the beat of this  frequency with the stellar spin frequency $\\ns$ (Strohmayer et al. 1996; Miller et al. 1998). However, this beat model is inadequate, for the detected  frequency separation ($\\dn \\equiv  \\nt - \\no$)   decreased  systematically with instantaneous \\mdot{} (see, e.g., van der Klis 2000). Later on, general relativistic effects were  invoked to account for kHz QPOs (Stella \\& Vietri 1999; Stella et al. 1999;  Psaltis \\& Norman 2000), which can satisfactorily explain the variation in kHz QPO separation $\\dn$. Moreover, the theory of epicyclic parametric resonance in relativistic accretion disks was proposed (Abramowicz et al. 2003), where the twin kHz  QPOs occur at the frequency of meridional oscillation and the radial epicyclic frequency in the same orbit, which can explain the frequency ratio 3:2 detected in black hole candidates. Although many other feasible ideas have  also proposed, such as the disk seismic model (Wagoner 1999), a two-oscillator model (Osherovich \\& Titarchuk) and the photon bubble model (Klein et al. 1996), no model has yet explained satisfactorily all observed  QPO  phenomena of LMXBs until now.  In this paper, the  MHD Alfven wave oscillation model is proposed, and its predictions and comparisons with the well detected sample sources are shown  in the  figures. ",
        "conclusions": ""
    },
    "0402/astro-ph0402444_arXiv.txt": {
        "abstract": "We present an out-of-core hydrodynamic code for high resolution cosmological simulations that require terabytes of memory.  Out-of-core computation refers to the technique of using disk space as virtual memory and transferring data in and out of main memory at high I/O bandwidth.  The code is based on a two-level mesh scheme where short-range physics is solved on a high-resolution, localized mesh while long-range physics is captured on a lower resolution, global mesh.  The two-level mesh gravity solver allows FFTs to operate on data stored entirely in memory, which is much faster than the alternative of computing the transforms out-of-core through non-sequential disk accesses.  We also describe an out-of-core initial conditions generator that is used to prepare large data sets for cosmological simulations.  The out-of-core code is accurate, cost-effective, and memory-efficient and the current version is implemented to run in parallel on shared-memory machines.  I/O overhead is significantly reduced down to less than 10\\% by performing disk operations concurrently with numerical calculations.  The current computational setup, which includes a 32 processor Alpha server and a 3 TB striped SCSI disk array, allows us to run cosmological simulations with up to $4000^3$ grid cells and $2000^3$ dark matter particles. ",
        "introduction": "Presently, one of the big tasks in cosmology is to determine the concordance model motivated by the results from the cosmic microwave background (CMB), large-scale structure (LSS), and supernovae.  The precision measurement of the matter power spectrum is one key scientific goal and observations of the Lyman alpha (Ly$\\alpha$) forest, the Sunyaev-Zeldovich (SZ) effect, and weak lensing of the LSS are expected to complement existing constraints.  In order to do cosmology with these probes, numerical modelling of the nonlinear physics must achieve a level of accuracy on par with upcoming precision observations.  Cosmological hydrodynamic and N-body simulations are standard tools for modeling nonlinear structure formation in the universe.  Numerical simulations must converge over a large range in scale, mass, and temperature.  Large box sizes are required to capture large-scale correlations and power while high resolution is needed to resolve small-scale, nonlinear structures.  Most simulations to date have sacrificed one for the other because of the limitations in computational resources.  While parallelized numerical codes and optimized mathematical libraries have significantly reduced computation time on high performance computing (HPC) systems, memory limitations remain the bottleneck in expanding the size of simulations.  We describe the implementation of an out-of-core hydrodynamic (OCH) cosmological code for high resolution simulations requiring terabytes of memory.  Out-of-core computation refers to the technique of using disk space as virtual memory and transferring data in and out of main memory at high I/O bandwidth.  Conventional memory is an expensive commodity and currently most HPC systems have a few hundred gigabytes at most.  However, disk space is relatively cheap and disk arrays can provide many more orders of magnitude in storage.  In order to be effective, out-of-core computation must avoid being I/O limited.  Striped SCSI disk arrays are capable of delivering on the order of $100-1000$ MB/s throughput and have sufficient bandwidth.  However, the slow latency presents a major problem for codes requiring non-sequential disk access.  Pioneering work has been done for computational astrophysics.  \\citet{1997SalmonWarren} developed a parallel, out-of-core tree N-body code to run an 80 million particle simulation on a distributed-memory Pentium cluster with a total of 16 processors, 2 GB of memory, and 16 GB of disk space.  The out-of-core paradigm remains a niche that has been relatively unexplored for numerical simulations.  We demonstrate that advances in algorithmic development now make it feasible to do out-of-core computation effectively.  The OCH code is designed for cosmological applications where high mass resolution is required at all scales.  The code is based on the Hydro\\&N-body implementation of \\citet{2004TracPenMACH}.  The hydrodynamics of the baryonic gas is captured using an Eulerian total variation diminishing (TVD) scheme \\citep{1983Harten} that provides high resolution capturing of shocks while preventing unphysical instabilities.  The gravitational evolution of the dark matter is tracked using the standard particle-mesh (PM) scheme \\citep{1988HockneyEastwood}.  The challenges in doing out-of-core computation stem from the requirement that the data domain be decomposed into blocks which can fit in memory.  The hydrodynamics of the ideal gas is a local process that is straightforward to solve on the decomposed domain with the addition of buffer regions.  However, gravity is a global force and Poisson solvers operate on the global density field.  Fast Fourier transforms (FFTs) are the optimal solvers in a standard PM scheme, but they involve global transposes and computing the transforms out-of-core will be intolerably slow.  This problem can be addressed by splitting the gravitational force into long and short-range components, similar to that done in P$^3$M \\citep{1988HockneyEastwood}, mesh-refined PM \\citep{1991Couchman, 1995CouchmanHydra}, and Tree-PM \\citep{1995Xu, 2002Bagla, 2003BodeOstrikerTPM, 2004DubinskiGOTPM} codes.  We implement a two-level mesh scheme where the short-range force is solved on a high-resolution, localized mesh while the long-range force is captured on a lower resolution, global mesh.  The two-level mesh gravity solver is memory-efficient and allows FFTs to be performed on data stored entirely in memory.  In this paper, we describe the out-of-core algorithm and highlight the steps required for doing out-of-core computation.  In particular, we discuss optimizations to reduce I/O overhead by performing disk operations and numerical calculations concurrently.  The two-level mesh gravity solver is described and its accuracy is compared to that of a standard one-level mesh solver.  In addition, we present an out-of-core initial conditions generator that can be used to construct large data sets for high resolution cosmological simulations. ",
        "conclusions": "We have developed an out-of-core hydrodynamic code and an out-of-core initial conditions generator for high resolution cosmological simulations that require terabytes of memory.  The OCH code utilizes a two-level mesh gravity solver and a multi-stepping scheme that significantly reduce the amount of disk operations.  In addition, we have managed to reduce I/O overhead down to less than 10\\% by performing disk operations concurrently with numerical calculations.  The code is cost-effective and memory-efficient.  It has been demonstrated to accurately simulate the nonlinear structure formation in the universe and will provide high mass resolution for cosmological applications.  The first application of the OCH code involves simulating the high redshift intergalactic medium in a WMAP cosmology (Trac \\& Pen 2004 in prep).  We are currently running an out-of-core simulation with $2016^3$ grid cells and $1008^3$ dark matter particles in a $50h^{-1}$ Mpc box.  This high resolution simulation has a comoving grid spacing of $\\Delta x=25h^{-1}$ kpc and a dark matter particle mass resolution of $\\Delta m=7.7\\times10^6h^{-1}M_\\odot$."
    },
    "0402/astro-ph0402391_arXiv.txt": {
        "abstract": "We combine the GMRT low frequency radio observations of SN 1993J with the VLA high frequency radio data to get  a near simultaneous spectrum around day 3200 since explosion. The low frequency measurements of the supernova determine the turnover frequency and flux scale of the composite spectrum and help reveal a steepening in the spectral index, $\\Delta \\alpha \\sim 0.6$, in the optically thin part of the spectrum. This is the first observational evidence of a break in the radio  spectrum of a young supernova. We associate this break with the phenomenon of synchrotron aging of radiating electrons. From the break in the spectrum we calculate the magnetic field in the shocked region independent of the equipartition assumption between energy density of relativistic particles and magnetic energy density. We  determine the ratio of these two energy densities and find that this ratio is in the range: $8\\times 10^{-6}-5\\times 10^{-4}$. We also predict  the nature of the evolution of the synchrotron break frequency with time, with competing effects due to diffusive Fermi acceleration and adiabatic expansion of the radiative electron plasma. ",
        "introduction": "The radio spectrum of a young supernova probes the conditions in the magnetized plasma where this radiation originates from relativistic electrons, which are believed to be accelerated in the interface region of the supernova blast-wave shock and the circumstellar medium.  The radio emission from the nearby supernova SN 1993J is clearly of a non-thermal nature and is argued to be due to synchrotron radiation in a magnetic field amplified in the interaction region and affected by synchrotron self absorption and external free-free absorption (see e.g.\\citet{fra98} and references therein). The most critical parameter of the plasma which affects the synchrotron radiation spectrum, is the strength of the magnetic field. This is often estimated indirectly under assumptions of equipartition of energy between the magnetic fields and that of relativistic particles or by fitting the radio flux density and turnover wavelength. In many classical radio sources, such as supernova remnants (SNRs) like the Crab or Cassiopeia A, or in luminous radio galaxies, the radio spectral index is found to steepen at high frequencies (see e.g. \\citet{kar62,eil97,mye85}). This is due to the so called synchrotron aging of the source, as during the lifetime of the source, electrons with high enough energies in a homogeneous magnetic field will be depleted due to efficient synchrotron radiation compared with the ones with lower energies. An observation of a synchrotron break can yield a measurement of the magnetic field  {\\it independent of the equipartition argument} if the age of the source is known (magnetic field in the crab nebula was measured by \\citet{pik56} (as quoted in \\citet{shk60}) using this technique).   Multi-frequency radio studies of a supernova like SN 1993J, that is bright enough in the radio bands can offer such a possibility. SN 1993J exploded on March 28, 1993. It was the archetypal type IIb supernova and provided a good opportunity to study the extragalactic supernovae in detail for being the nearest extragalactic supernova (3.6 Mpc). In this paper, we discuss the near simultaneous spectrum of SN 1993J obtained by combining the GMRT low frequency data with the VLA high frequency data around day 3200 since explosion. We find a steepening of the spectrum by $\\Delta \\alpha=0.6$ at radio frequency of 4 GHz. We associate this break with the synchrotron cooling. Synchrotron aging  of young supernovae has been discussed by \\citet{fra98} and also mentioned by \\citet{per02}. However we find the first clear observational signature of a spectral break in radio bands of the young supernova SN 1993J. With this break frequency and the independently known age of the SN 1993J, we determine the magnetic field in the supernova. Moreover, we predict how this break frequency will evolve with time, based on a quasi-static evolution of the electron spectrum under the combined effect of acceleration processes, synchrotron losses and adiabatic expansion of the shell where the electrons are confined.  We compare this observationally determined field with the best fit magnetic field under equipartition assumption and thence derive the fraction by which relativistic energy density deviates from  magnetic energy density.   We briefly describe the observations of SN1993J in section 2. In section 3 we combine the data with the VLA data and explore spectral fits with and without synchrotron cooling breaks. In section 4 we explore the cumulative effects of adiabatic expansion of the supernova envelope and energy gain undergone by electrons under diffusive particle acceleration upon the synchrotron cooling affected particle spectrum. In section 5 we discuss the evolution of the break frequency with time and the importance of wide-band spectra for modeling. ",
        "conclusions": "We see in the last section that the magnetic field calculated from the break in the spectral index is $0.33 $ Gauss, which is $\\sim$1.4 times higher than that expected from an extrapolation of \\citet{fra98} ($B=0.24$ Gauss at day 3200). If we took only the synchrotron cooling effect and neglected adiabatic expansion and diffusive Fermi acceleration, we get magnetic field $B=0.19$ Gauss, which is in closer agreement with that of \\citet{fra98}. However, at this young age of the supernova the effects of adiabatic expansion and diffusive Fermi acceleration are likely to be significant as seen in SN1987A \\citep{bal92} and hence these effects should not be neglected.  One can estimate the importance of the diffusive acceleration term, if one is able to follow how the break frequency evolves with time. In Eq. \\ref{synchr_freq}, first term is the contribution of the acceleration and second term is the contribution of adiabatic expansion and synchrotron losses.  We note that for the estimated value of ${\\kappa}_{\\perp}$ at the present epoch, diffusive Fermi acceleration dominates and will continue to dominate over the adiabatic losses until about 20 years since explosion. Therefore, at present epoch the break frequency evolves as ${\\nu}_{break} \\propto t^{-1}$. After 20 years when the acceleration will cease to dominate over adiabatic expansion and the break frequency will increase as: ${\\nu}_{break} \\propto t$.  Since we do not have an independent method to estimate the spatial diffusion coefficient ${\\kappa}_{\\perp}$ for SN 1993J, we have used the linearly scaled (by the respective  compression ratios, $\\rho$) value of ${\\kappa}_{\\perp}$ measured for SN 1987A from direct radio observations (see \\citet{bal92}). However, we can directly calculate the value of  ${\\kappa}_{\\perp}$ from the (measurable) rate of change of ${\\nu}_{break}$, i.e. from the expression: \\begin{equation} \\frac{d {\\nu}_{break}}{dt} = \\frac{2 \\times 10^{24}}{ B_0^{3}} \\left[\\frac{R^2 } {20 {\\kappa}_{\\perp}}t^{-1/2}\\,-2\\,t^{1/2}\\right]\\left[-\\frac{R^2 } {20 {\\kappa}_{\\perp}}t^{-3/2}\\,-2\\,t^{-1/2}\\right] {\\rm Hz/sec} \\end{equation} For the present epoch and with the estimated parameters as above, we calculate that the break frequency is changing at the rate of 1.2 GHz/year. Thus a few more multi-frequency spectral observations across GMRT and VLA bands, separated by a few years, will observationally determine the temporal variation in the break frequency and underline the importance of the diffusive acceleration effects.   The combination of multi-frequency radio spectrum across GMRT and VLA bands is critical for deriving the above results. In Fig. \\ref{fig2} we show a comparison of synchrotron self absorption model (with a single optically-thin power-law index) fitted only to the low frequency data (0.22 GHz to 1.4 GHz) versus such a model fit obtained with only the higher frequency data (1.4 GHz to 22.5 GHz). This comparison shows that while the model fitted only to the low frequencies over-predicts the flux density at high frequencies, the model fitted only to high frequencies on the other hand fails to account for both synchrotron cooling break and seriously under-predicts the low frequency flux densities. The comparison underscores the importance of broad band observations for determining the physical processes taking place in the supernova.    We note the uncertainity in ${\\kappa}_{\\perp}$ for SN 1993J as also emphasised by the referee who suggests that ${\\kappa}_{\\perp}$ may not be a constant, and points out a paper due to \\citet{rey98}. This paper {\\it assumes} that ${\\kappa}_{\\perp}$ is proportional to the particle energy. This dependence may affect the determination of the magnetic field (see Eq. \\ref{synchr_freq}). Only future observations of the break frequency evolution with time will directly put constraints on ${\\kappa}_{\\perp}$ for SN 1993J. We also note that our results are based on the assumption that the acceleration and synchrotron losses are taking place in the same region. However, if the regions of the two processes are not substantially overlapping, the synchrotron break frequency will not be affected by acceleration. Even in that case, the magnetic field is much higher than the equipartiton magnetic field and the plasma is still dominated by the magnetic energy density."
    },
    "0402/astro-ph0402672_arXiv.txt": {
        "abstract": "{\\object{LkH$\\alpha$\\,312} has been observed serendipitously with the ACIS-I detector on board the {\\sl Chandra X-ray Observatory} with 26\\,h continuous exposure. This H$\\alpha$ emission line star belongs to \\object{M\\,78} (\\object{NGC\\,2068}), one of the star-forming regions of the Orion~B giant molecular cloud at a distance of 400\\,pc. From the optical and the near-infrared (NIR) data, we show that LkH$\\alpha$\\,312 is a pre-main sequence (PMS) low-mass star with a weak NIR excess. This genuine T~Tauri star displayed an X-ray flare with an unusual long rise phase ($\\sim$8\\,h). The X-ray emission was nearly constant during the first 18\\,h of the observation, and then increased by a factor of 13 during a fast rise phase ($\\sim$2\\,h), and reached a factor of 16 above the quiescent X-ray level at the end of a gradual phase ($\\sim$6\\,h) showing a slower rise. To our knowledge this flare, with $\\sim$0.4--$\\sim$0.5\\,cts\\,s$^{-1}$, has the highest count rate observed so far with {\\sl Chandra} from a PMS low-mass star. By chance, the source position, $8.2\\arcmin$ off-axis, protected this observation from pile-up. We make a spectral analysis of the X-ray emission versus time, showing that the plasma temperature of the quiescent phase and the flare peak reaches 29\\,MK and 88\\,MK, respectively. The quiescent and flare luminosities in the energy range 0.5--8\\,keV corrected from absorption ($N_{\\rm H} \\approx 1.7\\,10^{21}$\\,cm$^{-2}$) are $6\\,10^{30}$\\,erg\\,s$^{-1}$ and $\\sim$10$^{32}$\\,erg\\,s$^{-1}$, respectively. The ratio of the quiescent X-ray luminosity on the LkH$\\alpha$\\,312 bolometric luminosity is very high with $\\log (L_{\\rm X}/L_{\\rm bol})= -2.9$, implying that the corona of LkH$\\alpha$\\,312 reached the `saturation' level. The X-ray luminosity of the flare peak reaches $\\sim$2\\% of the stellar bolometric luminosity. The different phases of this flare are finally discussed in the framework of solar flares, which leads to the magnetic loop height from 3.1$\\,10^{10}$ to $10^{11}$\\,cm (0.2-0.5\\,R$_\\star$, i.e., 0.5--1.3\\,R$_\\odot$). ",
        "introduction": "Since the first X-ray observations of star-forming regions (SFR) with the {\\sl Einstein Observatory}, young PMS low-mass stars, T~Tauri stars (TTS), are known to be variable in X-rays (\\cite{feigelson81}; \\cite{montmerle83}). TTS, as active stars, display X-ray flares triggered by magnetic reconnection events occurring in their stellar coronae (see reviews by \\cite{feigelson99}, and \\cite{favata03}). The {\\sl Chandra X-ray Observatory}, thanks to its highly elliptical orbit that permits continuous observation over many hours, has collected a real zoo of X-ray flares from TTS in several SFR (e.g., \\cite{preibisch02}, \\cite{feigelson02a}, \\cite{imanishi03}). The X-ray flares of TTS are generally impulsive, displaying a fast rise phase ($\\sim$2\\,h) corresponding to a heating phase, followed by an exponential decay corresponding to a cooling phase (e.g., \\cite{imanishi03}). Only a few long duration flares have been observed from TTS with the previous generation of X-ray satellites, {\\sl ROSAT} and {\\sl ASCA} (\\object{SR13}, \\cite{casanova94}; object{V773\\,Tau}, \\cite{skinner97}; \\object{P1724}, \\cite{stelzer99}). \\cite*{stelzer99} showed that these light curves can be reproduced by the rotational modulation of the exponential decay of a cooling flare. \\cite*{feigelson02a}, investigating the variability of young solar-like stars (M$_\\star$=0.7--1.4\\,M$_\\odot$) in the Orion Nebula Cloud, find several likely long-duration flares but the low count rates of these events prevent a time-dependent spectroscopy analysis.  M\\,78 (NGC\\,2068) is a reflection nebula illuminated by a B1.5V star (HD\\,38563 North; \\cite{mannion84}), and located in the northern part of the closest giant molecular cloud, Orion~B (L\\,1630), at a distance of $\\sim$400\\,pc (\\cite{anthony82}). The star-forming region M\\,78 has been observed with the ACIS-I detector aboard {\\sl Chandra} with 26\\,h exposure on October 18, 2000 (sequence number 200100). We report here the study of the brightest X-ray source detected during this observation, LkH$\\alpha$\\,312, located $\\sim10\\arcmin$ at the South-East of the optical emission nebula M\\,78. This source displays an intense X-ray flare with an unusually long rise phase. We describe the ACIS-I data reduction and the X-ray detection of LkH$\\alpha$\\,312 in \\S\\ref{observation}. The evolutionary status of LkH$\\alpha$\\,312 is determined from the optical and the NIR data in \\S\\ref{status}. The {\\sl Chandra} light curve of the flare is presented in \\S\\ref{variability}. The time-dependent spectra is investigated in \\S\\ref{spectroscopy}. Finally, the different phases of this flare are discussed in the framework of solar flares in \\S\\ref{discussion}. Concluding remarks are presented in \\S\\ref{conclusion}. ",
        "conclusions": "\\label{conclusion}  So far, LkH$\\alpha$\\,312 was considered from $\\sim$30-yr old H$\\alpha$ prism surveys as an unremarkable, weak-activity star at the periphery of M\\,78, until {\\sl Chandra} unveiled its extraordinary activity in X-rays. Our study based on the available optical and IR data shows that it is likely a $< 1$ Myr-old T Tauri star similar to the Sun ($M_\\star \\sim 0.7$--0.75\\,M$_\\odot, R_\\star \\sim 2.6$\\,R$_\\odot$), about as luminous ($L_\\mathrm{bol} \\sim 1.5$\\,L$_\\odot$), and surrounded by a hollow circumstellar disk. It is therefore perhaps a transition object between an active accreting `classical' T Tauri star, and a genuine disk-free `weak-line' T Tauri star.  In contrast with these fairly unoriginal properties of LkH$\\alpha$\\,312, our {\\sl Chandra} observations revealed a particularly long and intense X-ray flare, in addition to a bright quiescent emission. Because of the intensity of the flare ($\\sim 0.5$\\,{\\sl Chandra} cts\\,s$^{-1}$) and its favorable far off-axis position on the ACIS-I detector which prevented pile-up, the quiescent state as well as the rise phase of the flare could be studied at an unprecedented level of detail, yielding spectroscopy resolved on a time scale of down to 1\\,h out of a 26\\,h-long observation. The flare itself is quite remarkable, with its high X-ray luminosity (comparable to that of only a handful of TTS so far) and slow rise, but so is the level of the quiescent emission, which is saturated. The high intensity of the quiescent level was already perceptible in the {\\sl ROSAT/HRI} observation of 1997, so it is likely not a brief, transient phenomenon.  Quantitatively, the X-ray properties derived from our {\\sl Chandra} observations are as follows~: (i) the quiescent emission of LkH$\\alpha$\\,312 is stable over 18\\,h of observation and saturated at a high level of activity ($L_\\mathrm{X}/L_\\mathrm{bol} \\sim 10^{-3}$, with $L_\\mathrm{X} \\sim 6 \\times 10^{30}$\\,erg\\,s$^{-1}$); (ii) LkH$\\alpha$\\,312 also displayed an unusually long and intense X-ray flare, showing a long rise phase extending over several hours. This flare was bright, reaching $L_\\mathrm{X} \\sim 10^{32}$\\,erg\\,s$^{-1}$, i.e., $\\sim$2\\% of the bolometric luminosity. Perhaps more remarkably, the peak temperature was quite high ($kT \\sim 7$\\,keV). (iii) In terms of solar-type magnetic confinement of the plasma, the application of the scaling laws found by Shibata \\& Yokoyama shows that the flare corresponds to a fairly large volume (maximum height of about 1\\,R$_\\odot$), and typical solar coronal densities ($\\sim 10^{11}\\,{\\rm cm}^{-3}$).  The problem is to understand the reasons for these rather extreme X-ray properties, both in the quiescent state and in the flaring state. The first, ``standard'' way is to address the problem in the framework of solar physics. This is how we derived the flare properties (size and density of the confining magnetic loop, strength of the magnetic field). As discussed in \\S\\ref{quiescent}, one could invoke enhanced coronal processes and/or microflaring, 1,000 times more intense than the quiet Sun, or a volume filling factor in the LkH$\\alpha$\\,312 corona 1,000 times larger than in the solar corona, but the physical reasons for such an enhancement are unclear. This is in fact a general problem for pre-main sequence stars~: the mechanism for magnetic field generation leading to the high level of X-ray activity is still not well understood. For instance, the lack of an X-ray luminosity/rotation correlation on large samples like in Orion is perhaps due to the fact that an $\\alpha^2$ (turbulent) dynamo operates at the young TTS stage, as opposed to the standard rotation-convection induced $\\alpha-\\omega$ dynamo which holds in late-type main-sequence stars (\\cite{feigelson03}).  An alternative way is to consider that the enhanced X-ray activity of  LkH$\\alpha$\\,312, and perhaps more generally of at least some other, very X-ray luminous TTS, is triggered by an external cause such as a close,  faint companion like a brown dwarf or even a planet. Indeed, as shown by \\cite*{cuntz00}, {\\it planets} can cause enhanced magnetic activity if they are sufficiently massive and close to the parent star, like 51\\,Peg-like planet-bearing stars. The general idea is that tidal effects can enhance the dynamo by altering the outer convective zone, and/or, if the planet is magnetized, the star and the planet can interact magnetically via shearing and reconnections, for instance as in the star-disk magnetic interaction mechanism proposed by \\cite*{montmerle00} to explain the triple X-ray flare of the YLW15 protostar. In such a framework, the rotation information would be important~: although a slow rotation would not be conclusive, a fast rotation (periods of days or less) could be indicative of a tidally locked, RS CVn-like binary system. We note that the circumstellar disk of LkH$\\alpha$\\,312 is hollow, with no accretion taking place, which is indeed consistent with the possible presence of a planet clearing the inner disk (e.g., \\cite{nelson03}).  Both ground and space follow-up observations will be needed to investigate the coronal properties of LkH$\\alpha$\\,312 as a genuine TTS and its influence on its circumstellar material~: medium resolution optical spectroscopy, to confirm its youthfulness with Li\\,{\\small I}~6707 absorption line measurement, to check that the H$\\alpha$ emission has not changed, and to estimate its projected equatorial velocity;  to determine its rotational period from the photometric modulation of starspots; high-resolution IR spectroscopy to detect the possible presence of a very low mass companion; near-IR high resolution imagery to study a possible transition circumstellar disk; and of course high resolution X-ray spectroscopy to probe the density of its coronal plasma."
    },
    "0402/astro-ph0402358_arXiv.txt": {
        "abstract": "Because of the semi-collisional nature of the solar wind, the collisionless or exospheric approach as well as the hydrodynamic one are both inaccurate. However, the advantage of simplicity makes them useful for enlightening some basic mechanisms of solar wind acceleration. Previous exospheric models have been able to reproduce winds that were already nearly supersonic at the exobase, the altitude above which there are no collisions. In order to allow transonic solutions, a lower exobase has to be considered, in which case the protons are experiencing a non-monotonic potential energy profile. This is done in the present work. In this model, the electron velocity distribution in the corona is assumed non-thermal. Parametric results are presented and show that the high acceleration obtained does not depend on the details of the non-thermal distributions. This acceleration seems, therefore, to be a robust result produced by the presence of a sufficient number of suprathermal electrons. A method for improving the exospheric description is also given, which consists in mapping particle orbits in terms of their invariants of motion. ",
        "introduction": "\\label{introduction}  Most cosmic bodies eject matter into space, but the solar wind is the first and, up to now, only stellar outflow to have been measured in situ \\citep{neu97}. Numerous sophisticated models have been developed since Parker's pionnering papers \\citep{par58, par60} using complicated mechanisms. The different acceleration mechanisms, the origin of the high-speed solar wind as well as the associated problem of coronal heating have been recently reviewed in a comprehensive way by \\citet{cra02}. However the solar wind acceleration and its properties are still far from being well understood. The reason of this difficulty is that the solar wind is neither a collision-dominated medium nor a collisionless one. The Knudsen number $K_n$, which is defined as the ratio of the particle mean free path and the density scale height, is close to unity at Earth's orbit (see e.g. \\citet{hun72}). This means that neither the hydrodynamic approach nor the pure collisionless one (also called exospheric) are fully appropriate to model the solar wind expansion and to explain its observed properties.  The classical fluid approach is applicable for the extreme regime when $K_n\\ll1$, that is when the medium is collision-dominated. In this case, the particle velocity distribution functions (VDFs) are Maxwellians as the medium is assumed to be at local thermodynamic equilibrium. The Euler or Navier-Stokes approximations are applicable and produce a thermally driven wind out of the hot solar corona. There are two problems with this approach. Firstly, the particle VDFs might not be Maxwellians at the base of the solar wind. Secondly, the particle VDFs $are$ $not$ Maxwellian in the solar wind. There is an increasing number of both theoretical \\citep{vin00, leu02}, and observational evidences \\citep{ess00}, which tend to show that non-thermal VDFs can develop and exist in the high corona and even in the transition region. This is because in a plasma the particle free paths increase rapidly with speed $(\\propto v^4)$, so that high energy tails can develop for Knudsen numbers as low as $10^{-3}$ \\citep{sho83}, i.e. even in collisional plasmas. A fortiori, high energy tails can be expected to be found in the weakly collisional corona and solar wind acceleration region. Indeed, it is well known that the solar wind electron VDFs permanently exhibit non-thermal tails that can be modelled by a halo Maxwellian population (for e.g. \\citet{fel75}) or by the power law part of a generalized Lorentzian or Kappa function \\citep{mak97b}. In the frame of the fluid approach, which intrinsically cannot handle suprathermal tails, the effect of non-thermal VDFs on the solar wind acceleration can be understood through an increase of the heat flux \\citep{hol78, olb81}.  An alternative way of taking into account the possible effects of coronal non-thermal distributions is to use a kinetic approach. Among the various kinetic approaches for the solar wind, the simplest one is probably the exospheric one, which totally neglects binary collisions between particles above a given altitude called the exobase. The first solar wind model of this type was developed by \\citet{cha60} by analogy with the evaporation of planetary atmospheres. This first exospheric model, modelling the radial expansion of the solar corona from the thermal evaporation of hot coronal protons out of the solar gravitational field, produced a solar breeze. The subsonic speed obtained by the theory was partially due to an inadequate assumption: the electrostatic field was taken so as to ensure hydrostatic equilibrium \\citep{pan22, ros24}, which is inconsistent with an expanding atmosphere. The improved exospheric models by \\citet{jen63} and \\citet{bra66} were the first to be able to reproduce supersonic solar wind flows. In these models, multiple exobase locations were assumed, in order to take into account the variable mean free path of the particles as a function of their velocity. However these models still used the inadequate Pannekoek-Rosseland electric field as an imposed external solution. The actual (outward) ambipolar electric field, which ensures plasma quasi-neutrality and zero electric current is greater, thereby accelerating protons to greater speeds. Models using this correction \\citep{lem71a, joc70} produced supersonic winds, but too small speeds for explaining the fast solar wind $(\\sim 700-800\\ km\\ s^{-1})$.  More recently, \\citet{mak97a} have generalized these calculations by considering non-Maxwellian velocity distribution functions for the electrons in the corona, e.g. generalized Lorentzian or Kappa functions. With such non-Maxwellian distributions having suprathermal electron tails, a higher electrostatic potential is needed to ensure zero charge and current, therefore producing larger terminal bulk speeds. In essence, this comes about because the electron tail tends to increase the escaping electron flux, so that, to preserve quasi-neutrality, the electrostatic potential increases in order to trap more electrons, which in turn accelerates the protons outwards. This model yields a reasonable description of bulk solar wind properties, giving densities, temperatures and speeds within the ranges observed at 1 AU, even though the details of the VDFs are not reproduced, as expected since collisions are neglected. Its major interest is the prediction of high speeds without assuming extremely large coronal temperatures and/or additional heating of the outer corona, as is needed in hydrodynamic models. More basically, the main achievement of exospheric models is to furnish a possible driving mechanism for the fast solar wind, with a single assumption: the suprathermal electron VDF at the exobase.  However, the \\citet{mak97a} model cannot be applied for low altitude exobases, which is the case of coronal holes, from where emanates the fast solar wind. In these deep coronal layers, the gravitational force acting on the protons is stronger than the electric one, so that the total potential for the protons is attractive out to some distance where the two forces balance each other. Farther out, the outward electric force dominates. This means that the total potential energy for the protons is not monotonic, presenting a maximum at a certain distance from the exobase \\citep{joc70} and therefore not all the protons present at the exobase are able to escape. The presence of such a maximum is not taken into account in the \\citet{mak97a} model, nor in the \\citet{lem71a} one, since both models started beyond the vicinity of the sonic point.  The purpose of the present paper is to study the effect of non-thermal electron VDFs in the frame of a transonic exospheric solar wind model. We use a special technique described by \\citet{joc70} and \\citet{kha98} consisting in mapping particle orbits in terms of the invariants of motion. The same problem has been recently considered by \\citet{lam03} using a different formulation, which involves an approximation relative to the escaping particles rate (see Appendix \\ref{emucalcul}). The present work, which uses the same approximation, sets the basis of an exact exospheric description of the solar wind acceleration. The reader should, however, have in mind that the present model is still a very simplified one, which does not pretend to describe all solar wind properties. Instead, it may be useful to determine some basic aspects and a possible driving mechanism of the solar wind, avoiding ad hoc assumptions on energy dissipation and using as few free parameters as possible.  In section \\ref{basics} we recall the basics of an exospheric solar wind model. In section \\ref{protons} we outline the difficulties arising when dealing with a non-monotonic potential energy for the protons and describe the technique used to calculate the interplanetary electrostatic potential. In section \\ref{electrons} we deal with non-Maxwellian electron distribution functions and consider three cases: a Kappa, a sum of two Maxwellians (a core and a halo) and a sum of a Maxwellian and a Kappa. The latter distribution is rather general and reproduces the main features of VDFs observed in space plasmas: a Maxwellian profile at low speeds and a high energy tail with a power law shape. In section \\ref{results} we describe the results of our model and compare them with observations. A summary and final remarks are given in section \\ref{conclusion}. ",
        "conclusions": "\\label{results}  As explained in section \\ref{basics}, the exobase location is approximated to be at $r_0=1R_\\odot$ with no serious impact on the results. We assume for the temperatures at the exobase $T_{e0}=10^6\\ K$ and $T_{p0}=2T_{e0}$, in the range of values observed in coronal holes \\citep{cra02}. The density at the exobase does not affect the results of velocity or temperature and is just a multiplicative factor in the density profiles.  Let us first consider a Kappa VDF for the electrons as described in section \\ref{kappasect}. The calculated electric potential and the total potential energy of the protons are plotted in Figures \\ref{phikappa} and \\ref{psikappa} respectively for different values of $\\kappa$ ranging from $\\kappa=6$ to $\\kappa=2.5$, a case with a conspicuous suprathermal tail. Note that we use $\\kappa>2$ in order for the energy flux to be finite. One sees that the value of the maximum of potential increases and its distance $r_{max}$ decreases as $\\kappa$ decreases. This is because with more suprathermal electrons, a stronger electric potential is needed to preserve quasi-neutrality. For a Maxwellian VDF $(\\kappa\\rightarrow\\infty)$, the total proton potential energy increases monotonically and tends asymptotically to zero (remaining always negative).  The bulk speed - the ratio between flux and density - is shown in Figure \\ref{vitkappa} for a Kappa distribution. A high terminal bulk speed $(>700\\ km\\ s^{-1})$ is obtained when the suprathermal tail is conspicuous $(\\kappa=2.5)$. This is due to the large value of the maximum in ion potential energy $(\\approx14k_bT_{p0})$, which is transformed into kinetic energy of the escaping protons as they are accelerated above $r_{max}$. An important remark is that the major part of this high terminal bulk speed is obtained within a small heliocentric distance $(\\approx10R_\\odot)$; this is due to the large acceleration represented by the large slope of the potential above $r_{max}$. Note that this is the largest terminal bulk speed obtained by this model with a Kappa VDF.  The density profiles are shown in Figure $\\ref{densekappa}$. One sees that they are nearly independent of the value of $\\kappa$. The density at $1AU$ depends on the one taken at the exobase. In this figure we assumed an exobase density $n_0=1.8\\cdot 10^{13}\\ m^{-3}$, which corresponds to the density of a coronal hole extrapolated to $r_0=1R_{\\odot}$ as given by \\citet{kou77} in line with recent studies on atmospheric and coronal electron densities \\citep{ess99}. With this density, the model yields a density of about $6\\ cm^{-3}$ at $1AU$, of the same order as observed in situ. Note that the rest of the results do not depend on the density.  This analysis bears out previous results of an exospheric model using Kappa VDFs \\citep{mak97a}. There are however two basic differences. In the present work the velocity profiles span the whole domain from the subsonic to the supersonic regime, which was not the case when the exobase was located above $r_{max}$. The second difference is that we obtain high bulk speeds with more reasonable temperatures at the exobase. Note however that a direct comparison cannot be done because of a slightly different exobase definition resulting in different proton temperatures at the exobase. In any case both models can produce high bulk speeds without assuming an additional (ad hoc) heating mechanism in the outer corona, as is generally postulated in hydrodynamic models. Furthermore, the fact that a faster wind is obtained with low values of $\\kappa$ agrees with observations showing that VDFs have large suprathermal tails in the fast solar wind but are closer to a Maxwellian in the slow wind \\citep{mak97b}.  The large suprathermal tails for low values of $\\kappa$ have another important consequence. They make the electron temperatures increase considerably with distance up to a maximum $(\\approx7$x$10^6\\ K)$ within a few solar radii. This maximum in electron temperature is smaller for larger values of $\\kappa$ and disappears as $\\kappa\\rightarrow\\infty$ $(Fig.\\ref{tempekappa})$ as does the maximum in the total potential energy of protons. This temperature increase is a direct consequence of filtration of the non-Maxwellian VDF by the attracting electrostatic potential \\citep{scu92}. This large temperature increase is not observed, which suggests that Kappa functions may not be adequate to model VDF having suprathermal tails in the corona.  Let us now consider the results obtained with electron distributions made of a sum of two Maxwellians or a sum of a Maxwellian core and a Kappa halo. On the whole the results are rather similar. For the same acceleration we obtain approximately the same temperature increase as in the Kappa case as we can see in Figure \\ref{tempecontrib} $(\\alpha_0=0.03, \\tau_0=5$ and $\\kappa=2.5)$, which corresponds to a terminal bulk speed of $\\sim770\\ km\\ s^{-1}$. We deduce that the temperature increase is not an artefact of Kappa VDFs, but a general behavior of non-thermal distributions. For a given terminal bulk speed the filtration mechanism results in the same temperature increase. It is important to note that collisionless models are expected to give correct electron temperatures \\citep{mey98}, because collisions with other particles do not significantly affect the electron energy, whereas collisions between electrons do not change their total temperature. In any case we should remind that the present model is still a zero-order one, which is intended to explore the basic physics of the wind acceleration, but should not be expected to reproduce all observations in a detailed way, because it involves very few free parameters.  Figure \\ref{tempecontrib} shows also the contributions of the different particle species to the total electron temperature. At large distances the temperature profile is the sum of a term $\\propto r^{-4/3}$ plus a constant. The $r^{-4/3}$ term comes from the isotropically distributed electrons (ballistic and trapped) confined by the heliospheric electric potential, which is found to have the same radial variation at large distances. The constant term comes from the parallel temperature of the escaping electrons. This agrees with analytical results by \\citet{mey98} that do not depend on the VDFs in the corona, but were obtained with a monotonic proton potential profile. When the proton potential energy is non-monotonic, the asymptotic electron temperature profile is still the sum of a term varying as $r^{-4/3}$ plus a constant \\citep{mey03}, but the relative importance of these terms is not necessarily the same.  In Figure \\ref{contourmm} we show results for a sum of two Maxwellians. The diagram shows contours of the terminal bulk speed (at 1AU) as function of $\\alpha_0$ and $\\tau_0$, where we can see that this kind of VDF is able to explain the values of the fast wind speed $(~700-800\\ km\\ s^{-1})$. The terminal bulk speed increases with increasing $\\tau_0$, which is not surprising since the halo temperature increases. Concerning the parameter $\\alpha_0$, the terminal bulk speed behaves differently. One sees that the terminal speed has a maximum for some value of $\\alpha_0$ (for a given $\\tau_0$). This is reasonable since for $\\alpha_0\\rightarrow 0$ and $\\alpha_0\\rightarrow\\infty$ we have just one Maxwellian with temperature $T_{e0}$ and for all values of $\\alpha_0$ between these limits, the electron VDF is non-thermal giving rise to a more important acceleration because of the velocity filtration mechanism. In addition, close to these limits the terminal bulk speed becomes independent of $\\tau_0$ as there is just only one VDF. That makes the contour lines to be vertical.  When using a sum of a Maxwellian core and a Kappa halo (for instance with $\\kappa=2.5$) the contour plot of the terminal bulk speed is quite similar to the previous one as is shown in Fig.$\\ref{contourmk}$. The main difference is the higher acceleration, which is due to the use of the Kappa function. The maxima for a given $\\tau_0$ are now displaced to the right (to larger values of $\\alpha_0$). This is due to the fact that for $\\alpha_0\\rightarrow\\infty$ the VDF is now just a single Kappa accelerating the wind more than a single Maxwellian. It is important to note that we can obtain very high wind speeds even without using a Kappa VDF (Fig.\\ref{contourmm}). This shows that the acceleration is not just a consequence of the Kappa function, but results from non-thermal distributions, as expected. There are no restrictions on $\\alpha_0$ and $\\tau_0$ (as in the case of the Kappa VDF where we have to take $\\kappa>2$), but one should constrain these parameters by coronal observations or by future in situ measurements close to the Sun."
    },
    "0402/astro-ph0402629_arXiv.txt": {
        "abstract": "We present new \\x\\ results on the field around the NGC\\,346 star cluster in the SMC. This continues and extends previously published work on {\\it Chandra} observations of the same field. The two \\x\\ observations were obtained, respectively, six months before and six months after the previously published  {\\it Chandra} data. Of the 51 X-ray sources detected with \\x, 29 were already detected with {\\it Chandra}. Comparing the properties of these X-ray sources in each of our three datasets has enabled us to investigate their variability on times scales of a year. Changes in the flux levels and/or spectral properties were observed for 21 of these sources. In addition, we discovered long-term variations in the X-ray properties of the peculiar system \\hd, a luminous blue variable star, that is likely to be a colliding wind binary system, which displayed the largest luminosity during the first \\x\\ observation. ",
        "introduction": "The giant \\hii\\ region N66 of the Small Magellanic Cloud (SMC) is the largest star formation region of that galaxy. It notably harbors NGC\\,346, a young cluster containing a wealth of massive stars, and \\hd, a peculiar system that underwent a LBV-type eruption at the end of the last century. \\ro\\ and {\\it ASCA} observations have also revealed the presence of a few X-ray binaries (XRBs) in this field \\citep[see e.g. ][]{ha00,ts99}, some of which were later found to harbor pulsars \\citep[for a review, see ][]{ha04}.  Recently, a new generation of powerful X-ray observatories has been launched and the XMEGA\\footnote{http://lheawww.gsfc.nasa.gov/users/corcoran/xmega/xmega.html} consortium used this opportunity to observe the X-ray emission from the NGC\\,346 field in greater detail. The results of a 100~ks $Chandra$ observation have been given in two previous articles. In Naz\\'e et al. (2003a, hereafter Paper I), we reported the first detections in the X-ray domain of the cluster and \\hd. The X-ray emission from NGC\\,346 appeared tightly correlated with the cluster's core, while that of \\hd\\ was found to be bright and variable on the short timescale. In Naz\\'e et al. (2003b, hereafter Paper II), we analyzed the X-ray properties of the 75 point sources discovered in the field, and found possible counterparts to 32 of these sources. We refer the reader to these papers for detailed results and a thorough introduction to the importance of the NGC\\,346 field. This third Paper of the series is meant to supplement the two previous ones. Here, we continue our investigation of the field with the analysis of two \\x\\ datasets, which provide a wider field compared to the {\\it Chandra} data and allow us to examine the issue of source variability on timescales of seconds to months.  The observations and data reduction are presented in \\S2. Next, the properties of the sources detected by \\x\\ are discussed in \\S3, where they are also compared to previous data available. In \\S4, we focus the discussion on \\hd, the supernova remnant SNR\\,0056$-$72.5, and the XRBs. We conclude in \\S5. ",
        "conclusions": "In this third Paper on the X-ray emission from the NGC\\,346 field, we have analyzed \\x\\ data taken six months before and after our {\\it Chandra} observations. 51 sources were detected with \\x, 29 of them being in common with the {\\it Chandra} data analyzed in Paper II. A comparison of the X-ray observations of the field has revealed the variations of 21 of these 51 X-ray sources, 10 of them being known as X-ray binaries or XRB candidates. Another varying source is \\hd, which appears brighter during the first \\x\\ observation. However, the exact nature of these changes (secular or phase-locked variations ?) is not yet known, and requests additional X-ray data to be elucidated."
    },
    "0402/astro-ph0402135_arXiv.txt": {
        "abstract": "We introduce and implement two novel ideas for modeling lensed quasars.  The first idea is to require different lenses to agree about $H_0$.  This means that some models for one lens can be ruled out by data on a different lens.  We explain using two worked examples.  One example models 1115+080, 1608+656 (time-delay quads) and 1933+503 (a prospective time-delay system) all together, yielding time-delay predictions for the third lens and a 90\\%-confidence estimate of ${H_0}^{-1}=14.6_{-1.7}^{+9.4}$~Gyr ($H_0=67_{-26}^{\\,+\\;9}\\rm\\;\\,km\\;s^{-1}\\,Mpc^{-1}$) assuming $(\\Omega_M=0.3,\\Omega_\\Lambda=0.7)$. The other example models the time-delay doubles 1520+530, 1600+434, 1830-211, and 2149-275, which gives ${H_0}^{-1}=14.5_{-1.5}^{+3.3}$~Gyr ($H_0=67_{-13}^{\\,+\\;8}\\rm\\;\\,km\\;s^{-1}\\,Mpc^{-1}$). Our second idea is to write the whole modeling software as a highly interactive Java applet, which can be used both for coarse-grained results inside a browser and for fine-grained results on a workstation.  Several obstacles come up in trying to implement a numerically-intensive method thus, but we overcome them. ",
        "introduction": "Some aspects of modeling lensed quasars are much as they were just after the discovery of the first double quasar 0957+561.  In one of the earliest lens-modeling papers, \\cite{young81} are concerned with some now-very-familiar issues: the effect on image positions of both the main lensing galaxy and other galaxies, the time delays predicted by the models, the non-uniqueness of the models despite the adequacy of the data, and the desirability of supplementary data about the lens, such as velocity dispersions or X-rays.  But other aspects of the subject these days would have been unimaginable in 1981.  The first double quasar has been joined by dozens of others: the CASTLES survey compilation \\citep{castles} currently lists 76 secure multiple-image galaxy-lens systems, with image positions measured to the mas level, even more precisely if there are compact radio sources.  And the time delay for 0957+561, for which Young et al.'s preliminary estimate was about 5 years, is now measured as $423\\pm1\\,\\rm days$ \\citep{oscoz01}, along with time-delay measurements for eight other systems \\citep{schechter97,lovell98,biggs99,burud00,cohen00,% burud02a,burud02b,fassnacht02,hjorth02}.  These excellent data demand new, more automatic, and more portable software tools for modeling the lenses.  Thus motivated, we have developed a new code, {\\em PixeLens,} which we present in this paper.  {\\em PixeLens\\/} works by reconstructing a pixelated mass map for the lens, an idea we first implemented in \\cite{sw97}.  Most lens modeling codes work by fitting a parametric functional form---for example, \\cite{young81} fitted King models; {\\em gravlens\\/} by \\cite{keeton01} is a modern example, offering the user a large choice of parametric models. There are other possibilities too: \\cite{trotter00} reconstruct lenses non-parametrically as we do, but use multipole expansions rather than pixels.  {\\em PixeLens\\/} generates large ensembles of models rather than one or a few mass maps, as a way of addressing the non-uniqueness problem.  We introduced this strategy in \\cite{ws00} for pixelated models.  \\cite{kw03} use a somewhat similar strategy, but with parametrized models.  But {\\em PixeLens\\/} also brings two completely new features, one astrophysical, one computational: \\begin{itemize} \\item It can model several lenses simultaneously, enforcing consistency of $H_0$ across different time-delay lenses. As a result, lenses can be used to constrain other lenses in an interesting and surprising way; \\item The code is highly portable and can run without change as a standalone program, or as a Java applet inside a web browser.  In the online version of this paper, it is available as an alternative version of Figure~\\ref{gui}. \\end{itemize}  The best way to explain what we have implemented and what problems remain is through an example.  So in the following section we will work through a simultaneous reconstruction of three lenses: the time delay quads 1115+080 and 1608+656, and the ten-image system 1933+503. This will be the main part of the paper.  We turn to some computational issues after that. ",
        "conclusions": "This paper continues our work on modeling lens quasars using pixelated mass maps.  We elaborate on ideas introduced in earlier papers \\citep{sw97,ws00,sw01,rsw03}: (a)~formulating lens reconstruction as an undetermined linear inverse problem, (b)~searching through model space to infer $H_0$ (or alternatively, predict time delays) together with systematic uncertainties, (c)~predicting the main features of any Einstein ring, and if observed, estimating the size of the host galaxy.  But more importantly, we introduce and implement two new ideas: (i)~reconstructing several time-delay lenses simultaneously while requiring them to agree about $H_0$, (ii)~making the software so portable that it can be run inside a web browser while reading this paper.  The results in this paper are from two worked examples of seven lenses in all. \\begin{enumerate} \\item First we model the time-delay quads 1115+080 and 1608+656, and the `dec' 1933+503.  The two time-delay lenses help check against our earlier work using a different code. By coupling them with 1933+503 we can predict time-delays for the latter that are not conditional on $H_0$.  We also obtain $$ \\matrix{ {H_0}^{-1}=14.6_{-1.7}^{+9.4} \\rm\\ Gyr & (H_0=67_{-26}^{\\,+\\;9} \\rm\\ local\\ units) & \\hbox{at 90\\% confidence.} \\cr }$$ \\item Then we model the time-delay doubles 1520+530, 1600+434, 1830-211, and 2149-275, and here the main result $$\\matrix{ {H_0}^{-1}=14.5_{-1.5}^{+3.3} \\rm\\ Gyr & (H_0=67_{-13}^{\\,+\\;8} \\rm\\ local\\ units) & \\hbox{at 90\\% confidence} }$$ \\end{enumerate} The reference cosmology has $(\\Omega_M=0.3,\\Omega_\\Lambda=0.7)$. Note that the two worked examples contain no lenses in common, so the close similarity of the two derived values of $H_0$ (1 and 2 above) is very encouraging.  The well-known correlations---(i)~steeper lenses give higher $H_0$, and (ii)~more mass in the annulus covering the images gives lower $H_0$---are present for most of the lenses, and of these (ii) is better. However, such simple correlations appear to get weaker as $H_0$ is better constrained, which happens as a consequence of coupling lenses.\\footnote{A famous comment by R.O. Redman [quoted by \\cite{gmbp}] comes to mind: hearing ``after all, a star is a pretty simple thing,'' Redman retorted ``at a distance of 10 parsecs {\\it you'd\\/} look pretty simple''.} Furthermore, lenses with no time delays, like 1933+503, as well as lenses with very asymmetric mass distributions, like 1608+656, are not expected to show the above correlations.  At this stage we remain somewhat cautious about the Hubble-time estimates.  As we explained in subsection~\\ref{prior-sec}, the distribution of models we obtain is conditional on the prior we use. In the past, when the uncertainties on the Hubble time were very large, fine-tuning the prior was not so important.  But now, with the uncertainties shrinking down from better data and improved modeling, improving the prior is probably the next priority.  \\appendix"
    },
    "0402/astro-ph0402245_arXiv.txt": {
        "abstract": "We use physical constrains imposed from the H-Theorem and from the negative nature of the heat capacity of self-gravitating thermodynamically isolated systems to investigate some possible limits on the stellar polytrope index $n$ within the domain of a classical non-extensive kinetic theory. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0402/astro-ph0402590_arXiv.txt": {
        "abstract": "We investigated the time lags and the evolution of the cross spectra of Z source GX~5-1, observed by the Rossi X-ray Timing Explorer (RXTE), when it is in the horizontal branch oscillations. We showed that the time lags of 3 horizontal branch oscillations are related to the position on the hardness intensity diagram. All of the three QPOs were shown to have hard time lags. However on the cross spectra, one is in a `dip', one in a `bump', the other has no so obvious characteristic. The time lags of two of the QPOs decrease with QPO's frequency, while the other has a trend increasing with its frequency. Moreover, in the normal branch, we found no significant time lags in the present observational data. ",
        "introduction": "The neutron-star low-mass X-ray binaries (LMXBs) can be divided into two types on the basis of pattern traced out in the X-ray intensity-hardness or color-color diagram (HID or CCD, respectively; Hasinger \\& van der Klis 1989). One is Z source which traced out a ``Z\" pattern on HID or CCD, the other is atoll source which traced out a ``C\" pattern on CCD. From the top to the bottom of the Z pattern, the three limbs of the pattern traced out on HID or CCD is called the Horizontal Branch (HB), Normal Branch (NB), and the Flaring Branch (FB), respectively. The temporal properties of Z source depend on the position on the Z pattern. On the HB and the NB, there are varying quasi-periodic oscillations (QPOs) with frequency less than 100 Hz, called horizontal branch oscillation (HBO) and normal branch oscillation (NBO).  Recently, Jonker et al. (2002) studied in detail the power-density spectra (PDS) of the Z source GX~5-1. They showed that there are four QPO components on the HB, which are HBO and its three harmonics. Using the EXOSAT data, van der Klis et al. (1987) had studied the time lags (or phase lags) of the HBO and the low frequency noise (LFN) by cross-correlation spectra, showing a complex timing behavior. There are hard time lags decreasing with the HBO frequency and soft time lags in LFN increasing with the Fourier frequency. They argued that the hard time lags of HBO are due to Comptonization of soft photons, and the soft time lags of LFN due to the evolution of energy spectrum of the shot. Using the Ginga data, Vaughan et al. (1994) studied one harmonic with higher Fourier frequency and the relation between time lags of NBO and its energy, in which a jump at 3.5 keV was found. Their results are similar to that from van der Klis et al. (1987).  Thanks to the large effective collective area and high time-resolution of the Rossi X-ray Timing Explorer (RXTE), we can investigate the cross-spectrum and PDS of GX~5-1 in more details. In this paper, we report the time lags of HBO and its harmonics of GX~5-1 basing on the RXTE data. We found that the time lag behaviors are related to the HBO fundamental frequency. In section 2, we describe the method of data analysis. In section 3, the results are presented. We discuss the results briefly in the last section. ",
        "conclusions": "We have performed the analyzed of the RXTE/PCA data of GX~5-1 when it was on the HB and the NB. Beside that the soft time lags of the LFN and the hard time lags of QPOs were found to be similar to the results of Vaughan et al. (1994), we discovered some finer characteristics in the cross spectra around the centroid frequency of the QPOs. The time lag of the HBO is in a dip while the harmonic in a bump. The sub-HBO also shows a hard time lag. Due to the contamination of the LFN, the time lag of the sub-HBO increases with the centroid frequency of the sub-HBO. When the effect of the LFN was corrected, this trend disappear. The QPOs in the horizontal branch show hard time lag. In present RXTE data and energy bands, we didn't find any significant time lags in the NBO's frequency range. It may be caused by that the lowest energy channel for calculating cross spectra is above the jump energy (3.5 keV) of time lags (Vaughan et al. 1999).  If box series numbers in the HID represent the mass accretion rate of the source, the evolution of the cross spectrum along the track suggests that the cross spectrum should vary with the mass accretion rate.  There are several kinds of models proposed to interpret the time lags: 1) the Comptonization models, e.g., the uniform corona model (Payne 1980), 2) the non-uniform corona model (Kazanas, Hua, \\& Titarchuk 1997), and 3) the drifting-blob model (B\\\"ottcher \\& Liang 1999), they only explain the hard time lags and the energy dependence of the time lags. The soft time lags cannot be explained. However, those models have some problems physically. For example, for producing measured time lags in a static Compton cloud, a hot corona with a large radius (for example, for Cyg~X-1, the radius extent of the hot corona exceeds $10^{10} r_g$), it is physically unrealistic (see the review by Poutanen 2000). The coherence function might reflect the dynamical properties of the corona in the corona models(Nowak et al. 1999a, 1999b, Ji et al. 2003). The loss of coherence in high energy channels might show that the corona of GX~5-1 is dynamical. This is consistent with the above suggestions.  To explain the soft time lags and the time lag evolution in GRS~1915+105 and neutron-star binary Cir~X-1 (Cui 1999; Reig et al. 2000, Qu, Yu \\& Li 2001), a two-layer corona model has been proposed (Nobili et al. 2001, Qu, Yu \\& Li 2001). In this model, the evolution of time lags of GX~5-1 with mass accretion can be explained, but the increase of the time lag near the transition between the HB and the NB can not be explained.  Both shot profile properties and Comptonization of photons can introduce time lags. The numerical simulations showed that the energy-dependent shot profiles can produce low-energy phase lags in the cross spectrum at typical frequencies for the shot time scale (tenths of a Hz to a few Hz) without noticeably affecting the cross spectrum at higher frequencies(Shibazaki et al. 1988). The shots are thought to originate near or at the neutron star surface as material falls through the magnetosphere and onto the surface of the neutron star. The delays should be of the order of the free-fall time. But the free-fall time is not longer than 0.5 ms. Although the shot models can produce almost any time-delay spectrum like Comptonization models, they also have problems on physical grounds (see Vaughan 1994, Poutanen 2000, Qu et al. 2002)."
    },
    "0402/astro-ph0402073_arXiv.txt": {
        "abstract": "s{ An equation of state for 2-flavor quark matter (QM) with diquark condensation under the conditions for compact stars -$\\beta$-equilibrium, charge and color neutrality- is presented. Trapped antineutrinos prevent the formation of the diquark condensate at moderate densities above a critical value of the antineutrino chemical potential $\\mu_{\\bar \\nu_e}^c$. The following consequences are presented: 1) The star develops a 2-phase structure ($\\mu_{\\bar \\nu_e}\\geq \\mu_{\\bar \\nu_e}^c$): a color superconducting QM core and a normal QM shell. 2)During the cooling, when the temperature is small enough  ($T<1$ MeV) the antineutrino mean free path becomes larger than the thickness of the normal QM shell and the antineutrinos get untrapped in a sudden burst. The energy release is estimated as $\\simeq 10^{52}$ erg and an antineutrino pulse is expected to be observed. } ",
        "introduction": "It has been proposed that cold dense quark matter should be in a superconducting state with the formation of a diquark condensate\\cite{Alford:2000sx,Blaschke:uj}. The consequences of the diquark condensation for the configuration and the cooling behaviour of compact stars have been broadly studied \\cite{Blaschke:1999qx,Page:2000wt,Blaschke:2000dy,Blaschke:2003yn} and the question if this phase can be detected by the signatures still remains \\cite{Blaschke:2003rg}.  Also the engine for the most energetic phenomena in the universe like supernova explosions and gamma ray burst does not have a satisfactory explanation yet \\cite{Piran:2002gc}and it has been proposed that the energy involved could be related with the occurence of the color superdonductivity phase \\cite{Ouyed:2001cg,Hong:2001gt}.  Since the pairing energy gap in quark matter is of the order of the Fermi energy, the diquark condensation gives a considerable contribution to the equation of state (EoS) that is estimated of the order of $({\\Delta}/{\\mu})^2$. Disregarding relativistic effects, the total binding energy release in the core of a cooling protoneutron star has been estimated as $({\\Delta}/{\\mu})^2M_{{\\rm core}}\\simeq~10^{52}$ erg. But, if relativistic effects are considered, the gravitational mass defect of the cooling star decreases  when diquark condensation is included and there is no explosive process \\cite{Blaschke:2003yn} possible since the color superconductivity transition is second order.  In this work a new mechanism of releasing the energy in an explosive way is presented (for the original idea see \\cite{Aguilera:2002dh}). During the collapse of a protoneutron star antineutrinos are produced by the $\\beta$-processes and remain trapped due to the small mean free path. This increases the asymmetry in the system and therefore the diquark condensation is inhibited at moderate densities. So, a two-phase structure developes in the star: a superconducting interior and a sorrounding shell of normal quark matter, the latter being opaque to  antineutrinos for $T\\geq 1$ MeV \\cite{Reddy:1997yr}. In the cooling process the antineutrino mean free path increases above the size of this normal matter shell and  an outburst of neutrinos occurs releasing an energy of about $10^{51}-10^{52}$ erg. This first order phase transition leads to an explosive phenomenon in which a pulse of antineutrinos could be observed.  \\subsection{Equation of state for 2-flavour quark matter}  A nonlocal chiral quark model for 2-flavour $\\{u,d\\}$ and three color $\\{r,b,g\\}$ superconducting (2SC) quark matter in the mean field approximation is used, for details see \\cite{Blaschke:2003yn,Grigorian:2003vi}. The order parameters are the mass gap $\\phi_f$ and the diquark gap $\\Delta$ for the chiral and superconducting phase transitions respectively. As in \\cite{Grigorian:2003vi}, the following chemical potentials are introduced: $\\mu_q =  (\\mu_u+\\mu_d)/2$  for  quark number, $\\mu_I = (\\mu_u-\\mu_d)/2$ for isospin asymmetry and  $\\mu_8$ for color charge asymmetry. The deviation in the color space is considered $\\mu_8 \\ll \\mu_q$, so the effect of considering $\\mu_8$ is neglected.  The quark thermodynamic potential is expresed as \\cite{Kiriyama:2001ud} \\begin{eqnarray} \\lefteqn{\\Omega_q(\\phi,\\Delta;\\mu_q,{\\mu_I},T)+\\Omega_{vac} = \\frac{\\phi^2}{4G_1}+ \\frac{\\Delta^2}{4G_2}} \\nonumber\\\\ & & - \\frac{2}{2\\pi^2}\\int^\\infty_0dqq^2(N_c-2)\\{2E_{\\phi} +\\omega~[E_{\\phi}-\\mu_q-{\\mu_I},T] \\nonumber\\\\ & & +\\omega~[E_{\\phi}-\\mu_q+{\\mu_I},T] +\\omega~[E_{\\phi}+\\mu_q-{\\mu_I},T] +\\omega~[E_{\\phi}+\\mu_q+{\\mu_I},T] \\}\\nonumber\\\\ & & -\\frac{4}{2\\pi^2}\\int^\\infty_0dqq^2\\{E_{+}+E_{-} +\\omega[E_{\\phi}^{-}-{\\mu_I},T] \\nonumber\\\\ & & +\\omega[E_{\\phi}^{-}+{\\mu_I},T] +\\omega[E_{\\phi}^{+}-{\\mu_I},T] +\\omega[E_{\\phi}^{+}+{\\mu_I},T] \\} \\end{eqnarray} with \\begin{eqnarray} \\omega[E,T] = T\\ln\\left[1+\\exp\\left(-E/T\\right)\\right]~. \\end{eqnarray}  The dispersion relations for the quarks of unpaired  and paired colors are respectively, \\begin{eqnarray} {E_\\phi}^2&=& q^2+{(m+F^2(q)\\phi)}^2\\\\ {E_\\phi^\\pm}^2&=&(E_{\\phi}\\pm\\mu)^2+F^4(\\bar q){\\Delta}^2 \\end{eqnarray}  The interaction between the quarks is implemented via a Gaussian formfactor function $F(q)$ in the momentum space (Gaussian types give stable hybrid configurations \\cite{Blaschke:2003rg}) as  $F(q)=\\exp(-q^2/\\Lambda^2)~$.  The parameters $\\Lambda = 1.025$ GeV, $G_1= 3.761~\\Lambda^2$  and $m_u=m_d=m=2.41$ MeV are fixed by the pion mass, pion decay constant and the constituent quark mass at $T=\\mu=0$ \\cite{Schmidt:1994di}. The constant $G_2$ is a free parameter of the approach and is fixed as $G_2=0.86~ G_1$.   \\subsubsection{Stellar matter conditions}  The stellar matter in the quark core  of compact stars is considered to consists of $u$ and  $d$ quarks and leptons (electrons $e^-$  and antineutrinos $\\bar \\nu_e$) under the following conditions  \\begin{itemize} \\item $\\beta$-equilibrium $\\quad\\quad d \\longleftrightarrow u+e^-+\\bar \\nu_e,$ $\\quad\\quad \\mu_e +\\mu_{\\bar \\nu_e} = -2\\mu_I,$ \\item Charge neutrality $\\quad\\quad \\frac{2}{3}n_u-\\frac{1}{3}n_d-n_e = 0,$ $\\quad\\quad n_B+n_I-2n_e = 0,$ \\item Color neutrality $\\quad\\quad n_8=0,$ $\\quad\\quad 2n_{qr}-n_{qb}=0,$ \\end{itemize} where $n_j= \\frac{\\partial \\Omega}{\\partial \\mu_j}\\bigg|_{T,\\phi=\\phi_0,\\Delta=\\Delta_0}$ are the number densities corresponding to the chemical potential $\\mu_j$ defined above.  The lepton contributions ($l=e,\\bar \\nu_e$) as ideal Fermi gases \\begin{eqnarray} \\Omega_l(\\mu,T)= -\\frac{1}{12\\pi^2}\\mu^4 -\\frac{1}{6}\\mu^2T^2 -\\frac{7}{180}\\pi^2T^4 \\end{eqnarray} are added to the quark thermodynamical potential \\begin{eqnarray} \\Omega(\\phi,\\Delta;\\mu_q,\\mu_I,\\mu_e,T) = \\Omega_q(\\phi,\\Delta;\\mu_q,\\mu_I,T)+\\Omega_e(\\mu_e,T)+\\Omega_{\\bar \\nu_e}(\\mu_{\\bar \\nu_e},T). \\end{eqnarray}  The baryon chemical potential $\\mu_B = 3\\mu_q-\\mu_I$ is introduced as the conjugate of the baryon number density  $n_B$.   The $\\Omega$  function can have several minima in the $\\phi$, $\\Delta$ plane, an example is shown is Fig. \\ref{Ome_cuts}. The global minimum represents the  stable equilibrium of the system and the minima search is perfomed solving the gap equations  \\begin{eqnarray} {\\partial \\Omega \\over \\partial \\phi}\\bigg|_{\\phi=\\phi_0;\\Delta=\\Delta_0}={\\partial \\Omega \\over \\partial \\Delta}\\bigg|_{\\phi=\\phi_0;\\Delta=\\Delta_0}=0 \\end{eqnarray}  under the conditions that are mentioned above for the stellar interior.  \\begin{figure}[h] \\centerline{ \\psfig{figure=Ome_cuts.ps,height=9cm,width=13cm,angle=-90}} \\caption{Cuts of the thermodynamic potential $\\Omega(\\phi,\\Delta;\\mu_B,\\mu_I,T=0)$ in the planes $\\Delta = const$ (on the left) and $\\phi = const$ (on the right) for two different constant values of $\\mu_B$ and the corresponding $\\mu_I$. For $\\mu_B = 933$ MeV (upper panel) two degenerate minima can coexist at the values: $\\phi=331$ MeV, $\\Delta=0$ (solid lines) and $\\phi=107$ MeV, $\\Delta=98$ MeV (dashed lines). For $\\mu_B = 1100$ MeV (lower panel) the minimum with a nonvanishing diquark $\\Delta=121$ MeV and $\\phi=54.8$ MeV (dashed lines) is preferable. This corresponds to a first order transition from the vacuum to a superconducting phase. In this example $G_2/G_1=1$ was taken.} \\label{Ome_cuts} \\end{figure}  The thermodynamics of the system, e.g. pressure $P$, energy density $\\epsilon$, number density $n$ and entropy density $s$, is defined via this global minimum   \\begin{eqnarray} \\Omega(\\phi_{0},\\Delta_{0};\\mu_B,\\mu_I, T) = \\epsilon -Ts -\\mu_B n_B -\\mu_I n_I  = -P~. \\end{eqnarray}  To fulfill the charge neutrality condition (see Fig. \\ref{fig:volfraction}, right) a mixed phase  between a subphase with diquark condensation (subscript $\\Delta>0$) and normal quark matter subphase (subscript $\\Delta=0$) is defined via the Glendenning construction. The Gibbs condition for equilibrium at fixed $T$ and $\\mu_B$ is that the pressure of the subphases should be the same  \\begin{eqnarray} P =P^{\\Delta>0}(\\mu_B,\\mu_I,\\mu_e,T)=P^{\\Delta=0}(\\mu_B,\\mu_I,\\mu_e,T)~. \\end{eqnarray}  \\begin{figure}[h] \\vspace{-0.5cm} \\centerline{ \\psfig{figure=analisis3b.ps,height=7cm,width=7cm,angle=-90} \\psfig{figure=chi_neu2.ps,height=7cm,width=7cm,angle=-90}} \\caption{{\\bf Left}: Solutions of the gap equations and the charge neutrality condition (solid black line) in the $\\mu_I$ vs. $\\mu_e$ plane. Two branches are shown: states with diquark condensation on the upper right ($\\Delta>0$) and states from normal quark matter ($\\Delta=0$) on the lower left. The plateau in between corresponds to a mixed phase. The lines for the $\\beta$-equilibrium condition are also shown (solid and dashed red lines) for different values of the (anti)-neutrino chemical potential. The stellar matter should satisfy both conditions (intersection of the corresponding lines) and therefore for $\\mu_{\\bar \\nu_e}=0$ a mixed phase is preferable. {\\bf Right}:Volume fraction $\\chi$ of the phase with nonvanishing diquark condensate obtained by a Glendenning construction of a charge-neutral mixed phase. Results are shown for two different values of $\\mu_{\\bar \\nu_e}$.} \\label{fig:volfraction} \\end{figure}  The volume fraction $\\chi$ that is occupied by the subphase with diquark condensation is defined by the charge $Q$ in the subphases, \\begin{eqnarray} \\chi = Q_{\\Delta>0}/(Q_{\\Delta>0}-Q_{\\Delta=0})~, \\end{eqnarray}  and is plotted on the right panel in Fig. \\ref{fig:volfraction} for different antineutrino chemical potentials as a function of  $\\mu_B$. In the same way, the  number densities for the different particle species $j$ and the energy density are given by \\begin{eqnarray} n_j = \\chi n_{j_{\\Delta>0}}+ (1-\\chi)n_{j_{\\Delta=0}}~,\\\\ \\epsilon = \\chi \\epsilon_{\\Delta>0}+ (1-\\chi)\\epsilon_{\\Delta=0}~. \\end{eqnarray}  \\subsubsection{Equation of state with trapped antineutrinos}  Increasing the antineutrino chemical potential $\\mu_{\\bar \\nu_e}$ increases the asymmetry in the system and this shifts the onset of the superconducting phase transition to higher densities.  Above a critical value of $\\mu_{\\bar \\nu_e}\\geq \\mu_{\\bar \\nu_e}^c \\simeq 30$ MeV (see critical value in Fig. \\ref{fig:volfraction}, on the left)  first a normal quark matter phase occurs and  then the phase transition to superconducting matter takes place, see Fig. \\ref{fig:GE}, on the left.  The consequences for the equation of state can be seen on the right of the Fig. \\ref{fig:GE}: the onset of the superconductivity in quark matter is shifted to higher densities and the equation of state becames harder.  \\begin{figure}[h] \\vspace{-0.5cm} \\centerline{ \\psfig{figure=GE_NEU_0862.ps,height=7cm,width=7cm,angle=-90} \\psfig{figure=EoS_NEU_0862.ps,height=7cm,width=7cm,angle=-90}} \\caption{{\\bf Left}: Solutions of the gap equations and $\\mu_I$  as a function of $\\mu_B$. Increasing the antineutrino chemical potential increases the asymmetry in the system and the superconducting phase is inhibited at moderates densities. {\\bf Right}: Equation of state for different values of $\\mu_{\\bar \\nu_e}$ of trapped antineutrinos. As $\\mu_{\\bar \\nu_e}$ increases the equation of state becomes harder. } \\label{fig:GE} \\end{figure}  \\subsection{Quark stars and antineutrino trapping}  The configurations for the quark stars are obtained by solving the Tolman-Oppenheimer-Volkoff equations for a set of central quark number densities $n_q$ for which the stars are stable.  In Fig. \\ref{fig:Conf} the configurations for different antineutrino chemical potentials are  shown. The equations of state with trapped antineutrinos are softer and therefore this  allows more compact configurations. The presence of antineutrinos tends to increase the mass of the star for a given central density.  A reference configuration with total baryon number  $N_B = 1.51~ N_{\\odot}$ (where $N_{\\odot}$ is the total baryon number of the sun) is chosen and the case with (configurations $A$ and $B$ in Fig. \\ref{fig:Conf}) and without antineutrinos ($f$ in  Fig. \\ref{fig:Conf}) are compared. A mass defect can be calculated between the configurations with  and without trapped antineutrinos at constant total baryon number and the result is shown on the right panel of  Fig. \\ref{fig:Conf}). The mass defect could be interpreted as an energy release if the configurations $A,B$ with antineutrinos are initial states and the configuration $f$ without them is the final state of a protoneutron star evolution.  \\begin{figure}[h] \\vspace{-0.5cm} \\centerline{ \\psfig{figure=smzoom_0862.ps,height=7cm,width=7cm,angle=-90} \\psfig{figure=DeltaM_0862.ps,height=7cm,width=7cm,angle=-90}} \\caption{{\\bf Left}: Quark star configurations for different antineutrino chemical potentials $\\mu_{\\bar \\nu_e}=0, ~100,~150$ MeV. The total mass $M$ in solar masses ($M_{{\\rm sun}}\\equiv M_\\odot$ in the text) is shown as a function of the radius $R$ (left panel) and of the central number density  $n_q$ in units of the nuclear saturation density $n_0$ (right panel). Asterisks denote two different sets of configurations (A,B,f) and (A',B',f') with a fixed total baryon number of the set. {\\bf Right}: Mass defect $\\Delta M$ and corresponding energy release $\\Delta E$ due to antineutrino untrapping as a function of the mass of the final state $M_f$. The shaded region is defined by the estimates for the upper and lower limits of the antineutrino chemical potential in the initial state $\\mu_{\\bar \\nu_e}=150$ MeV (dashed-dotted line) and $\\mu_{\\bar \\nu_e}=100$ MeV (dashed line), respectively. } \\label{fig:Conf} \\end{figure}   \\subsubsection{Protoneutron star evolution with antineutrino trapping}  After the collapse of a protoneutron star the star cools down by surface emission of photons and antineutrinos. Antineutrinos are trapped because they were generated by the direct $\\beta$-process in the hot and dense matter and could not escape due to their small mean free path. The region of the star where the temperature falls below the density dependent critical value for diquark condensation, will transform to the color superconducting state which is almost transparent to (anti)neutrinos. But nevertheless due to the trapped antineutrinos there is a dilute normal quark matter shell which prevents neutrino escape from the superconducting bulk of the star, see Fig. \\ref{fig:kugel} and Fig. \\ref{fig:phasediag}. The criterion for the  neutrino untrapping transition is to cool the star below a temperature of about 1 MeV when the mean free path of neutrinos becomes larger than the shell radius \\cite{Prakash:2001rx}. If at this temperature the antineutrino chemical potential is still large then the neutrinos can escape in a sudden outburst. If it is small then there will be only a gradual increase in the luminosity. An estimate for the possible release of energy within the outburst scenario can be given via the mass defect  defined in the previous subsection between an initial configuration with trapped neutrinos (state $A$ or $B$) and a final configuration without neutrinos (state $f$).  \\begin{figure}[h] \\vspace{-0.5cm} \\centerline{ \\parbox{6cm}{\\psfig{figure=esf1.epsi,width=2.5cm,angle=-90}} \\parbox{6cm}{\\psfig{figure=esf2.epsi,width=2.5cm,angle=-90}} \\parbox{6cm}{\\psfig{figure=esf3.epsi,width=2.5cm,angle=-90}}} \\caption{Left graph: Quark star cooling by antineutrino and photon emission from the surface. Middle graph: Two-phase structure developes due to the trapped antineutrinos: a normal quark matter shell and a superconducting interior. Right graph: Antineutrino untrapping and burst-type release of energy.} \\label{fig:kugel} \\end{figure}  \\begin{figure}[h] \\vspace{-0.5cm} \\centerline{ \\psfig{figure=PDN2.ps,width=10cm,angle=-90} } \\caption{Star evolution corresponding to Fig. \\ref{fig:kugel} plotted in the phase diagram} \\label{fig:phasediag} \\end{figure} ",
        "conclusions": "The effects of trapped antineutrinos on the diquark condensates in quark star configurations are investigated. At fixed baryon number the energy release in the antineutrino untrapping transition is of the order of $10^{52}$ erg. This is a first order transition and leads to an explosive release of energy that could help to explain energetic phenomena in the universe like gamma ray bursts or supernova explosions."
    },
    "0402/astro-ph0402559_arXiv.txt": {
        "abstract": "{Orbital periods and other parameters of symbiotic binary systems in the LMC and SMC are presented and discussed. In particular, the symbiotic stars in the MCs are compared with those in the Milky Way.}    \\addkeyword{Magellanic Clouds} \\addkeyword{Stars: symbiotic} \\addkeyword{Stars: binary systems} \\addkeyword{Stars: parameters}   \\begin{document} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0402/hep-ph0402220_arXiv.txt": {
        "abstract": "Recently, the SPI spectrometer on the INTEGRAL satellite observed strong 511 keV line emission from the galactic bulge. Although the angular distribution (spherically symmetric with width of $\\sim 9^{\\circ}$) of this emission is difficult to account for with traditional astrophysical scenarios, light dark matter particles could account for the observation. In this letter, we consider the possibility that decaying axinos in an R-parity violating model of supersymmetry may be the source of this emission. We find that $\\sim 1-300\\, \\rm{MeV}$ axinos with R-parity violating couplings can naturally produce the observed emission. ",
        "introduction": "The SPI spectrometer on the INTEGRAL (INTErnational Gamma-Ray Astrophysics Laboratory) satellite has made the observation of a bright ($9.9^{+4.7}_{-2.1} \\times 10^{-4}\\,\\rm{ph}\\,\\rm{cm}^{-2}\\,\\rm{s}^{-1}$), 511 keV gamma-ray emission line from the galactic bulge \\cite{511}. The emission is consistent with being spherically symmetric, with a full-width-half-maximum of about $9^{\\circ}$($6^{\\circ}-18^{\\circ}$ at 2$\\sigma$ confidence). The 3 keV width of the line is dominated by $e+ e^-$ annihilations via positronium formation \\cite{milne}. These findings are in agreement with earlier observations, such as those by the OSSE experiment \\cite{previous}.  This observation of bright 511 keV emission from the galactic bulge has been quite difficult to explain with traditional astrophysics. Most potential sources considered do not produce a sufficient number of positrons (such as neutron stars, black holes \\cite{compact}, radioactive nuclei from supernovae, novae, red giants or Wolf-Rayet stars \\cite{stars}, cosmic ray interactions with the interstellar medium \\cite{ism}, pulsars \\cite{pulsars} or stellar flares) and those which may possibly be capable of producing the required number (type Ia supernovae \\cite{debate,casse} or hypernovae \\cite{sn2003,casse}), may not be capable of filling the entire galactic bulge \\cite{starbursts}. In particular, the rate at high altitude is likely to be too low to explain the observed extension of the 511 keV source \\cite{pohl}.  Given this difficulty, alternative explanations should be considered. In Ref.~\\cite{boehm}, it was suggested that light (1-100 MeV) dark matter particles annihilating in the galactic bulge could explain the observed emission. In this letter we, instead, consider light {\\it decaying} dark matter particles. In particular, we consider the supersymmetric partner of the axion\\cite{axion}, the axino, in R-parity violating supersymmetric models.   In the Minimal Supersymmetric Standard Model (MSSM), a $Z_2$ symmetry, R-parity\\cite{rparity}, is usually imposed to forbid dimension four operators which lead to fast proton decay \\cite{protondecay}. A by-product of an exact R-parity is that the Lightest Supersymmetric Particle (LSP) is stable and often a good candidate for cold dark matter. Although R-parity is an elegant way of suppressing proton decay, it does not have to be the only way that nature can choose to do so. In particular, baryon parity is sufficient to make the proton stable. In this case, the coupling strengths of other R-parity violating operators, such as $\\lambda_{ijk}L_i L_j E^c_k$, are  much less constrained. For a summary of R-parity violating coupling constraints, see Ref.~\\cite{rpv}.   A typical dark matter candidate provided by R-parity conserving supersymmetry is a neutralino LSP. Neutralinos are the superpartners of the neutral gauge bosons and Higgs bosons and have masses constrained by direct searches to be greater than $\\sim 30$ GeV, too heavy to produce a large flux of thermal positrons. Of course, with sizable R-parity violation, the neutralino will have a short lifetime and cease to be a good candidate for cold dark matter.  With the Peccei-Quinn (PQ) mechanism \\cite{axion} providing a natural solution to the strong-CP problem in a supersymmetric theory, the existence of an axino is inevitable. The axino's mass is expected to be considerably lower than the electroweak scale, perhaps in the keV to several GeV range and is capable of providing the observed quantity of dark matter given a low reheating temperature \\cite{axinodark,kim01,axinodecay,covi99,covi01}. Since the axino is in the same supermultiplet as the axion, its couplings to matter fields are generically suppressed by $f_a^{-1}$, where $f_a$ is the PQ symmetry breaking scale $\\sim 10^9 - 10^{12}$ GeV. As we shall see in detail in this paper, due to this large suppression, axinos could be considered {\\it stable} during the history of the universe even in the presence of sizable R-parity violating couplings. Furthermore, a long-lived MeV-GeV axino, such as we consider in this letter, would be heavy enough to constitute a good candidate for cold dark matter \\cite{covi99,covi01,axinoreheat}.  A possible concern for axinos in the early universe is their effect on the light element abundances \\cite{lightelementscorrect,reheating,lightelements}. With a long-lived axino, however, axino decays do not threaten these observations. Furthermore, with the introduction of R-parity violation in our scenario, we allow for relatively fast Next-to-Lightest Supersymmetric Particle (NLSP) decays into Standard Model particles. Thus the epoch of SUSY particle decays is over prior to nucleosynthesis and the light element abundances remain unaffected \\cite{lightelementsrviolation}.  Before going into a more detailed study of decaying axinos, we briefly consider here the possibility of decaying gravitinos. Gravitinos also have highly suppressed couplings to matter ($M_P^{-1}$ in the kinematical regime of interest) and could be a good candidate for cold dark matter.  We have found, however, that for the case of trilinear R-parity violating terms, the gravitino lifetime is too long to account for the observed 511 keV emission (see section II). In particular, the gravitino lifetime is estimated to be $\\tau_{3/2} \\sim 10^{31}\\,\\mbox{sec}\\, \\left(m_{\\tilde{l}}/100\\,\\mbox{GeV} \\right)^4 \\, \\left(0.1\\,\\mbox{GeV}/m_{3/2}\\right)^7 \\, \\left(0.1/\\lambda\\right)^2$. ",
        "conclusions": "In this letter, we have demonstrated that a light axino (1-300 MeV), in either the KSVZ or DFSZ axion models, with trilinear R-parity violating couplings could be responsible for the 511 keV line emission observed from the galactic bulge. In this scenario, axinos constitute the major component of the cold dark matter and are present in the galactic halo with a cusped distribution ($\\gamma \\sim 1.2$).  At this time, we can not exclude the possibility that poorly understood conventional astrophysics is responsible for the observed positron production in the galactic bulge. To differentiate such a scenario from more exotic sources, such as light decaying particles, future tests must be made \\cite{dwarf}. Additionally, if decaying axinos are the source of the 511 keV emission, signatures of supersymmetry, and R-parity violation will likely be observed at the LHC.  As this letter was being completed, an article appeared which also discussed decaying particles as the source of the observed 511 keV emission \\cite{new511}. They considered decaying sterile neutrinos with rather constrained mixing parameters. They also discussed decaying scalars with gravitational strength interactions.  {\\it Acknowledgments}: DH is supported by the Leverhulme Trust. LW is supported in part by U.S. Department of Energy under grant DE-FG02-95ER40896. We would like to thank D. Chung and L. Roszkowski for useful discussions. \\vskip -0.5cm"
    },
    "0402/astro-ph0402009_arXiv.txt": {
        "abstract": "We present simulations of interferometric Sunyaev-Zel'dovich effect (SZE) and optical weak lenisng observations for the forthcoming AMiBA experiment, aiming at searching for high-redshift clusters of galaxies. On the basis of simulated sky maps, we have derived theoretical halo number counts and redshift distributions of selected halo samples for an AMiBA SZE survey and a weak lensing follow-up  survey. By utilizing the conditional number counts of weak lensing halos with the faint SZE detection, we show that a combined SZE and weak lensing survey can gain an additional fainter halo sample at a given false positive rate, which cannot be obtained from either survey alone. ",
        "introduction": "The thermal Sunyaev-Zel'dovich effect (SZE\\cite{Birkinshaw}) is a spectral distortion of the Cosmic Microwave Background (CMB) radiation due to the inverse-Compton scattering of CMB photons by high energy electrons in the intracluster medium (ICM). The most remarkable properties of the SZE is that its surface brightness is redshift independent, which make it as an ideal probe of the high-redshift universe. Further, since the SZE is proportional to the thermal energy content of the ICM, SZE imaging surveys allow us to select clusters over a wide range of the redshift with physically meaningful selection criteria. Weak gravitational lensing, on the other hand, probes the total mass projected along the line-of-sight, and hence provides complementary information on the mass of galaxy clusters.\\cite{Bartelmann} Array for Microwave Background Anisotropy (AMiBA\\cite{AMiBA}) is a 19-element interferometric array with full polarization capabilities operating at $95$GHz, specifically designed for the CMB observations. One of the main scientific goals of AMiBA is to conduct blind SZE surveys to search for high-redshift clusters. AMiBA will also conduct follow-up optical imaging observations with wide-field camera, MegaCam, at {\\it Canada France Hawaii Telescope} (CFHT). In this paper, we simulate the forthcoming AMiBA SZE experiment combined with the planned follow-up weak lensing observations to examine the expected cluster number counts for individual surveys, and explore the potential of a combined AMiBA SZE and weak lensing cluser survey. ",
        "conclusions": "We have examined the expected cluster number counts and redshift distributions of cluster samples for the forthcoming AMiBA SZE/weak lensing cluster survey on the basis of $\\Lambda$CDM cosmological simulations. By utilizing the conditional halo number counts $N(\\kappa>\\kappa_{\\rm lim} |S>S_{\\rm lim})$, we have demonstrated that a combined SZE and weak lensing survey can % gain an additional fainter halo sample at a given false positive rate, which cannot be obtained from either survey alone."
    },
    "0402/astro-ph0402523_arXiv.txt": {
        "abstract": "\\noindent By combining the 2--degree Field Galaxy Redshift Survey with the NRAO VLA Sky Survey at 1.4\\,GHz, the environments of radio loud AGN in the nearby Universe are investigated using both local projected galaxy densities and a friends--of--friends group finding algorithm. Radio--loud AGN are preferentially located in galaxy groups and poor--to--moderate richness galaxy clusters. The AGN fraction appears to depend more strongly on the large--scale environment (group, cluster, etc) in which a galaxy is located than on its more local environment, except at the lowest galaxy surface densities where practically no radio--loud AGN are found. The ratio of absorption--line to emission--line AGN changes dramatically with environment, with essentially all radio--loud AGN in rich environments showing no emission lines. This result could be connected with the lack of cool gas in cluster galaxies, and may have important consequences for analyses of optically--selected AGN, which are invariably selected on emission line properties. The local galaxy surface density of the absorption--line AGN is strongly correlated with radio luminosity, implying that the radio luminosities may be significantly boosted in dense environments due to confinement by the hot intracluster gas.  The environments of a radio--selected sample of star forming galaxies are also investigated to provide an independent test of optical studies. In line with those studies, the fraction of star forming galaxies is found to decrease strongly with increasing local galaxy surface density; this correlation extends across the whole range of galaxy surface densities, with no evidence for the density threshold found in some optical studies. ",
        "introduction": "The advent of large galaxy redshift surveys, especially the 2-degree Field Galaxy Redshift Survey (2dFGRS; Colless \\etal\\ 2001)\\nocite{col01} and the Sloan Digital Sky Survey (SDSS; York \\etal\\ 2000; Stoughton \\etal\\ 2002)\\nocite{yor00,sto02} have revolutionised our understanding of the effect of local environment upon the evolution of galaxies. Understanding how galaxy properties, such as luminosities, morphologies, star formation rates and nuclear activity, depend upon the environment that a galaxy inhabits can place important constraints on models of galaxy formation and evolution, and allow the intrinsic properties of the galaxies to be separated from those that have been externally induced (`nature {\\it vs} nurture').  It has been known for many years that star formation rates are strongly suppressed in the central regions of galaxy clusters (e.g. Dressler et~al 1985),\\nocite{dre85} even if account is taken of the different distribution of morphological types in cluster environments as compared with the field. The large redshift surveys have shown that this suppression of the star formation rate is not only restricted to the extreme cluster environments, but begins at much lower environmental densities.  Hashimoto \\etal\\ \\shortcite{has98} showed that the mean star formation rate shows a continuous correlation with local galaxy density, both inside and outside of clusters, using the Las Campanas Redshift Survey. Lewis et~al (2002; hereafter Lew02)\\nocite{Lew02} studied the fields around 17 clusters within the 2dFGRS, and found that the mean star formation rate of galaxies is relatively constant for projected galaxy surface densities below 1 galaxy per square Mpc, but that at higher surface densities star formation is increasingly suppressed, down to essentially zero at 50 galaxies per square Mpc. A similar study using SDSS data broadly supports these conclusions although with a weaker break (G{\\'omez \\etal\\ 2003; hereafter Gom03)\\nocite{Gom03}, and a study of SDSS field galaxies (Mateus \\& S{\\'o}dre 2003; hereafter MS03)\\nocite{mat03} is also in agreement at high surface densities, but suggests that the star formation rate remains sensitive to the local galaxy density even in more rarefied environments.  Regardless of the precise dependence at lower surface densities, environmental effects clearly function down to at least 1 galaxy per square Mpc, which is below both the mean galaxy surface density at the virial radius of relatively rich clusters and that of galaxy groups. Therefore, the physical processes that lead to this quenching of star formation are not intrinsic to cluster environments (e.g. ram--pressure stripping of the interstellar medium by the hot cluster gas; Gunn \\& Gott 1972),\\nocite{gun72} but also occur in smaller structures (cf. Mart{\\'\\i}nez \\etal\\ 2002, who find that star formation is diminished even within galaxy groups of mass $M \\sim 10^{13}$ in the 2dFGRS).\\nocite{mar02b} The combination of this environmental dependence and the build up of galaxies into groups and clusters with cosmic time may be one of the drivers behind the decline in the mean cosmic star formation rate since redshifts $z \\sim 1$ \\cite{mad98}.  An alternative handle on galaxy activity and environmental influence comes through studies of active galactic nuclei (AGN). It is now apparent that essentially all massive galaxies in the nearby Universe host a supermassive black hole at their centres, whose mass is roughly proportional to the spheroidal mass of the galaxy (e.g. see review by Kormendy \\& Gebhardt 2001)\\nocite{kor01}. This suggests that the build-up of the central black hole and that of its host galaxy are fundamentally linked; the similarity of the cosmic evolution of the mean global star formation rate to that of the rate at which gas is accreted onto black holes (as estimated from the radio luminosity function; Dunlop 1998),\\nocite{dun98} at least out to redshifts $z \\sim 1$-2 where both are well-determined, provides further evidence for this. Investigating how the incidence of AGN activity depends upon local environment can provide valuable insight into the origin of this connection, and also into the physical processes that trigger AGN activity.  To produce a powerful AGN the necessary ingredients are a massive black hole, a supply of fuel, and a transport mechanism to connect the two. If AGN activity is driven predominantly by the availability of the cold gas (the same factor that drives star formation activity) then AGN activity, like star formation, should be greatly suppressed in cluster environments. If instead the probability of AGN activity is solely a function of the central black hole, independent of the local environment of the galaxy, the AGN fraction would simply trace the distribution of galaxy bulges. Alternatively, at least for the most powerful AGN it has frequently been proposed that galaxy interactions or mergers may trigger the AGN activity (e.g. Hutchings \\& Campbell 1983, Heckman \\etal\\ 1984, 1986),\\nocite{hut83,hec84,hec86} providing both a supply of gas and a mechanism for moving it to the central regions of the galaxy. In this case the dependence of AGN activity on environment on group-- or cluster--scales would be less clear--cut, with AGN favouring those environments in which conditions are optimal for galaxy interactions and mergers.  Studies of the environmental dependence of AGN activity have a long and chequered history. Dressler \\etal\\ \\shortcite{dre85} argued that AGN activity was suppressed in clusters, finding that only 1\\% of cluster galaxies showed AGN activity compared to 5\\% of field galaxies in their sample.\\footnote{Note that the fraction of galaxies which host AGN is very dependent upon both the depth of the observations and the details of the spectral definition of an AGN: Ho \\etal\\ \\shortcite{ho97} classify as many as 43\\% of the galaxies in their survey of nearby bright ($B_T < 12.5$) galaxies as active.} However, the redshift range of their cluster and field samples were very different, and the field sample contained many higher redshift AGN which were only bright enough to make it into their sample because of the magnitude boosting effect of the AGN itself. Biviano \\etal\\ \\shortcite{biv97} find that if a correction is applied for this magnitude bias then the lower incidence of AGN activity in clusters is consistent with being due solely to the difference in the morphological mix between cluster and field. Carter \\etal\\ \\shortcite{car01} similarly find little evidence for an environmental dependence of AGN fraction.  Using the 2dFGRS and the SDSS it is possible to go beyond a binary comparison of cluster against field and investigate the environmental dependence of AGN activity in much greater detail, as has been done for star formation activity. Kauffmann et~al \\shortcite{kau03} robustly selected a sample of AGN from the SDSS survey, by modelling the underlying stellar continuum to obtain accurate measurements of the Balmer emission lines, and then using emission line ratio diagnostics to separate the AGN from star forming galaxies. Kauffmann et al (in prep; private communication) find that these AGN, especially those with strong emission lines, are less common in regions of high galaxy surface density. However, Miller et~al \\shortcite{mil03a}, also using the SDSS, found that the fraction of galaxies hosting AGN remains roughly constant across all environments from cluster cores to the rarefied field.  The contradiction between these two sets of results indicates that selection of active galaxies by optical emission lines is neither straightforward nor uncontroversial. In addition, recent Chandra X-ray observations of galaxy clusters have identified a population of X-ray active cluster galaxies which would not be selected as active galaxies on the basis of either their optical or emission line properties (e.g. Martini \\etal\\ 2002)\\nocite{mar02a}, and radio studies \\cite{mil02a} have similarly found a population of dust--obscured active and star--forming galaxies towards the central regions of clusters. Therefore, selecting active galaxies by means other than their emission line properties may provide more robust results, and will certainly provide a valuable test of the results currently being derived from the SDSS. An efficient way to do this is to use radio-loud AGN: these are straightforward to locate and study, and large--area deep radio surveys are already available. It must be borne in mind, however, that only 10-20\\% of AGN are radio loud (e.g. Hooper \\etal\\ 1995)\\nocite{hoo95}, and these may represent a biased subset of the AGN population as a whole.  The environments of powerful radio sources have been widely studied at low redshifts (e.g. Prestage \\& Peacock 1988; Hill \\& Lilly 1991; Miller \\etal\\ 2002)\\nocite{pre88,hil91,mil02b} and appear to favour galaxy groups and weak clusters; they tend to avoid the densest environments except in the few cases where they are in the special location of being hosted by the central dominant galaxy of a cluster. To date, though, these studies have been both limited to the most powerful radio sources and, generally, based upon galaxy number count statistics or cross--correlation analyses, with little or no spectroscopic redshift information (the Miller \\etal\\ study is the exception to this, providing redshifts for a handful of galaxies in the $\\sim 20$ arcmin radius region around each of 25 low redshift radio sources). By comparing the 2dFGRS with a deep large--area radio survey such as the NRAO VLA Sky Survey (NVSS; Condon \\etal\\ 1998)\\nocite{con98} it is possible to overcome both of these deficiencies, studying the spectroscopically--determined large--scale environments of `typical' radio--loud AGN within the 2dFGRS, and comparing these to those of the galaxy population in general. This is the goal of the current paper.  In Section~\\ref{samples}, the galaxy and radio source samples used in the analysis are defined. The methods used to derive the properties of these galaxies are described in Section~\\ref{methods}. The results of the environmental analysis are described in Section~\\ref{results}, and these are discussed in Section~\\ref{discuss}. Section~\\ref{concs} summarises the results. Throughout the paper, the values adopted for the cosmological parameters are $\\Omega_m = 0.3$, $\\Omega_{\\Lambda} = 0.7$, and $H_0 = 65$\\,km\\,s$^{-1}$Mpc$^{-1}$. ",
        "conclusions": "\\label{concs}  The results of this paper can be summarised as follows:  \\begin{itemize} \\item The proportion of radio--selected star forming galaxies decreases strongly with increasing local galaxy surface density, in the same manner as found in optical studies of star forming galaxies. This correlation extends over the full range of galaxy surface densities, with no evidence for a lower density threshold.  \\item Radio--loud AGN activity shows little dependence on local galaxy surface density, except at the very lowest surface densities where little AGN activity is found. The larger scale environment is more important in determining AGN activity: AGN are preferentially found in moderate groups and poor clusters.  \\item The AGN activity traces neither the distribution of galaxy bulges nor the availability of cold gas in galaxies, meaning that an external influence is required to trigger the activity. The higher AGN fraction in environments where conditions are optimised for galaxy interactions and mergers indicates that these are likely to be an important mechanism.  \\item Where AGN are found in poor or moderate richness clusters they are almost invariably absorption--line AGN, and have relatively high radio luminosities. This likely reflects the lack of cool gas typically available for ionisation in cluster environments, and suggests that the radio luminosity of these sources may have been boosted by their dense surrounding environment.  \\item The substantial drop in the ratio of emission--line to absorption--line AGN in dense environments implies that, at very least, considerable care must be taken in selecting samples of AGN from their optical emission--line properties. Potentially these samples could miss a large fraction of cluster AGN, in which case results from AGN environmental studies based upon optically--selected AGN samples would be unreliable. \\end{itemize}  An investigation of the environments of radio--selected AGN over a much larger area, such as that which will ultimately be possible using the SDSS, will permit these environmental variations to be studied to a much greater degree, using a radio sample of sufficient size for more detailed statistical investigation of AGN subsamples. In addition, such studies would enable a detailed comparison between optically and radio selected AGN samples; this is the key to understanding the potential biases of each method.  It is also important to investigate more powerful AGN than those studied here: the most powerful radio source in the current study has a 1.4\\,GHz radio luminosity of $1.4 \\times 10^{25}$W\\,Hz$^{-1}$, while the most powerful nearby sources have radio luminosities of $10^{26-27}$W\\,Hz$^{-1}$, much more comparable to those found at higher redshifts. Studies of these sources would permit an investigation of the cosmic evolution of AGN environments, but such powerful AGN are very rare in the nearby Universe, and a dedicated redshift survey of their environments will be required to achieve this."
    },
    "0402/astro-ph0402379_arXiv.txt": {
        "abstract": "We study the simplest class of Bekenstein-type, varying $\\alpha$ models, in which the two available free functions (potential and gauge kinetic function) are Taylor-expanded up to linear order. Any realistic model of this type reduces to a model in this class for a certain time interval around the present day. Nevertheless, we show that no such model is consistent with all existing observational results. We discuss possible implications of these findings, and in particular clarify the ambiguous statement (often found in the literature) that ``the Webb results are inconsistent with Oklo''. ",
        "introduction": "In theories with additional spacetime dimensions \\cite{Polchinski} there are typically many light or massless degrees of freedom, which can give rise to a number of observable cosmological consequences. Noteworthy among these are  variations of the fundamental couplings \\cite{Essay,Uzan,VFC} (with the ensuing violations of the Equivalence Principle \\cite{Will,Damour2}) and contributions to the energy density budget of the Universe.  In recent years there has been a growing body of evidence for the presence of these two effects. Type Ia supernovae \\cite{Tonry}, the Cosmic Microwave Background (CMB) \\cite{Bennett} and lensing data \\cite{Bernardeau} are all consistent with the existence of the so-called dark energy component, whose gravitational behaviour is very similar to that of a cosmological constant, and which indeed appears to have become the dominant component in the energy budget of the Universe at a redshift $z\\sim1$.  On the other hand there is some (somewhat more controversial) evidence for the spacetime variation of the fine-structure constant $\\alpha$, coming from both quasar absorption systems (at redshifts $z\\sim1-3$, \\cite{Webb,Murphy,Petit}) and the Oklo natural nuclear reactor ($z\\sim0.1$, \\cite{Fujii}). There is also a further claim of a varying proton to electron mass ratio, also at $z\\sim3$ \\cite{Ivanchik}. While it is conceivable that hidden systematic effects are still contaminating some of these measurements, an unprecedented effort is being made by a number of independent groups and using a range of techniques, to search for such variations at various key cosmological epochs, which should soon clarify the situation. There is also a range of other constraints, either local \\cite{Marion} (from atomic clocks) or at low \\cite{Olive} (from Rhenium decay in meteorites) or high redshift \\cite{Avelino,Martins,Rocha} (from the CMB and BBN), with much stringent ones forthcoming \\cite{wmap1,wmap2}.  It goes without saying that from the point of view of a fundamental theory there is more than ample freedom to allow the dark energy and the varying couplings to be due to different degrees of freedom in the theory, and even to emerge through different physical mechanisms. Nevertheless, it is useful to study the simplest case in which the two have a common origin, as this will in principle have the fewest free parameters and can therefore be better constrained.  In what follows we will discuss a very simple toy model for this which, although arguably oversimplified from the particle physics point of view, has the advantage of having a minimal number of free parameters. The basic idea is to consider Bekenstein-type models \\cite{Bekenstein}, and reduce the freedom in the two free functions (the potential $V(\\phi)$ and the gauge kinetic function $B_F(\\phi)$) by Taylor-expanding them around the present day, and retaining only terms up to linear order. In fact, we will see that its free parameters are in some sense too few, so that the model is very tightly constrained by existing observations, and indeed ruled out if all of them are correct. Still the model can be useful in providing some guidance for the likely requirements of successful, fundamental theory inspired models. Other interesting analyses of this class of models can be found in \\cite{Sandvik,OlivePos,Anchordoqui,Parkinson,Copeland,Nunes}.  We will start in Sect. II with a brief overview of the Bekensein-type models. We then discuss in more detail the linearized regime that interests us (Sect. III) and discuss it in the context of existing observational data (Sect. IV). Finally Sect. V summarizes our results and briefly discusses further prospects. Throughout this paper we shall use fundamental units with $\\hbar=c=G=1$. ",
        "conclusions": "We have studied the simplest class of Bekenstein-type, varying $\\alpha$ models, and compared them to existing observational constraints. These are models in which the two available free functions (the potential and the gauge kinetic function) are Taylor-expanded around present-day values, with terms kept only up to linear order. Despite their apparent simplicity, they are interesting to the extent that any realistic model of this type should reduce to a model in this class for a certain time interval around the present day. Nevertheless, their simplicity means that very specific predictions ensue, that can be compared with existing data. We have shown that no such model is consistent with all the existing observational results. Hence either some of these observations are dominated by unknown systematics or our linearity assumption breaks down on a timescale significantly smaller than a Hubble time.  Given that a scalar field that produces a varying fine-structure constant can also make a significant contribution towards the dark energy of the universe, it's interesting to speculate on the possible relation between the above observation and hints for a time-varying equation of state of dark energy.Indeed in the latter context it has been argued that something analogous seems to happen: observational  data seem to disfavour not only a constant equation of state, but even a mildly varying one, say with a linear dependence in redshift \\cite{Bassett,Beca,Alam}. It is unclear if the two things are somehow related, but it has been said that a coincidence is always worth noticing---one can always discard it later if it turns out to be just a coincidence."
    },
    "0402/astro-ph0402653_arXiv.txt": {
        "abstract": "We have performed a non-LTE spectroscopic analysis using far-UV and UV data of the central star of the planetary nebula K1-26 (Longmore 1), and found $\\Teff = 120\\pm10$~kK, $\\logg = 6.7^{+0.3}_{-0.7}$~\\gunit, and $y \\simeq 0.10$.  The temperature is significantly hotter than previous results based on optical line analyses, highlighting the importance of analyzing the spectra of such hot objects at shorter wavelengths.  The spectra show metal lines (from, \\eg, carbon, oxygen, sulfur, and iron).  The signatures of most elements can be fit adequately using solar abundances, confirming the classification of \\star\\ as a high gravity O(H) object.  Adopting a distance of 800~pc, we derive $\\Rstar \\simeq 0.04 $~\\Rsun, $L \\simeq 250 $~\\Lsun, and $M \\simeq 0.6$~\\Msun.  This places the object on the white dwarf cooling sequence of the evolutionary tracks with an age of $\\tau_{evol} \\simeq 65$~kyr. ",
        "introduction": "\\label{sec:intro}  Longmore 1 (K1-26, PK 255-59 1, hereafter \\star) was originally discovered by \\citet{longmore:77} as a PN having a notably large angular size ($\\sim400$\\arcsec).  The spectra of its central star show both hydrogen and \\HeII\\ absorption features, with no evidence of a stellar wind in its UV or optical spectra \\citep{patriarchi:91, kaler:85, mendez:85}.  Because of its high galactic latitude ($b \\simeq -\\degree{60}$), the reddening toward \\star\\ is thought to be minimal \\citep{kaler:85}.  Based on its optical spectrum, \\citet{mendez:85} termed \\star\\ an ``hgO(H)'' star - a high gravity object with very broad Balmer absorptions.  Such objects can lie on the white-dwarf cooling tracks, but can also be non-post-AGB objects.  A distance of $D=800$~pc \\citep{ishida:87} implies a nebular radius of $\\sim0.8$~pc, suggesting that \\star\\ is a quite evolved CSPN (most PN have radii $\\lesssim 0.5$~pc --- \\citealp{cahn:92}).  Hot central stars of planetary nebulae (CSPN) emit most of their observable flux in the Far-UV range.  We have observed the central star of \\star\\ with the \\emph{Far Ultraviolet Spectroscopic Explorer} (FUSE) satellite in the 905---1187~\\AA\\ range.  Using this data as well as archive \\emph{International Ultraviolet Explorer} (IUE) data (1150---3300~\\AA), we determined the parameters of the central star through stellar modeling, and discuss evolutionary implications.  Parameters of \\star\\ compiled from previous literature are listed in Table~\\ref{tab:lit_params}. ",
        "conclusions": "\\label{sec:conclusions}  We have analyzed FUV and UV spectra of \\star, a hot CSPN notable for its relatively high galactic latitude and thus having a minimal reddening. Its FUSE spectrum, aside from showing hydrogen and helium lines, shows strong \\OVI\\ \\doublet 1032,38 signatures, perhaps indicating an oxygen-enriched object.  We have modeled the FUSE and IUE spectrum of this object to determine parameters of $\\Teff = 120$~kK, $\\logg = 6.7$~\\gunit, $\\Rstar = 0.04$~\\Rsun, $L = 250$~\\Lsun, and $M \\simeq 0.6$~\\Msun.  The temperature is much higher than that derived by \\citet{mendez:85} from an optical-line analysis ($\\Teff = 65\\pm10$~kK), and illustrates the importance of the FUV-UV range in the analysis of hot CSPN.  These parameters confirm the \\citet{mendez:85} classification of \\star\\ as a high-gravity O(H) star.  Comparison of our parameters to evolutionary tracks indicate a post-AGB age of $\\sim 65$~kyr.  We also measure $\\vrad \\simeq 100$~\\kms\\ for the \\star\\ PN system."
    },
    "0402/astro-ph0402186_arXiv.txt": {
        "abstract": "We have obtained near-infrared spectra of L dwarfs, L/T transition objects and T dwarfs using Subaru. The resulting spectra are examined in detail to see their dependence on the spectral types. One question is where the methane feature appears: We suggest that it appears at L8 and marginally at L6.5. The water bands at 1.1 and 1.4 $\\mu$m do not necessarily show steady increase towards later L types but may show inversion at late L types. This does not necessarily imply that the spectral types do not represent a temperature sequence, but can be interpreted as due to compensation of the increasing water abundance by the heavier of dust extinction in the later L types. We confirm that the FeH 0.99 $\\mu$m bands appear not only in the late L dwarfs but also in the early T dwarfs. We suggest that FeH could be dredged up by the surface convective zone induced by the steep temperature gradient due to the large opacity of the dust cloud itself and will be replenished constantly by the convection.  We have obtained bolometric luminosities of the objects with known parallaxes in our sample, first by integrating the spectra between 0.87 and 2.5 $\\mu$m, and second  by the $K$-band bolometric correction. Apart from an L3 dwarf, the bolometric luminosities obtained by both methods agree well and this implies that the $K$-band bolometric correction, which is  obtained by the use of the Unified Cloudy Models, can be applied to obtain the bolometric luminosities and effective temperatures of the L and T dwarfs with known parallaxes from the literature. The relation between the effective temperature and spectral type derived from the $K$-band bolometric correction shows  monotonic behavior throughout the L-T sequence. ",
        "introduction": "The very coolest stars and the warmer brown dwarfs require a new spectral class, known as `L' \\citep{mar97,kir99}. The L dwarf sequence is characterized by the disappearance of the red TiO and VO bands from the optical (0.6$-$1.0 $\\mu$m) spectrum, the increasing dominance at those wavelengths by broad absorption resonance lines of Na I and K I, and strong H$_2$O absorption bands and persistent CO overtone bands in the 1$-$2.5 $\\mu$m region \\citep{mar99,kir99,kir00,leg01,rei01b}. In terms of broadband colors, L dwarfs are characterized by very red infrared colors (e.g. $J-K>1.3$). Even cooler brown dwarfs require the additional class, `T' \\citep{kir99}. In T dwarfs, the CO bands are replaced by stronger and more extensive absorptions of CH$_4$ in the H and K bands, and there is further strengthening of water bands \\citep{opp95,geb96,str99,bur99}. T dwarfs are characterized by blue infrared colors (e.g. $J-K \\sim 0$).   Spectral classification of L and T dwarfs including L/T transition objects has been made by \\citet{geb02} and \\citet{bur02a}. L dwarfs are now classified from L0 to L9 and T dwarfs from T0 to T8. The classification scheme is purely observational and uses indices in the 1$-$2.5 $\\mu$m region related to H$_2$O, CO, CH$_4$ and near-infrared colors. We reexamine the spectral classification using the spectra obtained at Subaru. We pay special attention to the behavior of FeH and the features related to  H$_2$O and CH$_4$.  Once the classification scheme is  established, the next step is to find the correspondence between effective temperatures and spectral type and elucidate the physical meaning of the spectral classification. Some authors have derived relations between the effective temperatures and spectral type for L dwarfs \\citep{leg02,dah02} and for L and T dwarfs \\citep{bur01}, but there has not been a work in which the relation was obtained from effective temperatures based on bolometric correction derived from photospheric models which are applicable throughout  the L-T sequence.  Some of the objects in our sample have known parallaxes. We obtain bolometric luminosities and effective temperatures for these objects by integrating the observed spectra between 0.87 and 2.5 $\\mu$m and by the $K$-band bolometric correction using the Unified Cloudy Models (UCM) \\citep{tsu02,tsu03a}. The comparison of the two methods implies that the $K$-band bolometric correction leads to reasonable effective temperatures except for early L type. Encouraged by this analysis, we apply the $K$-band bolometric correction to the objects whose parallaxes are known from the literature \\citep{bur01,dah02,tin03}.    The paper is organized as follows. In \\S2, observations, data reduction and the sample for the analysis are described. Spectral classification is discussed and identification of spectral features is reexamined in \\S3. We confirm the validity of the $K$-band bolometric correction derived by UCMs with the integrated fluxes as noted above. We then obtain bolometric luminosities of the objects with known parallaxes and the relation  between the effective temperature and spectral type in \\S4. The summary of the paper is given in \\S5. ",
        "conclusions": "We have obtained near-infrared spectra of L dwarfs, L/T transition objects and T dwarfs using Subaru. The resulting spectra are examined in detail to see their dependence on the spectral types. As for methane, we have found  that it appears at L8 and marginally at L6.5. The water bands at 1.1 and 1.4 $\\mu$m do not necessarily show steady increase towards later L types but may show inversion at late L types. This does not necessarily imply that the spectral types do not represent a temperature sequence, but can be interpreted as due to compensation of the increasing water abundance by the dust extinction in the later L types. We confirm that the FeH 0.99 $\\mu$m bands appear not only in late L dwarfs but also in the early T dwarfs as was first noted by \\citet{bur02b}. We suggest that FeH could be dredged up by the surface convective zone induced by the steep temperature gradient due to the large opacity of the dust cloud itself and will be replenished constantly by the convection.  We have obtained bolometric luminosities of the objects with known parallaxes in our sample by integrating the spectra between 0.87 and 2.5 $\\mu$m and by the $K$-band bolometric correction. Apart from the L3 dwarf, 2MASS1146+22A, the bolometric luminosities obtained by both methods agree well and this implies that the $K$-band bolometric correction obtained by UCMs can be used to obtain the bolometric luminosities and effective temperatures of the L and T dwarfs with known parallaxes from the literature. The relation between the effective temperature and spectral type derived from the $K$-band bolometric correction shows  monotonic behavior throughout the L-T sequence.  There is another method to estimate the effective temperature of a brown dwarf, which is to compare the observed spectrum and model spectra. Such an analysis is also vital as a test of model photospheres. For these purposes, we analyze the observed spectra discussed in this paper with the model spectra obtained by an extended grid of UCMs in the separate paper \\citep{tsu03}."
    },
    "0402/astro-ph0402471_arXiv.txt": {
        "abstract": "We analyze UV spectra for a large sample of 578 Type 1 Active Galactic Nuclei and derive Eddington ratios, $L/L_{edd}$, from the bolometric luminosities and emission line widths for each object in the sample.  The sample spans five orders of magnitude in supermassive black hole (SMBH) mass, seven orders of magnitude in luminosity, and a redshift range from $0 \\leq z \\leq 5$. We include a sample of 26 low-redshift Narrow-Line Seyfert 1s (NLS1s) for comparative analysis.  The NLS1s have slightly larger than average $L/L_{edd}$ ratios (and smaller SMBH masses) for their luminosities, but those $L/L_{edd}$ values are still substantially below the average for luminous quasars.  A large fraction (27\\%) of the objects overall have $L/L_{edd} > 1$, which might be explained by non-spherically symmetric accretion.  We find no trend between $L/L_{edd}$ and either redshift or SMBH mass.  Composite spectra sorted by $L/L_{edd}$ show an unusual emission-line behavior: nearly constant peak heights and decreasing FWHMs with increasing $L/L_{edd}$.  This is in marked contrast to the emission-line behaviors with luminosity, SMBH mass, and FWHM(\\ion{C}{4}), which clearly show trends analogous to the Baldwin Effect: decreasing line peaks and equivalent widths with increasing luminosity, SMBH mass, and FWHM.  The origins of the unusual behavior with $L/L_{edd}$ are not understood, but one implication is that metallicity estimates based on emission line ratios involving nitrogen show no trend with $L/L_{edd}$ in the composite spectra created from different ranges in $L/L_{edd}$.  The NLS1 composite, however, shows a slightly high metallicity for its SMBH mass and luminosity.  Our earlier work suggests that host galaxy mass, correlated with SMBH mass and AGN luminosity, is the fundamental parameter affecting BLR metallicities.  Some secondary effect, not related to $L/L_{edd}$, must be enhancing the metallicities in NLS1s. ",
        "introduction": "The central engines of quasars, and more generally, active galactic nuclei (AGNs) are believed to be powered by supermassive black holes (SMBHs). Two of the fundamental properties of AGNs are the SMBH mass and the accretion rate of material onto the SMBH.  Several indirect methods have been devised to estimate SMBH masses.  One set of these methods assumes that the broad emission-line region (BLR) is in gravitational equilibrium with the central source, so that the SMBH mass can be estimated by applying the virial theorem, $M_{\\rm SMBH} = rv^2/G$, to the measured line widths (Peterson 1993; Peterson 1997; Wandel et al. 1999; Kaspi et al. 2000; McLure \\& Dunlop 2001; Vestergaard 2002).  In reverberation mapping studies, $R_{\\rm BLR}$, the radial distance between the central source and the BLR can be estimated from the lag time between continuum variations and the emission-line response (Peterson 1993; Peterson 1997; Wandel et al. 1999; Kaspi et al.  2000).  These reverberation mapping studies have demonstrated an observed relation of $R_{\\rm BLR}~\\propto~\\lambda~ L_{\\lambda}$(5100~\\AA)$^{0.7}$ that can be used to estimate $R_{BLR}$ for AGNs over a wide range of redshifts (Kaspi et al. 2000; McLure \\& Dunlop 2001; Vestergaard 2002; Corbett et al. 2003; Warner et al. 2003).  Netzer (2003) has argued that the slope is not known to an accuracy better than about 0.15.  A second set of methods is based on the tight correlation between the masses of SMBHs and the velocity dispersions, $\\sigma$, of their host galaxy spheroidal components (Ferrarese \\& Merritt 2000; Gebhardt et al. 2000; Merritt \\& Ferrarese 2001; Tremaine et al. 2002).  However, stellar velocity dispersions are not easy to measure for AGN hosts, especially at high redshifts.  Because of this, methods have been devised using proxies of the velocity dispersion, such as the width of the narrow emission line [\\ion{O}{3}] $\\lambda 5007$ (Nelson 2000; Boroson 2003; Shields et al. 2003) or the bulge luminosity, $L_{bulge}$ (Magorrian et al. 1998; Laor 1998; Wandel 1999). Early studies showed a large scatter, as much as two orders of magnitude between SMBH mass and $L_{bulge}$ (Ferrarese \\& Merritt 2000).  However, more recent studies that carefully model the bulge light profiles of disk galaxies and thus obtain more accurate values of $L_{bulge}$ show less scatter in $M_{\\rm SMBH}-L_{\\rm bulge}$, similar to that in the $M_{\\rm SMBH}-\\sigma$ relationship (McLure \\& Dunlop 2002; Erwin et al. 2002; Bettoni et al. 2003). Recently, SMBH mass has also been shown to correlate strongly with the global structure of bulges and ellipticals, such that more centrally concentrated bulges have more massive SMBHs.  This relationship is as strong as the $M_{\\rm SMBH}-\\sigma$ relationship with comparable scatter (Graham et al. 2001; Erwin et al. 2002).  Once the SMBH mass has been estimated, the Eddington luminosity can be calculated as $L_{edd} = 1.26 \\times 10^{38} M_{\\rm SMBH}$ (\\Msun) ergs s$^{-1}$ (e.g., Rees 1984; Peterson 1997).  Eddington luminosity is the limit in which the inward gravitational force acting on the gas exactly balances the outward radiation force induced by electron scattering.  It can be thought of as the maximum possible luminosity for an object of mass $M_{\\rm SMBH}$ that is powered by {\\it spherical} accretion (Peterson 1997).  The Eddington luminosity can be exceeded if accretion is not spherically symmetric (see \\S5, also Osterbrock 1989; Begelman 2002; Collin et al. 2002; Wang 2003). AGN luminosities should be directly proportional to the accretion rate, $L \\propto \\dot{M}_{acc}$, and therefore the ratio, $L/L_{edd} \\propto \\dot{M}_{acc}/M$, is an indirect measure of the accretion rate relative to the critical Eddington value.  Narrow Line Seyfert 1s (NLS1s) are a subclass of Seyfert 1s that exhibit distinct and unusual properties: very narrow broad emission lines ($H \\beta$ FWHM $< 2000$ km s$^{-1}$) with [\\ion{O}{3}] $\\lambda 5007$ / $H \\beta$ ratios of less than 3 (to exclude Seyfert 2s), strong \\ion{Fe}{2} emission, and unusually strong big blue bumps (Osterbrock \\& Pogge 1985; Kuraszkiewicz et al. 2000; Constantin \\& Shields 2003).  NLS1s also land at one extreme end of the Boroson \\& Green (1992) Principal Component 1 (PC1).  It has been suggested that PC1 is strongly correlated with $L/L_{edd}$ (Boroson \\& Green 1992; Boroson 2002; Shemmer \\& Netzer 2002; Constantin \\& Shields 2003).  Several studies have suggested that NLS1s have low SMBH masses for their luminosities, and thus very high Eddington ratios, near 1 (Mathur 2000; Kuraszkiewicz et al. 2000; Shemmer \\& Netzer 2002; Shemmer et al. 2003).  It has also been suggested that NLS1s have unusually high metallicities for their luminosities (see \\S4.4 and Figure 11 below, Mathur 2000; Shemmer \\& Netzer 2002; Shemmer et al. 2003). Shemmer \\& Netzer (2002) find that NLS1s depart from the nominal relationship between metallicity and luminosity in AGNs (Hamann \\& Ferland 1999; Dietrich et al. 2003, in prep), with some NLS1s indicating metallicities as high as those measured in high-luminosity, high-redshift quasars.  Because of their high metallicities and high Eddington ratios, Mathur (2000) proposed that NLS1s are analogs of high-redshift ($z \\gtrsim 4$) quasars, in that they may both be in an early evolutionary phase, residing in young host galaxies.   We have collected a large sample of 578 spectra of ``Type 1\" AGNs (quasars and Seyfert galaxies with broad emission lines) that span the rest-frame UV wavelengths needed for this study  (Dietrich et al. 2002).  We compute composite spectra from different ranges in the Eddington ratio, $L/L_{edd}$.  We include a composite spectrum produced from a subsample of 26 NLS1s for comparative analysis.  We present measurements of the emission lines in these composite spectra and investigate their relationship to $L/L_{edd}$. ",
        "conclusions": "We have examined a large sample of 578 AGNs that spans five orders of magnitude in SMBH mass, seven orders of magnitude in luminosity, and a redshift range from $0 \\leq z \\leq 5$.  We estimate SMBH masses using the virial theorem and formulae given in Kaspi et al. (2000), and then derive Eddington ratios.  To improve the signal-to-noise ratio and average over object-to-object variations, we calculate composite spectra for different ranges in $L/L_{edd}$. We include a composite spectrum of a sample of 26 NLS1s for comparative analysis.  Our main results are as follows.  1) We find that a large fraction (27\\%) of the objects in our sample have $L/L_{edd} > 1$.  These super-Eddington ratios may be explained by non-spherically symmetric accretion.  While NLS1s generally show high Eddington ratios for their luminosities, the objects with the highest $L/L_{edd}$ are high-luminosity, narrow-lined quasars.  2) There is no trend between $L/L_{edd}$ and either redshift or SMBH mass.  $L/L_{edd}$ does correlate positively with luminosity and negatively with FWHM(\\ion{C}{4}), but these trends may be attributed largely to our derivation of $L/L_{edd}$ from these quantities (see Equation 4).  3) There is no trend between the shape of the UV continuum and $L/L_{edd}$.  The NLS1 composite has a much steeper (softer) continuum than the $L/L_{edd}$ composites.  This is consistent with a trend between continuum shape and SMBH mass.  4) The composite spectra sorted by $L/L_{edd}$ exhibit an unusual emission-line behavior: nearly constant peak heights and decreasing FWHMs with increasing $L/L_{edd}$ (Figure 2).  The origins of this behavior are not understood, but it is in marked contrast to the emission-line behaviors in composite spectra sorted by luminosity, SMBH mass, and FWHM(\\ion{C}{4}) (Figure 8), all of which clearly show trends in the line REWs analagous to the Baldwin Effect.  5) The composite spectra show no trend between $L/L_{edd}$ and metallicity (Figure 10).  This is consistent with SMBH mass being related to the fundamental parameter affecting BLR metallicity (Warner et al. 2003).  6) The NLS1 composite exhibits several unusual behaviors.  It generally does not fit the trends between emission line REWs and $L/L_{edd}$ as defined by the $L/L_{edd}$ composites.  It also has a metallicity that is slightly high for its average SMBH mass and luminosity, although still well below the high metallicities exhibited by the most luminous quasars with the most massive central SMBHs.  The quasars with the highest $L/L_{edd}$, high-luminosity quasars with narrow \\ion{C}{4} emission, do not have high metallicities for their SMBH masses and luminosities.  Our earlier work (Warner et al. 2003) is consistent with the theory that host galaxy mass, correlated with SMBH mass (and AGN luminosity), is the fundamental parameter affecting BLR metallicity.  We conclude that i) there must be some secondary effect enhancing the metallicity in NLS1s, and ii) this secondary effect is not related to $L/L_{edd}$.   \\noindent {\\it Acknowledgements:} We are very grateful to Marianne Vestergaard for providing the UV Fe emission template for this study, and to Fred Chaffee, Anca Constantin, Craig Foltz, Vesa Junkkarinen, and Joe Shields for their direct participation in reducing or acquiring some of the ground-based spectra.  We acknowledge financial support from the NSF via grant AST99-84040 and NASA via grant NAG5-3234."
    },
    "0402/astro-ph0402192_arXiv.txt": {
        "abstract": "We report an attempt to  interpret the spectra of L and T dwarfs with the use of the Unified Cloudy Model (UCM). For this purpose, we extend the grid of the UCMs to the cases of log $g = 4.5$ and 5.5. The dust column density relative to the gas column density in the observable photosphere is larger at the higher gravities, and molecular line intensity is generally smaller at the higher gravities. The overall spectral energy distributions (SEDs) are $ f_{J} < f_{H} < f_{K} $ in middle and late L dwarfs, $ f_{J} < f_{H} > f_{K} $ in early T dwarfs (L/T transition objects), and finally $ f_{J} > f_{H} > f_{K} $ in middle and late T dwarfs, where $ f_{J}, f_{H}$, and $f_{K} $ are the peak fluxes at $J, H,$ and $K$ bands, respectively, in $f_{\\nu}$ unit. This tendency is the opposite to what is expected for the temperature effect, but can be accounted for as the effect of thin dust clouds formed deep in the photosphere together with the effect of the gaseous opacities including H$_2$ (CIA), H$_2$O, CH$_4$, and K I. Although the UCMs are semi-empirical models based on a simple assumption that thin dust clouds form in the region of $ T_{\\rm cr} \\la T \\la T_{\\rm cond}$ ($ T_{\\rm cr} \\approx 1800$\\,K is an only empirical parameter while $ T_{\\rm cond} \\approx 2000$\\,K is fixed by the thermodynamical data), the major observations including the overall SEDs as well as the strengths of the major spectral features are consistently accounted for throughout L and T dwarfs. In view of the formidable complexities of the cloud formation, we hope that our UCM can be of some use as a guide  for future  modelings of the ultracool dwarfs as well as for interpretation of observed data of L and T dwarfs. ",
        "introduction": "So far, few models are available for interpretation and analysis of the spectra of L and T dwarfs consistently. Especially, it is well recognized that dust forms in the photosphere of L dwarfs, but it is by no means clear how to take the effect of dust into account in the predictions of the spectra and the spectral energy distributions (SEDs). Our initial attempt simply assumed that dust forms everywhere so long as the thermodynamical condition of condensation is met (Tsuji, Ohnaka, \\& Aoki 1996). Although such models could  explain the spectra of late M dwarfs and early L dwarfs \\citep[e.g.][]{jon97, tsu00, sch01}, at least qualitatively, they failed to explain the spectra of cooler L dwarfs as well as of T dwarfs. In fact, the photospheres will soon be filled with dust if the simple thermochemical equilibrium including condensation is assumed, and the optical thickness of dust is so large that the predicted spectra from such a model  will simply be a blackbody radiation of  $T = T_{\\rm eff}$ for $ T_{\\rm eff} \\la 1500$K or so (Tsuji 2000, 2001). The fully dusty models by other authors \\citep[e.g.][]{all01} may have the same difficulty. On the other hand, cool T dwarfs, whose prototype is Gl 229B, show no evidence of dust in their spectra. A naive interpretation was  that the dust may have segregated from the gaseous mixture and precipitated below the photosphere \\citep[][]{tsu96b, mar96, all96, feg96}. However, a question is why such  segregation of the dust took place only in cool T dwarfs.  As a possibility to resolve such difficulties, we proposed a new model which we referred to as the unified cloudy model, UCM \\citep[][]{tsu01} and extended it to a grid (log $g$ =5.0 and $ 800 \\le T_{\\rm eff} \\le 2600$\\,K) for applications to L and T dwarfs (Tsuji 2002; hereafter referred to as Paper I).  In the UCMs, the segregation of dust from the gaseous mixture takes place in all the ultracool dwarfs including L and T dwarfs and at about the same temperature referred to as the critical temperature $T_{\\rm cr}$. Then, roughly speaking,  the dust will remain in the observable photosphere for the relatively warm dwarfs  with $T_{\\rm eff} > T_{\\rm cr}$ (note that $T \\approx T_{\\rm eff}$ at $\\tau \\approx 1$ and hence the region of $T \\ga T_{\\rm cr}$, where dust still survives,  is found in the optically thin region), and hence such warm dwarfs as L dwarfs will appear to be dusty. In the cooler dwarfs  with  $T_{\\rm eff} < T_{\\rm cr}$, on the other hand, the optically thin region (i.e. $\\tau < 1$ and hence $ T < T_{\\rm eff}$) will be cooler than  $T_{\\rm cr}$ and all the dust grains there will be segregated and precipitated. For this reason, such cool dwarfs as T dwarfs will appear to be dust-free. It is to be noted that this assumption behind the UCM  is physically more natural than to assume that dust once formed never segregate throughout the photosphere ( namely the fully dusty model of case B) or all the dust grains segregate as soon as they are formed (i.e.  fully dust-segregated model of case C).  The UCM, however, is by no means a self-consistent theoretical model, but rather it is a kind of semi-empirical model at present. It should be emphasized, however, that  empirical approach often plays an important role in modeling stellar photospheres and atmospheres, even in the more simple cases where dust plays no role. For example, empirical models are still widely used for the solar photosphere, not to speak of the solar atmosphere (i.e. whole the observable layers including the photosphere, chromosphere, CO-mosphere, transition layer, corona etc.) for which no fully theoretical model may yet exist. Once the phase transition occurs in the photosphere, it will introduce complicated phenomena such as those familiar in the meteorology, and it appears to be more difficult to build a fully theoretical model from the beginning. Instead, we hope to understand the basic features of the dust in L and T dwarfs with the simplest possible semi-empirical model which is consistent with the known observations as well as with the basic physics such as thermodynamics.  We notice that some attempts have been made in theoretical modelings of the dust formation in L and T dwarfs \\citep[e.g.][]{ack01, hel01, mar02, coo03, woi03}, but it is not yet clear if they provide consistent interpretation of the major observations throughout L and T dwarfs.  So far, we have already shown that the UCMs provide reasonable account for the L/T transition on the color-magnitude (CM) diagram \\citep[][] {tsu03a} as well as major observations such as infrared colors and spectra of ultracool dwarfs throughout L and T dwarfs (Paper I). This fact implies that the UCMs  may represent the physical structure of L and T dwarfs to some extent. As a next step in observational tests of the UCMs, we examine if the calibrated spectra observed with the Subaru Telescope, as detailed in a separate paper (Nakajima, Tsuji, \\& Yanagisawa 2004), can be fitted with the predicted spectra based on the UCMs.  For this purpose, we first discuss some details of the UCMs and extend them to cover the possible range of the surface gravities and effective temperatures (Sect.\\,2). Next, we discuss the dependence of the observable properties on the basic stellar parameters (Sect.\\,3). Then we focus our attention on interpreting the spectral energy distributions or the spectra of  L, L/T transition objects, and T dwarfs based on a single grid of UCMs (Sect.\\,4). Although we confirm that the observed spectra can reasonably be accounted for by the UCMs, many problems remain unsolved before a more detailed  confrontation between models and observations can be possible (Sect.\\,5). ",
        "conclusions": "We have shown that the spectra or SEDs of L and T dwarfs can be interpreted consistently by a single grid of UCMs. At present, we cannot yet achieve a fully self-consistent model photosphere of ultracool dwarfs because of the complexities due to the coupling of physico-chemical processes relating to the cloud formation and associated dynamical processes. Instead, we restricted ourselves to a semi-empirical approach which is based only on a simple thermodynamical constraint, and reduced all the possible complicated dynamical effects to a quasi-static model photosphere to be treated by the classical non-grey theory. It is to be noted that the model photosphere itself is not necessarily our final purpose, but our purpose is to understand the real astronomical objects, in this case, L and T dwarfs. The model photosphere is simply a means by which to help this aim, even though better models are certainly more useful for this purpose. Thus the aim of our UCMs is not to provide the exact quantitative fits to observed data at present. It is hoped that our semi-empirical approach can be of some help as a guide to interpret and analyze the observed data of ultracool dwarfs, and hopefully will provide a guideline by which a more physical model can be developed in the near future. To be of some use for this purpose, the numerical data of the UCMs, including the spectra and SEDs, are  made available through our Web site \\footnote{http://www.mtk.ioa.s.u-tokyo.ac.jp/\\~\\,ttsuji/export/ucm}."
    },
    "0402/astro-ph0402647_arXiv.txt": {
        "abstract": "A widely discussed explanation for the origin of the X-ray emission observed from knots in extended quasar jets with the {\\it Chandra X-ray Obseratory} is Compton-scattered CMBR by electrons with Lorentz factors $\\gamma^\\prime \\sim 10^2$.  This model faces difficulties in terms of total energy requirements, and in explaining the spatial profiles of the radio, optical, and X-ray knots in sources such as PKS 0637-752, 3C 273, or PKS 1127-145.  These difficulties can be resolved in the framework of one- and two-component synchrotron models.  We propose a model where the broad band radio to X-ray synchrotron emission in quasar jets is powered by collimated beams of ultra-high energy neutrons and gamma-rays formed in the sub-parsec scale jets. The decay of the neutral beam in the intergalactic medium drives relativistic shocks to accelerate nonthermal electrons out of the ambient medium.  A second synchrotron component arises from the injection of leptons with Lorentz factors $\\gg 10^7$ that appear in the extended jet in the process of decay of ultra-high energy gamma rays. This approach could account for qualitative differences in the extended X-ray jets of FR1 and FR2 galaxies. Detection of high-energy neutrinos from blazars and core-dominated quasars will provide strong evidence for this model. ",
        "introduction": "Three main radiation processes considered to account for the nonthermal X-ray emission in knots and hot spots of the extended jets discovered by the {\\it Chandra X-ray Observatory} are synchrotron, synchrotron self-Compton (SSC), and Compton scattering of external photons contributed mostly by the CMBR (e.g., \\citet{hk02,sta03}).  In X-ray knots of quasar jets  with projected lengths $\\gtrsim 100$ kpc, where the X-ray spectrum is not a smooth extension of the radio/optical spectrum, a currently favored interpretation is the external Compton (EC) model \\citep{tav00,cgc01}. In this model, the X-ray emission from knots such as WK7.8 of PKS 0637-752 is argued to be due to CMB photons that are Compton-upscattered by nonthermal electrons from kpc-scale emitting regions in bulk relativistic motion at distances up to several hundred kpc away from the central engine.  We note, however, that in the framework of this model it is problematic to explain a clear trend observed from many extended jets of radio quasars, such as of 3C 273 \\citep{mar01,sam01} or of PKS 1127-145 \\citep{siem02}, that the X-ray brightness of the knots decreases with distance along the jet while the radio flux is {\\it increasing}.  The Lorentz factors of the X-ray emitting electrons in the EC model are smaller than those of the radio- and optical-emitting electrons. It is therefore difficult to explain comparable knot sizes at optical and X-ray frequencies, but extended emission at radio frequencies. Moreover, the radiative cooling of these low-energy electrons is slow, resulting in high total energy requirements.  In this paper, we examine these difficulties and propose an alternative interpretation for X-ray emission from knots in FR2 radio galaxy jets and quasars where the radio through X-ray fluxes are explained by synchrotron radiation from one or two components of relativistic electrons. Both of these components are powered by neutral beams of ultra-high energy (UHE) neutrons and gamma-rays produced by the jet at the base of the central AGN engine \\citep{ad01,ad03}.  The momentum of the decaying beam of neutrons drives a relativistic shock that accelerates electrons from the surrounding medium to produce the main nonthermal electron component responsible for the broad band radio/optical/X-ray synchrotron emission. In such sources, a second ultra-relativistic lepton component can also be injected from neutron $\\beta$ decay and the pair production by UHE $\\gamma$ rays attenuated by the CMBR. The synchrotron emission of these pairs can contribute to or even dominate the synchrotron X-ray flux in the knots of core-dominated quasars.  Here we focus on knots in quasar jets, though the same 2-component synchrotron model can apply to hot spots of FR2 galaxies such as Cygnus A \\citep{wys00} or Pictor A \\citep{wys01}, where the SSC process could also play a role. For the X-ray jets in FR1 sources, with projected lengths of only $\\lesssim 5$ kpc, X-ray and optical emission are generally consistent with a smooth power-law continuation of the radio spectrum, indicating a synchrotron origin. Mild spectral hardenings at X-ray energies, such as those observed in the knots of M87's jet \\citep{wy02,mar02}, could still be explained within the context of a single component synchrotron model (Dermer and Atoyan, 2002; henceforth DA02).  In FR1 galaxies and BL Lac objects, the neutral beam power is considerably smaller than in FR2 galaxies and quasars because of the weaker external radiation field in the inner jet \\citep{ad01}, so that extended jets in FR1 galaxies would primarily result from the jet plasma expelled directly from the central nucleus.  In Section 2 we discuss difficulties with the EC model, and we propose one- and two-component synchrotron models for the X-ray jets in Section 3. The rationale for the model is discussed in Section 4, including testable predictions from X-ray and high-energy neutrino observations and a brief discussion of this model in the context of a scenario for AGN evolution. ",
        "conclusions": "The external Compton model for extended X-ray jets faces difficulties with large energy requirements, and in explaining the different spatial profiles at radio, optical, and X-ray frequencies. It also cannot be invoked to explain knots and hot spots where the X-ray spectral indices are significantly steeper than the radio spectral indices, in which case a synchrotron model is favored.  The large energy requirements in the EC model can be reduced to acceptable values only by assuming Doppler factors $\\delta \\geq 10$. In these cases, the jet would have to be directed towards us at very small angles of only $\\theta \\leq 6^\\circ $.  Although we cannot exclude that the number of X-ray jets with such small angles could still be significant, the probability for one of the 2 jets of any quasar to be directed within such an angle is only $P_\\theta = (1-\\cos \\theta)/2\\pi \\leq 8.7\\times 10^{-4} (\\theta/6^\\circ)^2$. Allowing for X-ray jets with $\\theta \\lesssim 20^\\circ$ will increase the probability to detect several such sources by several orders of magnitude.  A common model for X-ray knots would then have to be able to deal with sufficiently large $\\theta$ and to allow debeamed jets with $\\delta \\ll \\Gamma \\simeq 10$.  The energy demands for such jets become unrealistic in the EC model.  The large energy requirements are significantly relaxed for synchrotron models. Synchrotron X-rays are produced by very high energy electrons with short cooling times, which minimizes the injection power of electrons needed to explain the X-ray flux. For a two-component synchrotron model, unlike for the EC model, the magnetic field is not limited by the ratio of X-ray to radio fluxes and a given inclination angle (obtained, for example, by fixing the maximum possible Dopler factor). This minimizes the total energy accumulated in the form of radio emitting electrons and magnetic fields in the knot needed to explain the observed radio flux. The comoving equipartition fields in FR2 knots are typically at the level $\\sim 50$--$100\\,\\rm \\mu G$, which are higher than $B \\sim 10$-$20 \\,\\rm \\mu G$ deduced from EC models. The synchrotron interpretation of X-rays also allows one to have significantly debeamed jets and still remain within an acceptable range for the energy budget, as in the knots of 3C 273 (Fig.~1).   The relative advantage of synchrotron models for debeamed jets with $\\delta << \\Gamma$ is due essentially to the more narrow beaming diagram for the EC radiation than for the synchrotron radiation in the stationary frame \\citep{d95}. This is the case when the photons are detected at an angle $\\theta \\gg 1/\\Gamma$. In the comoving frame such photons are produced in Compton `tail-on' collisions with beamed (in the comoving frame) external photons at angles $(1-\\cos \\theta^\\prime)\\ll 1$; therefore the efficiency of the Compton process is strongly reduced. In the meantime, synchrotron radiation production in a quasi-isotropic knot magnetic field is isotropic in the comoving frame.   The steeper SED in the optical than X-ray bands in some knots, as A1 or B1 knots of 3C 273 in Fig.~1, can be explained in the framework of a single population of electrons accelerated to multi-TeV energies in the shocks forming the knots (DA02). The two-component synchrotron model appears more general. In particular, it can explain practically all types of SEDs detected from the knots and hot spots of multi-kpc scale jets of FR2 radio galaxies, such as the knot WK7.8 of PKS 637-752 where the optical flux is a factor $>10$ times below the power-law extrapolation of fluxes from the radio to X-ray bands.   We have argued that the origins of the two different components in the knots and hot spots of FR2 jets can be naturally explained within a unified model for large scale jets as manifestations of neutral beams \\citep{ad01,ad03}.  These beams, composed of UHE neutrons, gamma-rays and neutrinos, are efficiently produced in the compact relativistic jets of FR2 quasars by accelerated UHE protons in the process of photomeson interactions with the accretion disk radiation on $\\lesssim$ pc scales. The decay of UHE neutrons deposits momentum and energy of the neutral beam into the intergalactic medium, and drives surrounding plasma to relativistic motion (presumably with moderately relativistic speeds, though this will require a hydrodynamic study to quantify) through interactions with ambient magnetic fields and the generation of plasma waves.  The appearance of knots in continuous jets could be manifestations of shocks resulting from the increased activity of the central engine during some time in the past \\citep{ner02,ad03}. The first population of electrons is produced in the process of first-order Fermi acceleration of electrons from the surrounding intergalactic medium. The radio electrons cool slowly; therefore they drift and accumulate along the jet, which explains the increase of radio brightness.  A second population of electrons in the jet originates in the decay of UHE gamma-rays with Lorentz factors $\\gamma_0 \\gtrsim 10^{10}-10^{13}\\,\\rm eV$ \\citep{ad03}, and can be contributed even by $\\beta$ electrons from neutron decay. It is essential that the production spectrum of these electrons is cut off at lower energies, $\\gamma_c \\lesssim \\gamma_0 $. Because of the strong Klein-Nishina effect for the Compton radiation of UHE electrons, most of the electron energy will be deposited in the form of synchrotron radiation at MeV/GeV energies in ambient fields $B \\gtrsim 1 \\mu$G. This radiation can be effectively lost for an observer outside of the jet opening angle. These electrons will be cooled down to energies $\\gamma \\gtrsim 10^{8}-10^{9}$ before being overtaken by the relativistic shock.  Reprocessing of even a small fraction of the $\\gamma$-ray beam energy in the pair-photon cascade along the jet, however, will result in a very significant contribution to the total number of UHE electrons, with a broad spectrum above $\\gamma \\sim 10^8$ and a cutoff at lower energies. These electrons are deposited throughout the length of the jet, and will be found, in particular, in the jet fluid in front of the shock.  Convection of these electrons downstream of the shock with enhanced magnetic field results in a second synchrotron component that can explain the anomalously hard X-ray SEDs of many knots.  Note that this component of electrons is re-energized at the shock front in the process of adiabatic compression.  Moreover, because of the high efficiency of synchrotron radiation, in many cases even the energy of the neutral beam deposited in the $\\beta$-decay electrons alone could be sufficient for the explanation of the low X-ray fluxes observed.  We note that MHD waves excited by the beam, as in the shear boundary layer model of \\citet{so02}, could also contribute to the reacceleration and energization of a hard electron component.  The decline of the deposition rate of UHE electrons with distance $r$ from the AGN core could be as fast as $\\propto r^{-2}$ in case of a dominant $\\beta$-decay electron contribution. The decline could be somewhat slower in the case of a significant contribution from pair-photon cascades of UHE electrons. This would then explain the fast decline of X-ray emission of knots along the jet.  The final hot spot in our model would represent a termination shock of the relativistic flow when the jet runs out of UHE neutral-beam decay products that sustain the forward progress of the relativistic jet into the ambient medium. Compression and additional acceleration of the second component of electrons (cooled down to $\\gamma_{min} \\gtrsim 10^{6}$ on timescales up to $1\\,\\rm Myr$ for hot spots at distances $\\lesssim 300 \\,\\rm kpc$) and the increase of the magnetic field in the downstream region of this relativistic jet termination shock can explain the spectrum of the hot spot of Pictor A with $\\alpha_X \\geq 1$.  For very distant quasars with $z\\gtrsim 3$, such as reported recently \\citep{sch02,siem03}, an additional powerful channel contributing to the second electron component in large scale jet appears.  At such redshifts, the photomeson interaction length of super-GZK neutrons becomes comparable with and even less than the Mpc scale of the jets. In this case, the input from the super-GZK neutron component of the UHE neutral beam and the efficiency of the electromagnetic cascade will increase dramatically. In this way, the increased energy density of the CMBR can explain the detectability of X-rays even from quasars at large redshifts (compare different ideas by \\citet{sch02a} for the EC model).  We propose two testable predictions of this scenario. The first is variability of the X-ray flux from the knots of FR2 radio galaxies. This possibility, suggested by observations of X-ray variability in the knots of M87 \\citep{per03}, is a consequence of the short cooling time scale $t_{syn} \\simeq 500 \\sqrt{(1+z)/E_{keV} \\delta B_{30}^3}\\,$yr of X-ray emitting electrons. For the comoving magnetic fiels in the knots (and hot spots) reaching $B^\\prime \\gtrsim 100\\,\\rm \\mu G$ this time can be $\\lesssim 50/\\sqrt{E_{keV}} \\,\\rm yr$ even for knots/jets with moderate Doppler factors $\\delta \\lesssim 3$ moving at rather large angles $\\theta \\lesssim 20^\\circ$.  One could then expect that variations of the injection rate of electrons at the shock (swept up in the upstream region) on time scales $t_{var}\\gtrsim t_{syn}$ would result in variations of the X-ray flux on the same scales.  Detection of variability could be feasible for X-rays knots with transverse sizes $\\lesssim 1\\,$kpc, but which are sufficiently thin (flat) along our line-of-sight. The detection of a small-amplitude variability (at the level of up to several per cent over reasonable observation times) variability depends on the strength of the magnetic field downstream of the shock.  For knots detected with {\\it Chandra} at the X-ray count rates $\\gtrsim 100$ per hour, like from knot WK7.8 in PKS 0637-752 \\citep{cha00}, simple estimates show that a statistically significant flux variations could be expected for observations separated by several years.  X-ray variability of an FR2 knot or hot spot would be difficult to understand in the context of the EC model with the large cooling times of the emitting electrons and stationary target for Compton interactions.  A second prediction is that high-energy ($\\gtrsim 10^{14}$ eV) neutrinos will be detected from core-dominated quasars with km-scale neutrino telescopes such as IceCube.  Detection of every such neutrino from a flat spectrum radio quasar such as 3C 279 implies that a total energy $\\sim 5\\times 10^{54} \\delta^{-2} \\,\\rm erg $ is injected in the inner jet at sub-kpc scales \\citep{ad03}. Therefore even the detection of 1 neutrino per year would imply an UHE neutral beam power $\\sim 10^{46} \\delta^{-2} \\,\\rm erg \\, s^{-1}$.  Finally, we note that this model is developed in view of a scenario \\citep{bd02,ce02} where radio loud AGNs are formed by merging IR luminous galaxies, and evolve from a high luminosity FR2/quasar phase at early cosmic times to a lower luminosity FR1/BL Lac phase at later times, thus explaining the sequence of flat-spectrum radio quasar and BL Lac object SEDs \\citep{fos98}. The crucial feature in this scenario is the amount of broad-line emitting gas that fuels the AGN, and the corresponding intensities of the external radiation field in the vicinity of the inner jet. This radiation field determines the power of the neutral beam that produces the distinctive lobe-dominated morphologies of FR2 radio galaxies and, in exceptional cases such as Pictor A, a linear jet \\citep{wys01}. FR1 galaxies, where the neutral beam power is negligible, consequently display very different radio jet morphologies.  \\vskip0.5in We thank Herman Marshall, Eric Perlman, and Andrew Wilson for comments and discussions,  and the referee for useful comments. AA appreciates the hospitality of the NRL High Energy Space Environment Branch during his visit when this work has been started. The work of CD is supported by the Office of Naval Research and {\\it GLAST} Science Investigation No.\\ DPR-S-1563-Y."
    },
    "0402/astro-ph0402537_arXiv.txt": {
        "abstract": "{In dense stellar systems the frequent dynamical interactions between stars play a crucial role in the formation and evolution of compact binaries. We study these processes using a novel approach combining a state-of-the-art binary population synthesis code with a simple treatment of dynamical interactions in dense star cluster cores. Here we focus on the dynamical and evolutionary processes leading to the formation of compact binaries containing white dwarfs in dense globular clusters. We demonstrate that dynamics can increase by factors $\\sim 2 - 100$ the production rates of interesting binaries such as cataclysmic variables, ``nonflickerers'' (He white dwarfs with a heavier dark companion), merging white dwarf binaries with total masses above the Chandrasekhar limit, and white dwarf binaries emitting gravitational waves in the LISA band.} ",
        "introduction": "From the earliest observations of X-ray binaries in globular clusters it has always been clear that they must be very efficient factories for the production of compact binary systems \\citep{1975ApJ...199L.143C}. The overabundance of compact binaries in clusters, as compared to the field, must be a result of close stellar encounters. The key processes that affect the binary population in dense cluster environments include the destruction of wide binaries, hardening of close binaries (following ``Heggie's law'' \\citep{Heggie_75}), and exchange interactions, through which low-mass companions tend to be replaced by a more massive participant in the encounter. As a result of these processes, in the dense cores of globular clusters, binaries are strongly depleted and their period distribution is very different from that of a field population (Ivanova et al.\\ 2004). These processes also lead to an interesting and complex interplay between dynamics and binary evolution. For example, exchange interactions involving compact objects often produce systems that will evolve through a common-envelope (CE) phase and form very short-period binaries, which are much less common in field populations \\citep{2000ApJ...532L..47R}.  There are two possible approaches to the study of binary evolution and dynamics in globular clusters. One can start from $N$-body simulations and introduce various simplified treatments of binary star evolution. This has been the traditional approach for many years \\citep[for recent examples see][]{SharHurley_2002,2001MNRAS.321..199P}. Alternatively, one can start from a binary population synthesis code and add a treatment of dynamical interactions. This approach was pioneered by Portegies Zwart et al.\\ (1997) and has been adopted in our recent work. It has the great advantage that it is computationally much less expensive than $N$-body simulations, so that more exploration of the (enormous) parameter space is possible, and more realistic simulations, using sufficiently large numbers of stars and binaries, are possible today. In contrast, even when using special-purpose GRAPE computers, $N$-body simulations are still limited to smaller systems like open clusters with limited coverage of parameter space and with unrealistically small numbers of binaries \\citep[see][]{BinFrac_04, 2003MNRAS.343.1025W}.  In our code we combined {\\tt StarTrack}, a state-of-the-art  binary population synthesis code \\citep{Chris_02} and {\\tt FewBody}, a small-$N$-body integrator that we use to compute 3-body and 4-body interactions \\citep{Fregeau_FB2_04,Fregeau_FB1_04}. Currently we adopt a simple two-zone model, in which the cluster is partitioned into an inner core and an outer halo, with all interactions assumed to take place in the core. This background cluster model remains unchanged throughout the evolution \\citep{hut_1992}. In particular, the core density is assumed constant. However, our ultimate aim is to incorporate full dynamical Monte Carlo models \\citep{2003ApJ...593..772F}. In a typical simulation we start with  $N\\sim 10^5$ stars, with between 50\\% and 100\\% binaries. This high primordial binary fraction (much higher than assumed in all previous studies) is needed in order to match the observed binary fractions in globular cluster cores today \\citep{BinFrac_04}. ",
        "conclusions": ""
    },
    "0402/hep-th0402046_arXiv.txt": {
        "abstract": " ",
        "introduction": "Current cosmological observations suggest that it may be important to study the effective four-dimensional gravitational theory derivable from a fundamental theory, like M/string theory. If the effective theory has local supersymmetry, it is described by $d=4$, $N=1$ supergravity. For cosmological applications one is interested particularly in any possibility to find a de Sitter type configuration (dS) with broken supersymmetry to describe the currently accelerating universe as well as a slow-roll stage of the early universe inflation.   The potential in $N=1$ supergravity\\footnote{We will limit ourselves to gravitational, vector and chiral multiplets.} is well known: there is an F-term potential, constructed in a standard way from the K{\\\"a}hler potential and superpotential, in some cases there is also a non-trivial $D$-term potential, derivable in the standard fashion \\cite{Cremmer:1983en,Wess:1992cp,Binetruy:2000zx}. It is extremely important in the context of cosmological applications that the $D$-term potential is always positive, whereas the F-term in general has both positive and negative contributions.  Since the $D$-term potential is positive definite it may lead to de Sitter type solutions, particularly in presence of a constant FI term.  At present there is no known way to derive the effective  $d=4$, $N=1$ supergravity with constant FI terms from M/string theory. Only field-dependent $D$-terms have been identified so far (one should keep in mind that FI terms studied in the context of open string theory are not immediately relevant for  gravity and cosmology). There is no strict no-go theorem about the absence of constant FI terms in string theory, however, for all practical purposes, the situation is close to the existence of such theorem.\\footnote{We are grateful to S. Kachru and J. Maldacena for numerous discussions of this issue.}  Since constant FI terms in effective supergravity in 1+3 dimensions lead to dS spaces, the possibility to get such terms from M/string theory may require new developments in the understanding of string theory. It is worth reminding here that the second string revolution has allowed to treat M-theory and 11-dimensional supergravity  as leading to effective theories with 1+3 dimensional chiral fermions. Before  the compactifications on  orbifolds and orientifolds were studied, it was believed that it is impossible to get $d=4$ chiral fermions from 11-dimensional supergravity. At present we may only hope that some new possibility will realize  in M/string theory that will allow  to derive constant FI terms and dS and near dS spaces.   The purpose and the results of this paper can be summarized as follows:     \\begin{itemize} \\item  Firstly,   we will  study the general case when local $N=1$ supersymmetry in 1+3 dimensions admits constant FI terms and provide the supersymmetry rules in such theories.\\footnote{In globally supersymmetric theories the constant FI terms can be added   without any  constraints on the theory. However, in the local case this is not true anymore.} As an application of these rules we will revisit the $D$-term inflation model and correct the supergravity version of it to comply with the  restrictions on the superpotential required when  constant FI terms are present. We will also study $D$-term strings and their properties in supergravity with constant FI terms.  \\item Secondly, we will study the so-called anomaly generated FI terms originating from string theory and explain that a procedure of stabilization of certain moduli is required for these models to be used in the cosmological context in 1+3 dimensions.    \\item The recently suggested\\cite{Dvali:2003zh} equivalence between the $4d$ supergravity $D$-term strings and D-branes of type $II$ theory shows that brane-anti-brane systems in an effective $4d$ theory can be described as gauge theories with non-zero FI term. The axion shifting under the $U(1)$-symmetry is dual to the Ramond-Ramond form. This connection gives the possibility to establish a useful dictionary between the two descriptions and interpret many important properties of D$-\\bar{\\mbox{D}}$ systems in a simpler language of $4d$ supergravity gauge theories with non-zero $D$-terms, and vice versa. In particular, we can use our knowledge of the stability of supergravity vacua with non-zero $D$-terms for understanding the stability of the string vacua with brane-anti-brane systems and their various cosmological applications, such as $D$-brane inflation. % For instance, the shift of the axion in both cases restricts the possible forms of the moduli-stabilizing superpotential.  We provide some additional consistency checks of the correspondence between $D$-term-strings and D-branes and show that, not surprisingly, the instabilities of the two are closely related. Thus, not only the cosmic $D$-term-strings, formed after $D$-term inflation, do not cause any cosmological trouble, but in fact they may be potentially detected. Hence, the D-brane strings of type $II$ theory, could in principle be observed in the sky in the form of the supergravity $D$-term strings! \\end{itemize}  A standard expectation is that any K{\\\"a}hler potential and any holomorphic superpotential may define a version of $N=1$ supergravity in 1+3 dimensions. We will clarify here  the situation with constant FI terms, when this expectation is not valid and certain restrictions on the choice of the superpotential are required.   The very first version of supergravity with  locally supersymmetric extension of the FI term of the Abelian vector multiplet was constructed in \\cite{Freedman:1976uk}. It has positive cosmological constant, $\\Lambda >0$.  It was also shown there that local supersymmetry requires the axial gauging of  gravitino and gaugino (local $R$-symmetry). This theory involves only the gravitational supermultiplet and the vector supermultiplet and  has one-loop  axial anomalies  \\cite{DG}. More general classes of models with constant FI terms and scalar supermultiplets were constructed in \\cite{Stelle:1978wj,Barbieri:1982ac,Ferrara:1983dh,Kallosh:2000ve}. More recently there were few important developments in studies of some anomaly-free models with gauged $R$-symmetry and constant  FI terms in supergravity \\cite{Chamseddine:1995rs,Castano:1996ci}. At that time the main focus of such investigations was towards particle physics with vanishing cosmological constant.  On the other hand, in the cosmology community, the role of $D$-terms has become extremely important as a possible origin of de Sitter configurations and inflation in supergravity \\cite{Binetruy:1996xj,Halyo:1996pp}. It remains, however,  not well known that the presence of constant FI terms poses specific restrictions on supergravity theories (see however \\cite{Barbieri:2002ic,Arkani-Hamed:2003mz}).  The existing versions of supergravities with FI terms are mostly incomplete for our purposes. The $D$-term inflation model has an important property that in the unstable de Sitter vacuum as well as in the absolute Minkowski vacuum the superpotential vanishes, $W_{\\rm min}=0$. However, outside the minimum, the superpotential does not vanish, $W\\neq 0$. Thus formulations of supergravity \\cite{Cremmer:1983en,Ferrara:1983dh}, where the Lagrangian depends not on two functions, the \\Ka\\ potential ${\\cal K}(z, z^*)$ and the superpotential $W(z)$, but only on one combination ${\\cal G}(z, z^*) = -{\\cal K} (z, z^*) - \\ln |M_P^{-3}W|^2$ are not suitable\\footnote{In \\cite{Ferrara:1983dh} there is a short ``Note added'' how to treat the case with vanishing superpotential.} since they are not well defined at $W=0$.  The superspace approach with a non-singular dependence on the superpotential $W$ presented in \\cite{Barbieri:1982ac,Wess:1992cp} has all terms depending on constant FI. However, the holomorphic kinetic function $f_{\\alpha\\beta}(z)$ for the vector multiplets is the simplest one, equal to 1. On the other hand, in \\cite{Binetruy:2000zx} where there is an arbitrary scalar dependent $f_{\\alpha\\beta}(z)$, the constant FI terms are not introduced. The significance of a generic, scalar dependent $f_{\\alpha\\beta}(z)$ has to do with axial coupling $aFF^*$ which sometimes plays an important role in the mechanism of anomaly cancellation. \\bigskip  In section~\\ref{ss:superconfAction} we give a summary of the ingredients of the construction of the supergravity action with superconformal symmetry. For our purpose it is most useful to study the formulation of supergravity with the superconformal origin which was recently constructed in \\cite{Kallosh:2000ve}.  It has all 3 generic functions, the \\Ka\\ potential ${\\cal K}(z, z^*)$, the superpotential $W(z)$, and the kinetic function $f_{\\alpha\\beta}(z)$ for the vector multiplets  and the theory is regular at $W=0$. One furthermore has to define the symmetry transformations. This includes for any $U(1)$ factor the possible occurrence of a FI constant $\\xi _{\\alpha i}$.    In the superconformal approach, one constructs in a first step the action with full superconformal symmetry. It contains an extra chiral multiplet, which was often called `compensating multiplet', but was baptized `conformon' in~\\cite{Kallosh:2000ve} to reflect its significance. In the next step, the gauge symmetries that are not present in Poincar{\\'e} supergravity, such as local dilations, local chiral $U(1)$-symmetry\\footnote{In the context of superconformal $SU(2,2|1)$ symmetry  one often calls the superconformal chiral $U(1)$ symmetry ``$R$-symmetry'', since it rotates the supercharges (transforms the gravitinos). However, in the Poincar{\\'e} supergravity, after the superconformal $U(1)$ symmetry is fixed, there are sometimes other $U(1)$ symmetries which are  combinations of the superconformal $U(1)$ symmetry and some additional $U(1)$ gauge symmetries which were present in the superconformal theory. These  $U(1)$ symmetries  one also calls $R$-symmetries. In what follows, we will use the term ``local $R$-symmetries'' in the context of Poincar{\\'e} supergravity only.} and local $S$-supersymmetry, are gauge fixed. This is discussed in the beginning of section~\\ref{s:superconf}.  The formulation of the theory in \\cite{Kallosh:2000ve} has several advantages. For example, it simultaneously incorporates two different formulations of phenomenological supergravity depending on the gauge fixing of the local chiral $U(1)$-symmetry. The first  formulation, in a \\Ka\\-covariant gauge, which is more standard, corresponds to \\cite{Cremmer:1983en}. The other one, in a new gauge where the conformon is real, is closer to \\cite{Wess:1992cp,Binetruy:2000zx}, and has a non-singular dependence on the superpotential $W$.   The new formulation \\cite{Kallosh:2000ve}  allowed  to give a detailed explanation of the superconformal origin of FI-terms by including gauge transformations of the conformon field as first suggested in \\cite{Ferrara:1983dh}. The conformon field $Y$ is one of the extra superfields of the superconformal version of the theory, which gets fixed to remove the local dilatation and local chiral $U(1)$-symmetry. However, at the superconformal level before the gauge fixing  such a field may participate in gauge transformations, which turn out to provide the FI terms: \\begin{equation} \\delta _\\alpha \\Yrho =\\rmi  {g\\xi_\\alpha\\over 3M_P^2} \\Yrho\\,\\,,\\qquad \\delta _\\alpha z_i= \\eta_{\\alpha i}(z)\\,, \\label{delarhoz} \\end{equation} where $g$ is the gauge coupling constant. When $\\xi_\\alpha$ are some real constant terms in some of the $U(1)$'s, they turned out to be constant FI terms $\\xi_\\alpha$ in the related $U(1)$ in the supergravity theory.  All corrections to the supergravity action in such a case can be deduced from the original superconformal action.   The scaling of fields that allows a suitable rigid limit is discussed in section~\\ref{ss:rigidlimit}. It allows us to present the action and transformation rules of supergravity with chiral and vector multiplets in a simple form in section~\\ref{ss:simplAction}. The different contributions to $R$-symmetry and $D$-terms and the implications for the superpotential are collected in section~\\ref{ss:summRFIW}. There the difference between field-dependent and constant FI terms is clearly exhibited, but also it is shown how terms can be reinterpreted by performing K{\\\"a}hler transformations. The final part of section~\\ref{s:superconf} shows how effective constant FI terms may result from field-dependent FI terms by replacing a chiral multiplet by its constant value without breaking local supersymmetry. This procedure is obvious in rigid supersymmetry, but cannot be done in supergravity in general. We treat in section~\\ref{ss:susyRemoveCh} a case with a K{\\\"a}hler potential that splits in two parts.  Section~\\ref{s:DtermcstFI} shows how the $D$-term inflation is modified by this connection to $R$-symmetry. We present here  corrections to the action  proportional to $g\\xi/ M_P^2$, which are required for consistency of the most  general $N=1$ supergravity with FI terms. We give an explicit example of such corrections for the supergravity theory describing $D$-term inflation \\cite{Binetruy:1996xj,Halyo:1996pp}. We also show that such corrections vanish in the limit of rigid supersymmetry. Such corrections to the $D$-term inflation model have not been exposed so far. Therefore, we will revisit this model and present a corrected form of it as an example of the general supergravity with constant FI terms.  The supergravity theory with constant FI terms has recently been shown to have $D$-term string solutions with unbroken supersymmetry in~\\cite{Dvali:2003zh}. A short summary of the $D$-term string solution is presented in section~\\ref{ss:Dstringconfig}. In section~\\ref{s:zeromodesD} we study these solutions and their zero modes in an extended model in which the $D$-term string is coupled to an arbitrary number of chiral superfields. We present the fermionic zero modes coming from these superfields and verify explicitly, as well as deduce from the superalgebra that the bosonic ones are absent, despite the unbroken supersymmetry.   Then, we turn to field-dependent FI terms in section~\\ref{ss:FIfromanomU1}. We revisit the issues of the $D$-term inflation model for the case of string theory inspired anomalous FI terms. We discuss the the relation with the anomaly cancellation by the Green-Schwarz mechanism.  We show that the true derivation of such cosmological models from string theory requires to find a stabilization mechanism for the dilaton and/or volume moduli. We describe some preliminary efforts in this direction existing in the literature, which may eventually lead to a stringy version of $D$-term inflation. Furthermore we discuss  the scales of $F$ and $D$-terms.  The cosmological implications of $D$-term strings for D-brane systems and D-brane inflation are discussed in section~\\ref{s:cosmoApplicD}. Especially the stability is discussed first in the supergravity formulation (section~\\ref{ss:DstabilSG}). Then the relation of the instabilities of type $II$ D-strings and the $D$-term strings in supergravity is exhibited in section~\\ref{ss:DDinstab}.  The connection of the supergravity $D$-term description to the D-brane-anti-brane configuration is further deepened in section~\\ref{ss:FIDfromD} by mapping the the moduli stabilization issues in the two cases.  Appendix~\\ref{app:decompsusy} on the residual supersymmetry algebra of the $D$-term string configuration and appendix~\\ref{app:BosFerModes} that repeats the the relation between fermionic and bosonic modes are useful for section~\\ref{s:zeromodesD}. ",
        "conclusions": "In this paper we have clarified the status of constant FI terms $\\xi$ in $N=1$, $d=4$ supergravity in general and in examples. Their presence shows up in covariant derivatives of all fermions and in the supersymmetry transformation laws, since the relevant local $U(1)$ symmetry is a gauged $R$-symmetry. These new couplings proportional to ${g\\xi/ M_P^2}$  lead to gauge and gravitational anomalies. Under certain conditions it is possible to cancel the gauge anomalies.  One of the important restrictions on supergravity with constant FI terms is the following: the superpotential $W$  has to  transform under $U(1)$ gauge symmetry, $\\delta W= -\\rmi{g\\xi\\over M^2_{P}} W$, otherwise the constant FI term $\\xi$ has to vanish. This requirement is consistent with the fact that in the gauge theory at $M_{P}\\rightarrow \\infty$ the potential is $U(1)$ invariant. However, we consider local supersymmetry of the classical supergravity action, in which terms of the order ${g\\xi/ M_P^2}$ are all taken into account.   In the  example of a $D$-term inflation \\cite{Binetruy:1996xj}, it is possible to generalize the original model with rigid supersymmetry to exact local supersymmetry. A gauge theory potential of the $D$-term model $W=\\phi_0 \\phi_+\\phi_-$ is neutral under $U(1)$ symmetry in  gauge theory with constant FI terms. In this paper we  have found how to promote  this model to the supergravity level with constant FI terms: we required that the total charge of $\\phi_+$ and $\\phi_-$ fields does not vanish but is equal to $-{\\xi/ M_P^2}$. The gauge theory anomaly from gravitino, gaugino, original chiral multiplets $\\phi_0$, $\\phi_+$ and $\\phi_-$ and additional 3 chiral multiplets  can be cancelled.  In string theory there are no known examples of constant FI terms $\\xi$. Only moduli-dependent $D$-terms are available \\cite{Dine:1987xk,Burgess:2003ic}. In absence of constant FI terms, the rules of local supersymmetry require the superpotential to be invariant under the $U(1)$-gauged symmetry (for invariant K{\\\"a}hler potential). In such theories, where the cancellation of gauge anomaly is due to the Green-Schwarz mechanism with the shift of the axion and the coupling $aFF^*$, one has to stabilize the scalar partner of the axion to get the effective supergravity with constant FI terms. Some efforts in this direction have been made recently in \\cite{Kachru:2003aw,Kachru:2003sx,Burgess:2003ic,Hsu:2003cy,Firouzjahi:2003zy,Angelantonj:2003zx,Koyama:2003yc}. It is possible that a stringy version of the D-brane inflation \\cite{Dvali:1998pa} with improvement with respect to volume and dilaton stabilization will be derived in the future and that the problems with inflation in string theory, pointed out in \\cite{Kachru:2003sx}, will be resolved. This kind of string cosmology program requires a better understanding of the structure of 3+1 dimensional $N=1$ supergravity with constant FI terms as well as the one with field-dependent $D$-terms. This paper has clarified the properties of such theories.  In this paper we have investigated another interesting role that FI terms can play in string theory. This role is based on the recently-suggested \\cite{Dvali:2003zh} equivalence between the $4d$ supergravity $D$-term strings and D-branes of type $II$ string theory. According to it, brane-anti-brane systems in an effective $4d$ theory can be viewed as gauge theories with non-zero FI term, in which the axion shifting under the $U(1)$-symmetry comes from the RR sector. The tachyonic instability of the brane-anti-brane system is seen as the instability triggered by the FI-term. Thus, many important properties of D$-\\bar{\\mbox{D}}$ systems can be understood in the light of $4d$ supergravity with FI $D$-terms. Certain aspects of stability of some string compactifications with branes and anti-branes can be understood in the language of supergravity vacua with non-zero $D$-terms. For instance, the shift of the axion under the gauge $U(1)$ symmetry gives selection rules for the possible invariants of the stabilizing superpotential.  On the cosmic string front, we have provided some additional consistency checks of the conjectured correspondence \\cite{Dvali:2003zh} between $D$-term-strings and D-branes, by mapping the instabilities of the two.  Finally, we have studied in non-minimal model the zero mode content on the BPS $D$-term cosmic string solutions of $N=1$ supergravity with constant FI terms \\cite{Dvali:2003zh}. We have discovered a puzzling property that the numbers of bosonic and fermionic zero modes can be {\\it arbitrarily} different. We have explained this puzzling behaviour by unusual properties of unbroken supersymmetry.    \\subsection*"
    },
    "0402/astro-ph0402584_arXiv.txt": {
        "abstract": "The increase in luminosity with time of a main sequence star eventually can lead to substantial evaporation of the oceans  on an orbiting terrestrial planet.  Subsequently, the gas-phase H$_{2}$O in the planet's upper atmosphere can be photodissociated by stellar ultraviolet and the resulting atomic hydrogen  then may be lost  in a wind. This gaseous envelope may pass in front of the host star and produce transient, detectable ultraviolet absorption in the Lyman lines in systems older than 1 Gyr. ",
        "introduction": "The possible existence  of life on other planets is of central interest in  modern astronomy.  A standard working hypothesis is that liquid water is required for life, and here we describe an observational signature of evolved oceans on extra-solar terrestrial planets.  The occultation of HD 209458 by a Jovian planet has demonstrated that it is possible to study the atmospheres of extra-solar planets with absorption line spectroscopy during a transit (Charbonneau et al. 2000, 2002).  The Lyman ${\\alpha}$  spectrum of HD 209458 observed during its companion planet's transit can be explained as an outflow of hydrogen from this planet (Vidal-Madjar et al. 2003. Liang et al. 2003) that also has entrained other  gases (Vidal-Madjar et al. 2004). Below, we describe an extension of this technique to study gaseous winds from terrestrial planets.  To date, only Jovian mass extra-solar planets have been detected. The discovery of terrestrial planets by occultations of their host stars is the goal of future space missions such as Kepler (see, for example, Jenkins 2002).  In addition to transits discovered by photometry, it may be possible to perform  transient absorption line spectroscopy to study  gaseous outflows from terrestrial planets, and we suggest that atomic hydrogen originating from evolved oceans might be detectable by this method. It seems likely that Venus once had oceans that were lost through a wind  (Watson, Donahue \\& Walker 1981, Kasting \\& Pollock 1983), and computing the structure and composition of a gaseous outflow  from an analog to Venus in orbit around a  star with an age less than 1 Gyr is a goal of models being developed  by Parkinson et al. (2003). However, as described below, the detection of an outflow from an Earth-like planet orbiting a star older than 1 Gyr is more promising.   Although not necessarily true in the past (see Sackmann \\& Boothroyd 2003), currently, the Sun's luminosity is increasing with time  (see Girardi et al. 2000). Eventually, if significantly vaporized,  H$_{2}$O can dominate the composition of the Earth's atmosphere since the total mass of the oceans is 1.4 ${\\times}$ 10$^{24}$ g while the total mass of the current atmosphere is 5.1 ${\\times}$ 10$^{21}$ g (Schubert \\& Walterscheid 2000). In approximately 1 Gyr, the Earth will be sufficiently warm that a ``moist\" greenhouse may occur and the temperature and composition profiles of the atmosphere will evolve so that the fraction of H$_{2}$O in the upper atmosphere is increased (Kasting 1988).   When the Sun's luminosity is 1.4 times its current value, or when its age is 8 Gyr (see Girardi et al. 2000), then depending upon the effects of clouds, a runaway greenhouse may occur and the oceans will be totally vaporized (Kasting 1988).    As described by models for the early atmosphere of Venus, in the relatively water-rich upper atmosphere of the future, photodissociation of H$_{2}$O into OH and H (Ip 1983, Wu \\& Chen 1993) will be a major source of ultraviolet opacity. Subsequently,  the resulting atomic hydrogen atoms may escape from the Earth in a wind.    Below, we extend this scenario for the Earth's future evolution to extra-solar terrestrial planets, and we argue that in some circumstances, the wind of atomic hydrogen  from vaporizing oceans may produce observable absorption lines.  The Earth now is losing about 7 ${\\times}$ 10$^{26}$ H atoms s$^{-1}$ (see, for example, Pierrard 2003);  we propose that this rate may increase by a factor of ${\\sim}$10$^{3}$ in the future. ",
        "conclusions": "The oceans on a terrestrial planet may store the bulk of its hydrogen. When the host star's luminosity increases enough so that the oceans are substantially evaporated,  a wind of atomic hydrogen from the planet could be strong enough to produce detectable transient  Lyman absorption lines in the star's ultraviolet spectrum.  Systems older than 1 Gyr are particularly promising candidates to exhibit  this signature of  terrestrial planets with evolved oceans.                The referee made insightful and helpful comments. This work has been partly supported by NASA."
    },
    "0402/astro-ph0402298_arXiv.txt": {
        "abstract": "We have obtained high-resolution, high S/N near-UV-blue spectra of 22 very metal-poor stars ([Fe/H] $<-2.5$) with Subaru/HDS, and measured the abundances of elements from C to Th. The metallicity range of the observed stars is $-3.2 <$ [Fe/H] $< -2.4$. As found by previous studies, the star-to-star scatter in the measured abundances of neutron-capture elements in these stars is very large, much greater than could be assigned to observational errors, and in comparison with the relatively small scatter in the $\\alpha$- and iron-peak elements. In spite of the large scatter in the ratios of the neutron-capture elements relative to iron, the abundance {\\it patterns} of heavy neutron-capture elements ($56\\leq Z\\lesssim 72$) are quite similar within our sample stars. The Ba/Eu ratios in the 11 very metal-poor stars in our sample in which both elements have been detected are nearly equal to that of the solar system r-process component. Moreover, the abundance patterns of the heavy neutron-capture elements (56 $\\leq Z \\leq$ 70) in seven objects with clear enhancements of the neutron-capture elements are similar to that of the solar system r-process component. These results prove that heavy neutron-capture elements in these objects are primarily synthesized by the r-process.  In contrast, the abundance ratios of the light neutron-capture elements (38 $\\leq$ {\\it Z} $\\leq$ 46) relative to the heavier ones (56 $\\leq$ {\\it Z} $\\leq$ 70) exhibit a large dispersion. Our inspection of the correlation between Sr and Ba abundances in very metal-poor stars reveals that the dispersion of the Sr abundances clearly decreases with increasing Ba abundance. This trend is naturally explained by hypothesizing the existence of two processes, one that produces Sr without Ba, and another that produces Sr and Ba in similar proportions. This result should provide a strong constraint on the origin of the light neutron-capture elements at low metallicity.  We have identified a new highly r-process element-enhanced, metal-poor star, CS~22183--031, a giant with [Fe/H] = $-2.93$ and [Eu/Fe] = +1.2. We also identified a new, moderately r-process-enhanced, metal-poor star, CS~30306--132, a giant with [Fe/H] = $-2.42$ and [Eu/Fe] = +0.85.  The abundance ratio of the radioactive element Th ({\\it Z} = 90) relative to the stable rare-earth elements (e.g., Eu) in very metal-poor stars has been used as a cosmochronometer by a number of previous authors. Thorium is detected in seven stars in our sample, including four objects for which the detection of Th has already been reported. New detections of thorium have been made for the stars HD~6268, HD~110184, and CS~30306--132. The Th/Eu abundance ratios (log(Th/Eu)), are distributed over the range $-0.10$ to $-0.59$, with typical errors of 0.10 to 0.15 dex. In particular, the ratios in two stars, CS~31082--001 and CS~30306--132, are significantly higher than the ratio in the well-studied object CS~22892--052 and those of other moderately r-process-enhanced metal-poor stars previously reported. Since these very metal-poor stars are believed to be formed in the early Galaxy, this result suggests that the abundance ratios between Th and stable rare-earth elements such as Eu, both of which are presumably produced by r-process nucleosynthesis, may exhibit real star-to-star scatter, with implications for (a) the astrophysical sites of the r-process, and (b) the use of Th/Eu as a cosmochronometer. ",
        "introduction": "The very metal-poor stars, presently found in the halo of the Galaxy, are believed to have formed at the earliest times, shortly after it became possible for the Universe to make stars with sufficiently long main-sequence lifetimes (i.e., masses $\\le 0.8 M_\\odot$) to survive for $\\sim 14$ Gyr. The chemical compositions of these stars are thus expected to reflect a quite small number of nucleosynthesis processes, possibly as small as one, while the compositions of more metal-rich stars like the Sun reflect the cumulative (hence quite complex) results of the various processes that have been in operation during the entire history of Galactic chemical evolution. Recent abundance analyses for extremely metal-poor stars have provided quite valuable information on the individual nucleosynthesis processes involved \\citep[e.g.,][]{mcwilliam95b,ryan96, burris00, cayrel03}. In particular, detailed abundance studies of heavy metals in such stars have led to important progress in the understanding of the origin of neutron-capture elements in the early Galaxy.  Sneden et al. (1996, 2000, 2003) have studied the abundances of the extremely metal-poor star CS~22892--052, the first example of a growing class of metal-poor stars that exhibit very large excesses of r-process elements relative to iron ([r-process/Fe] $> +1.0$). Even from the first analysis, it was apparent that the relative abundance pattern of the heavy neutron-capture elements (56 $\\leq$ {\\it Z} $\\leq$ 76) in this star was identical (within observational errors) to that of the (inferred) solar system r-process component, which has been strengthened as more and better data have been acquired. This striking similarity may be surprising, considering that the abundance pattern of the Solar System has certainly been influenced by the integrated yields from a variety of nucleosynthesis sites. It should be pointed out that, in contrast, the abundance pattern of the lighter neutron-capture elements ($38 \\leq Z \\leq 48$) in CS~22892--052 exhibit clear deviations from that of the solar-system r-process component \\citep{sneden00,sneden03}.  Recent analyses of a small number of additional very metal-poor stars with moderate excesses of r-process elements ($0.5 \\le $ [r-process/Fe] $ < +1.0$) \\citep[e.g., ][]{westin00,johnson01,cowan02} have obtained similar results, that is, the abundance pattern of heavy neutron-capture elements ($56\\leq Z\\lesssim 72$) for each star is quite similar to that of the r-process component in solar-system material. These results imply that the neutron-capture elements in these very metal-poor stars were produced by the r-process, which produces quite similar abundance patterns at least for the range of $Z=56 \\sim 72$ (Schatz et al. 2002; Wanajo et al. 2002; Otsuki, Mathews, \\& Kajino 2003). By way of contrast, the abundance patterns of the lighter neutron-capture elements ($Z < 56$) are significantly different from that of the r-process component in the Solar System. This observational fact is interpreted as a result of the existence of (at least) two distinct classes of r-process events \\citep[e.g., ][ and references therein]{truran02}.  The search for, and subsequent detailed abundance studies of, r-process-enhanced, very metal-poor stars enable one to make estimates of the ages of these objects\\footnote{More appropriately, the time interval that has passed since the production of the r-process elements by the progenitor object(s) of these stars. For convenience, we will use the term ``age'' to refer to this time interval.} by use of the radioactive species that can be identified in them. The long-lived radioactive r-process elements, in particular Th and U, whose half lives (14 Gyr for $^{232}$Th; 4.5 Gyr for $^{238}$U) are comparable to or shorter than the cosmic age ($\\sim$ 15 Gyr), provide a powerful tool for determination of a lower limit on the age of the Galaxy, and hence of the universe. If Th and/or U is detected in a very metal-poor star, we can estimate the lapse of time from the era of the nucleosynthesis process that created these elements to the present. This is accomplished by measuring the present abundance ratios, either of U/Th, when they are both detected (which only applies to two stars thus far), or by measurement of the U or Th abundance ratio as compared to stable r-process elements, such as Eu, and comparing them to predictions of the initial production ratios from models (both site-independent, e.g., Schatz et al. 2002, and site-dependent, e.g., Wanajo et al. 2002, 2003; Otsuki, Mathews, \\& Kajino 2003).  The cosmochronometry technique (in stars) was pioneered by Butcher (1987), who unfortunately only had roughly solar-abundance stars to work with, hence he had to contend with the difficulties arising from complex continua and line-blending problems, which are not so severe for very metal-poor stars. One clear advantage of the cosmochronometric method is that the resulting age estimate is free from a host of uncertainties encountered by other estimates of the cosmic age, such as calibration of distance scales, detailed understanding of stellar evolution, etc. Furthermore, since it is thought that the large overabundances of r-process elements are likely to have been associated with a single production event, quite likely a Type II supernova explosion in the early Galaxy, one is not forced to model the entire complex history of Galactic chemical evolution in order to estimate the age. The difficulties in this method, aside from the rarity of the very metal-poor stars with detectable Th and U (presently estimated to be no more than $\\sim 3$\\% of giants with [Fe/H] $< -2.5$\\footnote{We use the usual notation [A/B]$\\equiv \\rm{log}_{10} (N_A/N_B) _*-\\rm{log}_{10}(N_A/N_B)_\\odot$ and log$\\epsilon(\\rm A)\\equiv \\rm{log}_{10}(N_A/N_H)+12.0$, for elements A and B. Also, the term ``metallicity'' will be assumed here to be equivalent to the stellar [Fe/H] value.}; Beers, private communication), arise from the need to accurately predict the initial production ratios, which in turn depend on having a detailed understanding of the nucleosynthesis pathways, nuclear mass models, and cross-sections for species that have not been measured adequately at present (see Schatz et al. 2002 for details).  In the seminal study by Fran\\c{c}ois, Spite, \\& Spite (1993), Th abundances were reported for the first time for stars with metallicities as low as [Fe/H] $\\simeq$ --2.6. Sneden et al. (1996) first succeeded in obtaining a clear detection of Th, along with other r-process elements, in CS~22892--052 (see above), a relatively bright, extremely metal-poor giant ([Fe/H] = --3.1) discovered during the HK survey of Beers and collaborators (see Beers et al. 1992). From these authors' measurement of the Th/Eu ratio, the age of this star was estimated to be 15.2 $\\pm$ 3.7 Gyr. In their analysis, they assumed the initial abundance ratio between Th and stable elements to be the same as that of the calculated initial solar-system abundance ratio. Subsequently, Westin et al. (2000) estimated the age of another r-process-enhanced metal-poor star, HD~115444, to be 14.2 $\\pm$ 4 Gyr from the Th/Eu ratio. Johnson \\& Bolte (2001) also measured Th abundances for five very metal-poor stars. Most recently, Cowan et al. (2002) studied the moderately r-process-enhanced, very metal-poor star, BD+17$\\degr$3248. This star, with [Fe/H] = --2.1, and [Eu/Fe] = +0.9, has the distinction of being among the most metal-rich stars in which the r-process-enhancement phenomenon has been observed. These authors, based on an average of a number of chronometer pairs involving U and Th, obtained an age estimate of $13.8 \\pm 4$ Gyr. Sneden et al. (2003) have assembled a definitive high-S/N set of spectra for CS~22892--052, drawing on space-based and ground-based observations. Their analysis has led to a revision of the Th/Eu age estimate for this star, to $12.8 \\pm 3$ Gyr. These results are consistent with the very recent estimate of the age of the Universe by WMAP (Bennett et al. 2003) to within the reported errors.  Recently, Cayrel et al. (2001) and Hill et al. (2002) reported the detection of U, as well as Th, in a high-quality VLT/UVES spectrum of the extremely metal-poor star CS~31082--001, a bright [Fe/H] = --2.9 giant discovered in the HK survey. They also found that the abundance pattern of neutron-capture elements with $56 \\leq Z \\leq 72$ closely mimics that of the r-process component in the Solar System. However, it is of interest that the situation seems to be different for nuclei near the heaviest elements. Th ($ Z = 90$) and U ($Z = 92$) show larger deviations from their expected levels, if we adopt $\\sim 14$ Gyr as the age of this object, as was estimated from the U/Th ratio by \\citet{hill02}. This result suggests that either CS~31082--001 is peculiar in some way that has affected the surface abundances of the actinides, or that the production ratios of Th and U, as compared to the stable elements (e.g., Eu) exhibit some real dispersion amongst very metal-poor stars. Further observational work is clearly required to answer this important question.  The primary purpose of this study is to obtain measurements of the abundance patterns of neutron-capture elements for very metal-poor stars, and to examine the age estimation of stars using the Th/Eu chronometer. In Honda et al. (2003; hereafter Paper I), we reported equivalent width measurements of absorption lines in high-quality spectra of 22 very metal-poor stars obtained with the Subaru Telescope High Dispersion Spectrograph \\citep[HDS, ][]{nogu02}. This sample is as large as those in previous studies by McWilliam et al. (1995), Ryan, Norris, \\& Beers (1996), and Johnson \\& Bolte (2001), and the quality of the spectra is quite high ($S/N > 100$ per resolution element) in the blue region. In particular, the sample was selected to include as many objects with excesses of neutron-capture elements as possible, excluding the objects affected by the s-process, to investigate the nature of r-process nucleosynthesis in the early Galaxy. In this paper, we present the abundance analyses for the neutron-capture elements, and discuss the observed abundance distributions in very metal-poor stars with excesses of r-process elements. In \\S 2 we describe the determination of atmospheric parameters for our program stars. The abundance analysis is described in \\S 3, where error estimates in our derived abundances, and comparisons with previous studies, are also discussed. In \\S 4 we discuss the patterns in the abundance ratios of neutron-capture elements produced by r-process nucleosynthesis, based on the results of our analysis. In this same section, the abundance patterns of other elements, including the actinide Th, are considered, as is the suitability of Th-based chronometers for age estimates. ",
        "conclusions": "The present analysis has shown that our objects indeed possess very low metallicities ($-3.2 <$ [Fe/H] $< -2.4$). The iron abundances derived by the present analysis are similar to the values which were estimated by the previous studies from lower dispersion spectroscopy (Beers et al. 1992; Bonifacio et al. 2000; Allende-Prieto et al. 2000). One exception is the star CS~22952--015, whose metallicity was estimated to be [Fe/H] $= -3.50$ by Beers et al. (1992), while our estimate is rather higher, [Fe/H] $= -2.94$. Another is CS~30306--132, whose metallicity was estimated to be [Fe/H]$=-3.1$ from lower-dispersion spectroscopy by Beers et al. (2003, in preparation), while our result is [Fe/H]$=-2.44$. This discrepancy may be due to the fact that this star has strong CH molecular bands.  The metallicity range of the stars in our sample should be kept in mind, as it is known that a large scatter in the abundance ratios in many neutron-capture elements appears when the metallicity drops to [Fe/H] = --2.5 or below \\citep[e.g., ][]{mcwilliam95b}. In this section, we discuss the relative abundances of the neutron-capture elements, the dispersion of the observed abundance ratios, and the origin of these elements (subsection 4.1). We discuss in detail the abundance pattern of neutron-capture elements for the seven stars in our sample for which Th is detected, and the impact of our new results on the cosmochronology technique based on Th (subsection 4.2). The abundances of the $\\alpha$- and iron-peak elements are discussed only for comparison purposes. Details on the behavior of these elements will be presented separately in a future paper in this series (Honda et al., in preparation).   \\subsection{Relative Abundances of the Neutron-Capture Elements}   \\subsubsection{The Heavy Neutron-Capture Elements ($56 \\leq Z \\leq 76$)}  Figure \\ref{fig:ba} shows the values of [Ba/Fe] as a function of [Fe/H] for all of the stars in our sample (filled circles with error bars), as well as the results reported by previous authors for comparison (open circles). Ba abundances have been reported for stars with lower metallicity than the abundances of other neutron-capture elements such as Eu. This is because Ba has strong Ba {\\small II} resonance lines (4554 and 4934~{\\AA}), which remain detectable even as the overall level of metallicity decreases. Previous studies have shown that [Ba/Fe] drops below the solar ratio, on average (e.g., Ryan et al. 1996; McWilliam 1998), at the lowest metallicities. This trend is thought to arise because of a change in the primary nucleosynthesis sources for Ba between stars with [Fe/H] $\\la -2.5$ and those with [Fe/H] $\\ga -2.5$. Ba in the more metal-rich stars is believed to originate primarily from the (main) s-process in low-mass or intermediate-mass stars (1-8 $M_\\odot$) with a comparatively small contribution by the r-process. The production of Ba at the lowest metallicities, however, is most likely due to the r-process alone, occurring in Type II supernova explosions of massive and hence rapidly-evolving stars. The contribution of the s-process to Ba in the early Galaxy is small, because the time-scale for the evolution of lower mass stars is long, and the ejecta from these stars contribute only to stars which formed somewhat later, with [Fe/H] $\\ga -2$~\\citep[e.g., ][ and references therein]{truran02}.  Our results, shown in Figure \\ref{fig:ba}, confirm the existence of an extremely large scatter (a factor of $\\sim$1000 over the entire range) in the Ba abundances of the most metal-poor stars, even larger than reported in previous studies. For instance, the standard deviation of the [Ba/Fe] values of our sample is 0.82~dex, while that found by \\citet{mcwilliam98} for 24 objects with [Fe/H] $<-2.5$ is 0.59~dex. One clear reason for this larger dispersion is that the detection limit for Ba lines in our work is lower than in previous programs, thanks to the high quality of the Subaru/HDS spectra. The other reason is likely due to our selection of candidate neutron-capture-enhanced stars, as mentioned in Paper I. In spite of this selection bias, the large dispersion in Ba abundances that exists in the metallicity range of $-3.0 \\la$ [Fe/H] $\\la -2.5$, and the possibly decreased scatter at lower iron abundances, provides an important clue to the sites and mechanisms of astrophysical neutron-capture processes \\citep{wasserburg00}.  Roughly 95\\% of the Eu in solar-system material is believed to be produced by the r-process \\citep[e.g.,][]{arlandini99,burris00}, hence this element is particularly suitable for investigating the characteristics of the r-process in the early Galaxy. However, measurements of Eu abundances in extremely metal-poor stars ([Fe/H] $\\sim -3.0$) are still quite limited, because there is no strong Eu line, unlike those of Sr and Ba. We have detected Eu in 11 objects in our sample; the derived [Eu/Fe] ratios for our sample are shown in Figure \\ref{fig:eu} as a function of [Fe/H] (filled circles), along with the results obtained by previous studies (open circles). Our results show a large scatter also in [Eu/Fe] in very metal-poor stars, as has been reported by previous authors \\citep[e.g.,][]{mcwilliam95b,burris00}.  The ratio [Ba/Eu] is useful for distinguishing the contributions of the r- and s-processes, because the expected ratios from the two processes are quite different. Figure \\ref{fig:baeu} shows [Ba/Eu] as a function of [Fe/H] for the 11 stars in our sample for which Eu has been detected. The dotted line indicates the value of [Ba/Eu] of the solar-system r-process component ([Ba/Eu] $= -0.69$, Arlandini et al. 1999), while the dashed line indicates that of the s-process component ([Ba/Eu] $= +1.15$, Arlandini et al. 1999). The [Ba/Eu] ratios exhibited by our stars are clearly associated with the r-process, rather than the s-process.  Similar results have been reported by previous authors. McWilliam (1998) found that the mean of the [Ba/Eu] ratios for 10 stars with [Fe/H] $\\leq$ --2.4 is --0.69, consistent with the value expected from pure r-process nucleosynthesis within the measurement uncertainties. Here we find that the ratio of [Ba/Eu] matches that of the solar-system r-process component for many stars at low metallicity. These results suggest that Ba, as well as Eu, is primarily produced by the r-process during the early history of the Galaxy, and that Ba can be used as a powerful tool to investigate the behavior of r-process elements in the early Galaxy, as it remains detectable in stars of metallicities that are far lower than those in which Eu is detectable.\\footnote{A number of stars with large excesses of s-process elements {\\it are} known to exist in this metallicity range \\citep[e.g, ][]{aoki02b}. The chemical compositions of these stars are interpreted as resulting from s-process nucleosynthesis in intermediate-mass, evolved stars followed by mass-transfer across the binary system. These stars do not exist in our sample, presumably because carbon-enhanced metal-poor stars have been specifically excluded in our sample selection (see Paper I).}. Though most of the stars shown in Figure \\ref{fig:eu} exhibit [Eu/Fe] $\\gtrsim0$, we expect that the [Eu/Fe] values of the stars with low Ba abundances are lower than zero, and in reality will, once detected, exhibit a similarly large dispersion as seen in [Ba/Fe]. This prediction should be confirmed by studies of Eu lines using higher quality spectra for objects with [Fe/H] $< -3$ (e.g., Ishimaru et al. 2004).  Large levels of scatter are also found in the abundance ratios of almost all neutron-capture elements with {\\it Z} $\\geq$ 56. Figure \\ref{fig:scatter} shows the average of the elemental abundances relative to iron [X/Fe], and the standard deviation as a measure of the scatter in the abundances, as a function of atomic number. The number of objects in which the species is detected depends upon the element: i.e., neutron-capture elements heavier than Ba are detected only in stars with excesses of neutron-capture elements. The dispersion of the abundances for elements which are detected in less than 12 objects is shown by the thin bars. Even though this limitation will make the dispersion of the abundances of heavy neutron-capture elements ($Z \\ga 57$) smaller, the dispersion of the abundances of neutron-capture elements is much larger than found for the $\\alpha$- and iron-peak elements.  The scatter in the abundances of heavy neutron-capture elements relative to iron found in these very metal-poor stars means that the nucleosynthesis site of iron-peak elements and the r-process elements are quite different, and that the mixing of the yields from the first supernovae into the ISM is incomplete in the early stages of the Galaxy. The large scatter of r-process elements appearing at [Fe/H] $\\sim -3$ should provide a constraint on the dominant site of r-process nucleosynthesis \\citep[e.g., ][]{ishimaru99,tsujimoto00}.  We have confirmed that CS~22892--052 and CS~31082--001 are extremely r-process-rich objects, and that HD~6268, HD~115444 and HD~186478 are moderately r-process-rich objects. In addition, we have found two new r-process-enhanced objects, CS~22183--031 and CS~30306-132. CS~22183--031 exhibits a large excess of r-process elements (e.g., [Eu/Fe] $\\gtrsim+1$). Unfortunately, we could not detect many lines because of the relatively low S/N in our spectrum of CS~22183-031. We are able to detect numerous lines in CS~30306-132, and found that this object has a moderate enhancement of r-process elements ([Eu/Fe] $\\sim$ +0.8). In \\S 4.2 we describe this object in detail.   \\subsubsection{The Light Neutron-Capture Elements ($38 \\leq Z \\leq 40$)}  Figure \\ref{fig:lightr} shows the abundance ratios of [Sr/Fe], [Y/Fe], and [Zr/Fe] as a function of [Fe/H]. The [Sr/Fe] ratios in our program stars are distributed over a very wide range, from $-1.7$ to $+0.5$, confirming the large dispersions in this ratio found by previous studies \\citep[e.g., ][]{mcwilliam95b,burris00}. The stars in our sample appear to exhibit a rather higher mean [Sr/Fe] ratio than those of previous studies, but this is likely because of the sample selection, as discussed in the previous subsection. The scatter that appears in the [Y/Fe] and [Zr/Fe] ratios is smaller than that in [Sr/Fe]. However, Y and Zr do not have such strong spectral lines as the \\ion{Sr}{2} resonance lines, hence these two elements are not detected in several objects in our sample. This may account for their smaller dispersion; clearly, this should be investigated by obtaining higher-quality spectra of stars that presently have only upper limits on their Y and Zr abundances.  In order to investigate the reason of the large dispersion in the Sr abundances, in Figure \\ref{fig:srbafe} we plot [Sr/Ba] as a function of [Fe/H]. As discussed previously by \\citet{mcwilliam98}, although the dispersion in [Sr/Ba] at very low metallicity is rather smaller than that of [Sr/Fe], the range is still almost 2~dex. This stands in stark contrast to the range of [Ba/Eu] exhibited by the stars in our sample, all of which have quite similar values. This result suggests that either (a) the process that contributed significant amounts of Sr in these metal-deficient stars did not yield similar amounts of Ba, or (b) the process that produced Ba at very low metallicity yielded a variety of Sr/Ba ratios.  To investigate the correlation between Sr and Ba, Figure~\\ref{fig:srbabafe} shows the ratio [Sr/Ba] for the stars in our sample as a function of [Ba/Fe]. Such a diagram was was also shown, for a different set of stars, by \\citet{truran02} and Sneden, Preston, \\& Cowan (2003). The sample of Sneden et al. (2003) includes s-process-rich stars, in which Pb is detected. By way of contrast, we have excluded stars known to exhibit large excesses of s-process elements and stars with [Fe/H] $>-2.5$ from this figure, in order to avoid possible contamination by s-process nucleosynthesis. (A possible contribution of the so-called weak s-process is discussed below.) One clear result found in this figure is that the dispersion in [Sr/Ba] decreases with increasing Ba abundance. This correlation is much clearer in our figure than in Figure 10 of \\citet{truran02}, presumably because their sample includes several rather metal-rich objects, which could well be affected by the contribution of s-process nucleosynthesis. Moreover, our new sample of stars has added a number of objects at the high end of the [Ba/Fe] range, hence this contributes to making the correlation in this range clearer.  Figure \\ref{fig:srbares} shows the abundances of Sr and Ba for very metal-poor stars. For the stars in common between our study and those of others, we have adopted the abundances derived in the present study. As a result, a total of 46 stars are shown in this figure. It is obvious that the dispersion in the Sr abundance decreases with increasing Ba abundance. To demonstrate this quantitatively, we divided the sample into three groups on the Ba abundance ($\\log \\epsilon$(Ba)$>-1$: 16 stars, $-1\\geq \\log \\epsilon$(Ba)$>-2$: 19 stars, and $-2\\geq \\log \\epsilon$(Ba): 11 stars), then measured the standard deviation of the Sr abundances for each group. The results are 0.25 dex, 0.38 dex, and 0.71~dex, respectively. Clearly, the dispersion increases with decreasing Ba abundances. While the standard deviation in the range of $\\log \\epsilon$(Ba)$>-1$ is of a similar level as the typical observational errors, in the range of $\\log \\epsilon$(Ba)$\\leq -2$ the dispersion is significantly larger than the measurement errors.  We also plot in Figure \\ref{fig:srbares} the values of Sr and Ba in solar-system material (Grevesse, Noels, \\& Sauval 1996, the dotted circle), as well as the solar-system r-process component (the filled square). Unfortunately, estimation of the s-process component for Sr (knowledge of which is required in order to obtain the r-process residual value) is quite uncertain, because of possible contamination from the so-called weak component of the s-process \\citep{kappeler89}. For the r-process component of Sr, we adopted the average of the values derived by \\citet{kappeler89} and by Arlandini et al. (1999, the estimate from their classical model), and show the difference by an error bar.  A simple model can be constructed for the enrichment of Sr and Ba, assuming initial abundances of Sr and Ba ($\\epsilon_{0}$(Sr) and $\\epsilon_{0}$(Ba)) and a constant Sr/Ba ratio ((Sr/Ba)$_{\\rm r}$) in the yields of the r-process, i.e.,  $\\epsilon$(Sr)=$\\epsilon_{0}$(Sr)$ + $(Sr/Ba)$_{\\rm r} x$  $\\epsilon$(Ba)=$\\epsilon_{0}$(Ba)$ + x$ .  The two solid lines in Figure \\ref{fig:srbares} show cases for different initial Sr abundances: $\\epsilon_{0}$(Sr) $= 3 \\times 10^{-3}$ and 5, respectively. $\\epsilon_{0}$(Ba)$ = 3 \\times 10^{-3}$ and (Sr/Ba)$_{\\rm r}=1.5$ is assumed for both cases. A glance at this figure shows that the observational data fill the range between the two lines. Hence, the distribution of the observed Sr and Ba abundances are quite naturally explained by the simple assumptions of a large dispersion of Sr abundances at $\\log \\epsilon$(Ba)$\\sim -2.5$ and enrichment of Ba and Sr with a constant Sr/Ba ratio. Moreover, if we extend the line representing the enrichment of Sr and Ba, the Sr and Ba abundance of the r-process component in solar-system material is also explained.  A similar scenario for Sr and Ba enrichment has already been proposed by previous studies \\citep[e.g., ][]{sneden00,truran02}. In particular, \\citet{truran02} selected several metal-deficient stars with high and low Ba abundances, and concluded that Ba-poor stars show high Sr/Ba ratios, while the Sr/Ba ratios of Ba-rich stars are similar to that of the r-process component in the Solar System. These authors suggested the existence of two processes (sites) that produce Sr. Our interpretation for the Sr and Ba abundance distributions is essentially the same as theirs. However, the enrichment of Sr and Ba is seen much more clearly in the present study, as the result of adding new measurements and excluding stars of higher metallicity.  The enrichment of Ba in the metallicity range around [Fe/H]$\\sim -3$ is sometimes referred to as the main r-process, which produces heavy (A $>$ 130) neutron-capture elements \\citep{wasserburg00}. Our interpretation for the variation of the Sr and Ba abundances results in the Sr/Ba ratio produced by this process, (Sr/Ba)$_{\\rm r}$, to be about 1.5. The absence of stars with Sr/Ba $<1$ in Figure \\ref{fig:srbares} means that (Sr/Ba)$_{\\rm r}$ is not significantly smaller than unity. This value is much higher than the Sr/Ba predicted by recent models of r-process nucleosynthesis (e.g., $\\sim 0.03$: Otsuki, Mathews, \\& Kajino 2003). The value derived by these models should be, however, a lower limit of the Sr/Ba ratio produced by a single site, as these models deal with only one condition which can produce heavy neutron-capture elements like Ba. The yields in a real r-process site, such as a Type II supernova explosion, should be an integration of the results of a variety of conditions including the cases with low neutron-to-seed ratios, and produce light neutron-capture elements like Sr. If our interpretation for the correlation between Sr and Ba abundances is correct, this places a quite strong constraint on the Sr/Ba ratio produced by the main r-process.  In contrast, the nucleosynthesis process that is responsible for the production of light neutron-capture elements (e.g., Sr) without producing the heavier elements (e.g., Ba) is as yet unclear. One possibility is the existence of an independent nucleosynthesis process which dominantly produces lighter neutron-capture elements, probably prior to the 'main' r-process which provides heavier ones (e.g., Truran et al. 2002). For instance, the above observational results are qualitatively explained by the assumption that different mass ranges of supernovae are responsible for the light and heavy neutron-capture elements (e.g., Qian \\& Wasserburg 2000). Two different explosion mechanisms, i.e., prompt and delayed explosion, are proposed for low-mass and high-mass supernovae, respectively, and numerical simulations of r-process nucleosynthesis have been made \\citep[e.g., ][]{hillebrandt84,woosley94,meyer92,otsuki00, sumiyoshi01,wanajo03}. Possible r-process nucleosynthesis in neutron-star mergers may also play a role in the enrichment of the light neutron-capture elements \\citep[e.g.,][]{rosswog99, freiburghaus99}.  On the other hand, another explanation by r-process nucleosynthesis in a single site may be possible if incomplete mixing of supernova ejecta with interstellar matter is assumed. For instance, some neutrino-heated wind models predict significantly large overproduction of light neutron-capture elements in the early phase ($\\sim 1$ second) of the explosion prior to the r-process which produces the heavier elements (e.g., Woosley et al. 1994). Cameron (2001) suggested that r-process nucleosynthesis may take place in the jets associated with gamma-ray bursts, which would be a possible mechanism for production of an inhomogeneous r-process in a single event. Stars with high Sr and low Ba abundances may form from the interstellar medium polluted by the ejecta enriched in light neutron-capture elements, if the ejecta is not well mixed with the matter ejected in the later phase of the r-process.  An alternative explanation for the enrichment of Sr without Ba may be the so-called weak s-process, which proposes a neutron-capture process proposed to occur in core He-burning massive stars \\citep[e.g., ][]{hoffman01}. This process is expected to be unimportant in metal-deficient stars, because $^{22}$Ne is believed to be the neutron source. However, if this process {\\it can} operate, it may contribute to light neutron-capture elements in metal-deficient stars.   Further systematic studies of stars with low Ba abundances, especially of the abundance patterns for elements around Sr, will provide important information to constrain the nucleosynthesis process responsible for their production, and their ejection dynamics as well.   \\subsection{Distribution of Neutron-Capture Elements}  One important discovery in recent abundance studies of the neutron-capture elements in metal-deficient stars is that the abundance pattern of heavy neutron-capture elements are very similar, essentially an exact match, within observational errors, to that of the r-process component in the Solar System. To examine details of r-process nucleosynthesis at low metallicity, it is essential to detect as many elements as possible over the entire range of atomic numbers ({\\it Z} = 31 $\\sim$ 92). Such a study is possible for very metal-poor stars with overabundances of neutron-capture elements ([X/Fe] $\\ga$ +0.5). In such stars the absorption lines of neutron-capture elements are relatively strong, and the blending from lines of other lighter elements is comparatively weak, due to the overall low metallicity of the star.  We examine the abundance pattern of neutron-capture elements in detail for the seven very metal-poor stars in which the Th absorption line (4019 {\\AA}) was detected (see Figures~\\ref{fig:th1} and \\ref{fig:th2}). All of these stars exhibit enhancements of their neutron-capture elements. Among them, CS~31082--001 and CS~22892--052, which have already been studied by Hill et al. (2002) and Sneden et al. (2003), are the most extreme cases, showing large enhancements of the neutron-capture elements (e.g., [Eu/Fe]=+1.7 and +1.5, respectively). The star HD~110184 does not exhibit a remarkable enhancement of neutron-capture elements relative to iron, but our high S/N spectrum makes it possible to study Th and other r-process elements in this star as well.   \\subsubsection{Abundance Pattern for Elements with 56 $\\leq$ {\\it Z} $\\leq$ 70}  Abundance studies of r-process-enhanced, metal-poor stars to date have shown that the abundance patterns of the neutron-capture elements with 56 $\\leq$ {\\it Z} $\\leq$ 70 agree very well with that of the solar-system r-process pattern (e.g., Sneden et al. 2003). In order to investigate this phenomenon further, using our extended sample, we now compare the scaled abundance patterns of our stars with the abundance distributions in solar-system material.  For each of the seven stars with detectable r-process elements, we use the solar-system r-process abundance pattern as a template to compare the heavy-element abundances of the stars on a common scale (Figure \\ref{fig:pattern}).  We scale the elemental abundances of our objects to match the solar-system abundances for elements between Ba and Yb ($56\\leq Z\\leq 70$). The logarithmic values of the scaling factor ($\\log f$) for individual stars are listed in the columns labeled 's.s. r-process' in Table \\ref{tab:scale}. For this analysis, we adopt the r-process fraction in the Solar System given by Burris et al. (2000). The total solar-system abundances were taken from \\citet{grevesse96}.  Figure \\ref{fig:pattern} shows that the abundance patterns of the elements with 56 $\\leq Z \\leq$ 70 for these very metal-poor stars agree very well with that of the solar-system r-process component.  We also attempted to compare the abundances of our metal-poor stars with the solar-system s-process distribution in the same manner. Take CS~30306--132, for example (Figure~\\ref{fig:pattern132}). As expected from the [Ba/Eu] ratios of these objects, the agreement between the abundance patterns of our objects and the s-process pattern is poor. We also compare the abundances of our subset of metal-poor stars with r-process enhancements with the total elemental abundance distribution of the Solar System (Figure \\ref{fig:pattern132}). The logarithmic values of the scale factor are given in the column labeled 's.s. total' in Table \\ref{tab:scale}. The agreement is clearly better than that in the case for the abundance pattern of the s-process component alone. This is because the r-process fraction dominates in the total solar-system abundances of elements with $62\\lesssim Z\\lesssim 70$. The scaled abundances of Ba, La, and Ce in our stars, shown in Figures \\ref{fig:pattern}, are less than the solar-system total abundances of these elements. Presumably, this arises because the s-process contribution to the abundances of Ba, La, and Ce in the Solar System is larger than for other elements like Eu (Burris et al. 2000). From these comparisons, we conclude that the abundance patterns of heavy neutron-capture elements in our objects agree best with the r-process component in the Solar System.  In Table \\ref{tab:scale} we also provide the standard deviation (1-$\\sigma$) of the difference between the scaled abundances of each star and the total solar-system abundances, as well as its r-process component. These standard deviations can be taken as indicators of the level of agreement between the abundance pattern of each star and the pattern of the solar-system abundances. The average of the observational errors for elements with $56\\leq Z\\leq 70$ ($\\sigma_{\\rm obs}$) are also provided in Table 8. Comparisons of the observed 1-$\\sigma$ values with $\\sigma_{\\rm obs}$ indicates good agreement between the scaled abundance pattern of our objects with that of the solar-system r-process component. We note that the 1-$\\sigma$ value for HD~110184 is slightly larger than those for other objects, and than its $\\sigma_{\\rm obs}$. One reason for this result may be that the r-process enhancement of this star (e.g., [Eu/Fe] = 0.06) is small, and a small contribution by the s-process may affect the abundance pattern of this star.  The excellent agreement between the abundance pattern of the heavy neutron-capture elements in very metal-deficient stars and that of the solar-system r-process component has already been reported for several very metal-poor stars by previous studies \\citep[e.g., ][]{sneden96,sneden00, westin00,johnson01,hill02,cowan02, sneden03}, and is sometimes referred to as the 'universality' of r-process nucleosynthesis. The apparent universality is ascribed to the fact that the predicted abundance patterns of elements with $56\\leq Z\\leq 70$ appear to be rather insensitive to variations in the parameters of the current r-process models (e.g., the entropy/baryon ratio). The nucleosynthesis paths in this mass range (i.e., $A \\sim$ 150) are quite similar among the r-process models that are used to predict the production of the actinide nuclei, even though the abundances of the actinides show an apparent variation in some stars \\citep[e.g., ][]{wanajo02,otsuki03}.   \\subsubsection{The Radioactive Element Th and the Impact on Cosmochronometry}  As mentioned in \\S1, the actinide Th is heavier than the elements at the third abundance peak produced by the r-process (Os, Ir, Pb, etc.), and is a key element for the understanding of this nucleosynthesis process. Application of the abundances of this element to cosmochronometry has also been discussed extensively in recent studies of very metal-deficient stars.  The abundances of Th measured for the seven stars in our sample are presented in Figure \\ref{fig:pattern}. For the Th abundance, the solid line here shows the {\\it initial} abundance of this radioactive element, as estimated by Cowan et al. (1999), rather than the {\\it present} Th abundance, which is shown by the dashed line in this figure. Since these very metal-poor stars are believed to be born in the early Galaxy, the Th abundances of our sample are expected to be lower than the value shown by the dashed line, if we assume that the abundance patterns of heavy neutron-capture elements, including Th, produced by the r-process, is indeed universal. This is in fact found for the stars HD~110184, HD~115444, HD~186478, and CS~22892--052. However, the Th abundances of the other three stars are {\\it higher} than would be expected from this logic.  In order to investigate this issue more clearly, we show the abundance ratios between Th and the stable r-process element Eu ($\\log$(Th/Eu) = $\\log \\epsilon$(Th) $- \\log \\epsilon$(Eu)) in Figure \\ref{fig:theu}, where we plot our results, and the results of previous studies \\citep{westin00,sneden00,johnson01, hill02,cowan02}. The average and standard deviation of the values of our seven stars are $-0.40$ dex and 0.17~dex, respectively. The standard deviation is as large as, or slightly larger than, the typical observation errors, which were estimated from the root-sum-square of the random errors of Th and Eu abundances (error bars in Figure \\ref{fig:theu}).  We have found that CS~31082--001 and CS~30306--132 have clearly higher Th/Eu ratios than the well studied star CS~22892--052 (Figure \\ref{fig:theu}). In Figure \\ref{fig:th2}, the observed spectra around the \\ion{Th}{2} line are shown. The \\ion{Nd}{2} 4018.6~{\\AA} line exists blueward of the \\ion{Th}{2} line.  Since the atmospheric parameters in these giants are quite similar, the ratio of the line strengths between \\ion{Th}{2} and \\ion{Nd}{2} directly reflects the Th/Nd abundance ratio. In CS~22892--052, the \\ion{Th}{2} line is as strong as the \\ion{Nd}{2} line. The ratios of the line strengths in HD~6268 and HD~115444 are rather similar to that of CS~22892--052. In contrast, in CS~31082--001 the \\ion{Th}{2} line is significantly stronger than the \\ion{Nd}{2} line. In addition, in CS~30306--132, the \\ion{Th}{2} is clearly detected, while there is no evidence for the presence of the \\ion{Nd}{2} line. These results indicate that the Th/Nd ratios in CS~31082--001 and CS~30306--132 are significantly higher than in the other three stars. As discussed in detail in subsection 4.2.1, Nd in these objects is regarded as a product of the r-process, and can be taken as representative of the stable r-process elements. Therefore, the above inspection suggests the existence of some dispersion in the abundance ratios between Th and the other neutron-capture elements ($56 \\leq Z \\leq 70$).  A similar result was already reported for CS~31082--001 by \\citet{cayrel01} and \\citet{hill02}, as mentioned in \\S 1. Our present study shows that this object is not unique, but that there is at least one other similar object that shares the same ``actinide boost,'' CS~30306--132, though its enhancement factor of r-process elements is much smaller than that of CS~31082--001. We might expect that, in the near future, as additional r-process-enhanced metal-poor stars are identified, such behaviors will be seen in additional stars.  If the conventional Th/Eu chronometer (Cowan et al, 1999) is simply applied to the Th/Eu abundance ratios, the ages of a few stars, those exhibiting actinide boosts, are estimated to be shorter than that of the present age of the Sun. In particular, the derived ages of CS~31082--001 and CS~30306--132 are as low as zero. Even with a level of uncertainty of as much as 5 Gyr in the age estimation, the ages derived from the above method appear unrealistic. We note that the average of the Th/Eu ratios of our seven stars ($<\\log$(Th/Eu) $>=-0.42$) is quite similar to the value of solar-system material ($\\log$(Th/Eu) $=-0.46$ \\citep{grevesse96}). That is, the average of the ages derived from application of the Th/Eu chronometer is similar to the age of the Sun, and hence is also unrealistic.  The observations seem to suggest that some very metal-poor stars had {\\it higher} initial Th abundances than expected from the solar-system r-process abundance pattern. In other words, even though the abundance pattern of the elements with $56\\leq Z\\leq 70$ agrees with the abundance pattern of the solar system r-process component, the initial abundance ratios of the heaviest elements, like Th, to the lighter ones (e.g., Eu) are not necessarily the same as those expected from the r-process component in the Solar System. In order to apply the abundance ratios between Th and other stable elements as chronometers, estimates of the initial abundance ratios for these elements are essential, hence a deeper understanding of the r-process nucleosynthesis for wider mass ranges is necessary.  One possible alternative cosmochronometer is the U/Th ratio, recently applied by Cayrel et al. (2001) and Hill et al. (2002) for the extremely r-process-enhanced, metal-poor star CS~31082--001. Since $^{232}$Th and $^{238}$U are neighboring actinide nuclei, their production rate is expected to be quite similar. This justifies the assumption that the initial abundance ratio is the same as in the initial solar-system abundance ratio, as shown by recent theoretical calculations \\citep[e.g., ][]{wanajo02,otsuki03}. At present, U/Th is expected to be the best available chronometer.  We would like to point out that, even though we conclude from our analysis that there exists a real scatter in the abundance ratios between Th and other neutron-capture elements with $Z \\sim 60$, the level of this scatter is at most a factor of three, much smaller than the ratios between light ($Z \\sim 40$) and heavy ($Z \\sim 60$) neutron capture elements, which are as large as a factor of 10 (see Figure \\ref{fig:pattern}). Recent models of the r-process showed that the abundance ratios between the elements at the second and the third r-process peaks are quite sensitive to the parameters in the calculation \\citep[e.g.,][]{hoffman97,otsuki00,wanajo02,otsuki03}. The small dispersion of Th/Eu ratios found in our stars, as well as in other previously studied stars, should be an important constraint on modeling the r-process, as we are beginning to place limits on the possible range over within which the initial production ratio can fall. In order to derive a clear conclusion, it is required to systematically study the Th abundances for a larger sample of very metal-poor stars, based on high S/N spectra, which would enable one to detect the Th line even in stars with lower Th abundances."
    },
    "0402/astro-ph0402121_arXiv.txt": {
        "abstract": "In simulations of (magnetized-)fluid dynamics in physics and astrophysics, the visualization techniques are so frequently applied to analyse data that they have become a fundamental part of the research. Data produced is often a multi-dimensional set with several physical quantities, that are usually complex to manage and analyse. JETGET is a visualization and analysis tool we developed for accessing data stored in Hierarchical Data Format (HDF) and ASCII files. Although JETGET has been optimized to handle data output from jet simulations using the Zeus code from NCSA, it is also capable of analysing other data output from simulations using other codes. JETGET can select variables from the data files, render both two- and three-dimensional graphics and analyse and plot important physical quantities. Graphics can be saved in encapsulated Postscript, JPEG, VRML or saved into an MPEG for later visualization and/or presentations. An example of use of JETGET in analysing a 3-dimensional simulation of jets emanating from accretion disks surrounding a protostar is shown. The strength of JETGET in extracting the physics underlying such phenomena is demonstrated as well as its capabilities in visualizing the 3-dimensional features of the simulated magneto-hydrodynamic jets. The JETGET tool is written in Interactive Data Language (IDL) and uses a graphical user interface to manipulate the data. The tool was developed on a LINUX platform and can be run on any platform that supports IDL. JETGET can be downloaded (including more information about its utilities) from http://www.capca.ucalgary.ca/software. ",
        "introduction": "Today the scientific community is confronted with the problem of understanding and analysing large masses of numeric scientific data thanks to the availability of super-computing systems. One solution to this problem is scientific visualization: converting the numeric data into pictures more readily understandable by scientists. However, while visualization helps understand the overall behavior of the simulations, it is still a daunting task to extract the underlying physics. Ideally one would like to use a tool where both visualization and analysis are built in. Particularly, in many cases in computational (magnetized-)fluid dynamics in physics and astrophysics, extensive data processing is required to obtain meaningful information.  JETGET is  a data visualization and analysis software specifically designed to deal with data from such simulations. JETGET is IDL (Interactive Data Language) based. It provides a set of modules we built to aid the scientist when analysing and visualizing 2-dimensional (2-D and 2.5-D) and 3-dimensional (3-D) data. It can handle large datasets allowing both their graphical representation and analysis. JETGET allows users to interact with time histories, profiles, contour and surface plots on the same window, and quickly perform mathematical manipulations or combinations of data.   Here we describe the basic functionalities of JETGET. Note that  a help file for each module is provided which contains more detailed information. This paper is organized as follows. We start in section 2 by a description of the {\\it Setup} modules where specific information about the data files is specified. In sections \\ref{fieldlinessection}-\\ref{vectorssection} the different modules of JETGET are described. In particular, we show how the modules can be used to extract various and crucial information from the simulations. A discussion of the different output methods is given in section \\ref{outputsection} before concluding in section \\ref{conclusionsection}.   \\begin{figure} \\begin{center} \\resizebox{12cm}{!}{\\includegraphics{jan.eps}} \\caption{The {\\it Main} JETGET module. This is the main JETGET window. From here, all of JETGETs modules are accessible.} \\label{main} \\end{center} \\end{figure} ",
        "conclusions": "\\label{conclusionsection}  JETGET can be used to analyse and visualize most data output from simulations of (magneto-)hydrodynamic fluids, specifically simulations of jets. Besides the reconstruction of the three dimensional shapes and features of the jets, JETGET allows us to extract vital information concerning their energetics, dynamics and stability. This leads to a deeper understanding of jet physics. JETGET is free for download and can be found at\\\\ http://www.capca.ucalgary.ca/software .  \\begin{ack} Thanks to D. A. Clarke, W. Dobler, Ch. Fendt and J. Stil for helpful input. Few routines used in JETGET (namely, ``vcolorbar\\_\\_define.pro'', ``fsc\\_droplist.pro'' and ``xpanel.pro'') were kindly provided to us by David Fanning. The ``ve\\_\\_define.pro'' routine was kindly provided to us by RSI. \\end{ack}"
    },
    "0402/astro-ph0402317_arXiv.txt": {
        "abstract": "A puzzling feature of the {\\it Chandra}--detected quasar jets  is that their X-ray emission decreases faster along the jet than their radio emission, resulting to an outward increasing  radio to X-ray ratio. In some sources this behavior is so extreme that the radio emission peak is located clearly downstream of that of the X-rays. This is a rather unanticipated behavior given that the inverse Compton nature of the X-rays  and the synchrotron radio emission are  attributed to roughly the same electrons of the jet's non-thermal electron distribution. In this note we show that this morphological behavior can result from the gradual deceleration of a  relativistic flow and that the offsets in peak emission at different wavelengths carry the imprint of this deceleration. This notion is consistent with another recent finding, namely  that the jets feeding the terminal hot spots of powerful radio galaxies and quasars are still relativistic with Lorentz factors $\\Gamma \\sim 2-3$. The picture of the kinematics of powerful jets emerging from these considerations is that they remain relativistic as they  gradually decelerate from Kpc scales to the hot spots, where, in a final collision with the intergalactic medium, they slow-down rapidly to the subrelativistic velocities of the hot spot advance speed. ",
        "introduction": "The superior angular resolution and sensitivity of {\\it Chandra} has led to the discovery of X-ray emission from a number of quasar jets. Schwartz et al. (2000) were the first to note that the X-ray emission  from the knots of the jet of the superluminal quasar  PKS 0637-752 (Chartas et al. 2001), at a projected distance $\\sim 100$ Kpc from the quasar core, is part of a spectral component separate from its synchrotron radio-optical emission and it is too bright to be explained through Synchrotron - Self Compton (SSC) emission from electrons in energy equipartition with the jet magnetic field (note however that  in the innermost knot  of some sources -- e.g. 3C 273 (Marshall et al. 2001), PKS 1136-165  (Sambruna et al. 2002) -- a synchrotron X-ray contribution is possible). These properties are apparently common to other quasars jets, as indicated by the mounting observational evidence (Sambruna et al 2001; Marshall et al. 2001; Jester et al. 2002; Sambruna et al. 2002; Siemiginowska et al. 2002; Jorstad, Marscher,  \\& McHardy 2003; Siemiginowska et al. 2003; Yuan et al. 2003; Cheung 2004,  Sambruna et al. 2004).   To account for the level of the observed  X-ray emission, Tavecchio et al. (2000) and Celotti et al. (2001) proposed that this is due to  External Compton (EC) scattering of the Cosmic Microwave Background (CMB) photons off relativistic electrons in the jet, provided that the jet flow is sufficiently relativistic ($\\Gamma \\sim 10$) to  boost  the CMB energy density in the flow frame (by $\\Gamma^2$)  at the level needed  to reproduce  the observed X-ray flux.    In all these sources the radio-to-X-ray ratio increases downstream along the jet, an unexpected behavior, given that the cooling length of the EC X-ray emitting electrons ($\\gamma \\sim$ a few hundreds) is longer than that of the radio emitting ones ($\\gamma \\sim$ a few thousands) and comparable or  longer than the size of the jet, which would lead to a constant X-ray brightness as far out as the hot spots, contrary to observations. More surprisingly, in some jets (e.g. 3C 273 in Sambruna et al. 2001 and  Marshall et al. 2001; PKS 1136-135 and 1354+195  in Sambruna et al. 2002; PKS 1127-145 in Siemiginowska et al. 2002; 0827+243 in Jorstad, Marscher,  \\& McHardy 2003) the X-ray and radio maps are anti-correlated, with the X-ray emission peaking closer to the core, gradually decreasing outward, and the radio emission increasing outward to peak practically at the very end of the radio jet.  To explain the reduction of the X-ray flux along the jet Tavecchio, Ghisellini \\& Celotti (2003) suggested that the X-ray emission originates in a collection of micro-knots undergoing  adiabatic expansion sufficient to produce the desirable electron cooling. However, this would suppress also the radio emission, leading to practically indistinguishable radio and X-ray morphologies, contrary to their observed spatial anti-correlation. An elegant suggestion by Dermer \\& Atoyan (2002) that the X-rays are  synchrotron emission from electrons cooling in the Klein--Nishina regime naturally produces shorter sizes  in X-rays than in radio. However, it  also  produces larger optical than X-ray sizes, contrary to observations, and it seems therefore not to be applicable in this particular context.  Our view is that, although X-ray producing electrons are present throughout the jet, the X-ray brightness decreases because the CMB photon energy density in the flow  frame decreases along the jet as a result of a decelerating relativistic jet flow. The decrease in the flow Lorentz factor leads to a decrease in the comoving CMB photon energy density and hence to a decrease in the X-ray brightness along the jet. At the same time, the flow deceleration leads to a compression of the magnetic field; this results to an enhanced radio emission with distance, which gets thus displaced downstream of the EC X-ray emission.    Based on the radio, optical, and X-ray jet maps, we argue in this note, that powerful extragalactic jets are relativistic and gradually decelerating. In \\S 2 we formulate the synchrotron and EC emission process from a decelerating jet flow and present our results. In \\S 3 we discuss our findings and  touch upon some  open issues. ",
        "conclusions": "We have proposed that the increase of the radio-to-X-ray flux ratio along the length of the jets of powerful quasars, as well as the occasional offset of the jet images in these wavelengths, are naturally accounted for in terms of relativistic flows that decelerate over the entire length of the jet. Despite this deceleration, the jets remain relativistic ($\\Gamma \\sim$ a few) to their terminal hot spots (Georganopoulos \\& Kazanas 2003; hereafter GK03), within which they eventually attain sub-relativistic speeds. Our proposal provides, for the first time, a means for deducing the jet kinematics through simple models of their multiwavelength images; these can then be checked for consistency when coupled to the detailed hot spot emission, which as shown in earlier work (GK03), depends on the value of $\\Gamma$ at this location. We note here that a change of the Doppler factor can be also produced if the jet curves monotonically away from the line of sight without actually decelerating. However, in this case the CMB comoving photon energy density would remain constant along the jet, and the X-ray and radio emission would decrease along the jet in a similar manner, contrary to observations.   The scenario we propose here finds further support from modeling jets with  several sequential individual knots: Sambruna et al. (2001) noted the need for a gradual decrease of the Doppler factor and/or an increase of the magnetic field in order to reproduce the emission from the knots along the jet of 3C 273 with simple one zone models. These knot to knot variations can be naturally incorporated within the context of  a decelerating collimated flow, as we propose.  These same maps indicate also the need or not of distributed particle reacceleration along these jets, when the EC loss length scale of their radio emitting electrons (proportional to $(1+Z)^{-4} \\Gamma^{-2}$) is shorter than the observed length of the radio jet. This appears to be the case with sources like PKS 1127-145 at $Z=1.187$ (Siemiginowska et al. 2002) and possibly GB 1508+5714 at $Z=4.3$ (Yuan et al. 2003; Siemiginowska et al. 2003; Cheung 2004), which show their peak radio emission displaced from that of the X-rays: in the  absence of reacceleration, EC losses of the radio emitting electrons would produce radio jets shorter than the X-ray ones.  Much about the reacceleration process at these jets can be inferred from their morphology at different frequencies. For example, the  knotty optical jet morphology shows that reacceleration is not strong enough to offset the EC-dominated losses of the optically emitting synchrotron  electrons ($\\gamma \\sim 10^{6-7}$).  The radiative efficiency of these jets is less than a few percent (e.g. Tavecchio et al. 2000) and the energy lost in deceleration must be either used to heat up the matter in the jet, or must be transferred to material that is entrained by the jet. While the first option would result to an expansion of the jet, contrary to what is seen, entrainment from an external medium would load the jet  with baryonic mass while decelerating it. Entrainment would produce velocity gradients across the jet, and in this sense a faster spine and a slower sheath are to be expected. However, the observed X-ray and radio jet morphologies suggest that  the dominant effect must be a deceleration along the jet. A fast spine that  does not decelerate substantially and carries most of the jet power would produce  a constant X-ray EC and radio synchrotron flux along the jet, in disagreement with observations. A consequence of entrainment would be that even if the jet did not start as a baryonic one,  entrainment would gradually enrich it with baryons and eventually a fraction $\\Delta\\Gamma/\\Gamma$ of its power, where  $\\Delta \\Gamma$ is the decrease of the Lorentz factor, would be carried by the entrained baryonic matter. This in turn would increase the radio lobe equipartition energy content estimates derived under the assumption of a leptonic composition. Our findings point to a picture where powerful relativistic jets decelerate, depositing  most of their power in their surroundings in the form of kinetic energy, their observed radiation being  only the tip of the  iceberg."
    },
    "0402/astro-ph0402637_arXiv.txt": {
        "abstract": "{ An analytical MHD model of a normal-polarity prominence with compressible flow is presented. The exact solution is constructed via a systematic nonlinear separation of variables method used to calculate several classes of MHD equilibria in Cartesian geometry and uniform gravity. Although the model is 2D, a third magnetic/velocity vector field component is included and the highly sheared fields observed in prominences are reproduced. A description is given of the balance of gas pressure gradient with gravity and the Lorentz or inertial forces acting along and across the prominence. It is found that the flow does not significantly influence the heating profile.  The analyzed model has dimensions, plasma density, temperature and velocity profiles which agree with those in the observations literature. ",
        "introduction": "\\label{introduction}  The term {\\sl prominence} is used to describe various objects ranging from relatively stable ones with lifetimes of many months to transient phenomena lasting hours or less (Tandberg-Hanssen,~1995).  When they are seen in absorption against the disk they are referred to as {\\sl filaments}.  The long-lasting structures observed to last from days to months away from active regions are often called {\\sl quiescent prominences}.  They are long, cool, dense, sheet-like structures near-vertical to the solar surface supported by a series of arches whose feet are anchored in the photosphere.  In and around active regions, a different kind of shorter-lived prominences exist, refered to as {\\sl active region prominences}.  On the disk their appearance is like that of quiescent prominences except that they are generally smaller.  The category of active region prominences can be further subdivided into plage filaments, which are relatively stable prominences found above magnetic polarity inversion lines in or bordering active regions, and more dynamic phenomena such as surges, sprays and flare loops (Tandberg-Hanssen,~1995).  Quiescent prominences are structures of cool plasma suspended in the chromosphere or corona, usually above photospheric polarity inversion lines.  A prominence is said to be of normal or inverse polarity depending on whether the prominence field points in the same direction across the prominence as the field of the bipolar region below, or in the opposite direction. Before the crucial role played by magnetic fields in prominence physics was understood, prominences were regarded as cool objects in hydrostatic equilibrium with the hot corona.  However, Menzel (Bhatnagar et al.,~1951) argued that coronal magnetic fields could support prominences in static equilibrium. Different formulations of this problem were given by Dungey~(1953), Kippenhahn \\& Schl\\\"uter~(1957) and Brown~(1958).  In these models the dense prominence material is supported against gravity mainly by the Lorentz force.  Since then many models of prominence support have been developed, most of them two-dimensional because prominences are observed to be long, straight and reasonably uniform along their long axes structures.  Such normal and inverse polarity models include those by Anzer~(1972), Kuperus \\& Raadu~(1974), Lerche \\& Low~(1977), Malherbe \\& Priest~(1983), Anzer \\& Priest~(1985), Hood \\& Anzer~(1990), Fiedler \\& Hood~(1992), Low \\& Hundhausen~(1995), Low \\& Zhang~(2002), Fong et al.~(2002) and Low et al.~(2003).  Because of the mathematical complexity of the full 3D magnetohydrostatic equations most prominence modelling is 2D.  However, important exact 3D magnetostatic prominence models have been calculated by Low~(1982, 1984, 1992). Since the full 3D MHD equations are not amenable to analytical treatment (but see Petrie \\& Neukirch, 1999) we will focus on the basic macroscopic structure of a prominence. Full MHD normal prominence models have rarely been attempted in the past (in 1D by Tsinganos \\& Surlantzis,~1992; in 2D by Ribes \\& Unno, ~1980; Del Zanna \\& Hood,~1996).  As well as support against gravity, also important is the energy balance within a prominence.  The energy balance in prominences has been modelled using radiative transfer theory by many authors, e.g. Poland et al.,~(1971), Heasley \\& Mihalas~(1976), Heinzel et al.,~(1987), Paletou et al.,~(1993), Gontikakis et al.,~(1997) and Anzer \\& Heinzel~(1999, 2000) while Poland \\& Mariska~(1986) have modelled 1D normal prominences with a unidirectional flow and asymmetric heating. In this project we include a full MHD momentum balance and a simple treatment of the energy balance in a prominence model for the first time. The focal object of the present study is to investigate the effect of non-isothermal compressible flows in prominence dips for the first time and, following on from Paper 2, to check if these flows influence the heating as such flows do in coronal loops.  The paper is organised as follows.  The analytical modelling technique is outlined in Sect. \\ref{analyticalmodel}.  A model fitted to typical observed physical parameter values is presented in Sect. \\ref{models} and the results are summarized in Sect. \\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  We have modelled a prominence by using a two-dimensional compressible equilibrium solution of the full ideal steady MHD equations with a consistent heating included in the model for the first time.  Our model generalises known self-similar prominence models, such as those by % Hood \\& Anzer~(1990), and Del Zanna \\& Hood~(1996). The heating model takes into account non-LTE radiation for the first time in a full MHD model by exploiting a radiative transfer model by Kuin \\& Poland~(1991).  Although the model is 2D, a third component of the magnetic and velocity vector fields allow us to model the highly sheared fields observed in prominences.  Unlike the coronal loop model in Paper 2, the heat in/out of the flow does not influence the energy equation significantly.  The model is consistent with an ionisation ratio of order unity according to the radiation model of Kuin \\& Poland~(1991).  This is consistent with several observations. Both magnetic diffusion and cross-field diffusion of neutrals are found to be insignificant within the time scales of interest so that an ideal MHD description is reasonable.  The modelled prominence dip must be very shallow for the physical parameters to stay within reasonable bounds. This is also consistent with observations.  Within the prominence the plasma is so dense that the gas pressure bears most of the burden of the prominence weight.  The supporting role of the magnetic field may be more important underneath the prominence where the plasma is less dense and the magnetic field may be compressed, and therefore stronger and more dipped, than inside the prominence.  We were unable to add self-consistently a surrounding hotter arcade solution separated from the cooler prominence either with an MHD discontinuity, because this would imply a huge enthalpy change there, or with a tangential discontinuity between thermally isolated prominence and coronal field lines. Such a global prominence model remains a challenge for the future."
    },
    "0402/astro-ph0402292_arXiv.txt": {
        "abstract": "We investigated the brightness evolution of 7 FU\\,Ori systems in the $1-200\\,\\mu$m wavelength range using observations from the {\\it Infrared Space Observatory} (ISO), 2MASS and MSX data. The SEDs were compared with earlier ones derived from the IRAS photometry and ground-based observations around the epoch 1983. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0402/astro-ph0402547_arXiv.txt": {
        "abstract": "Since the IAU XXIV meeting in 2000, the CMB anisotropy has matured from being one of a number of cosmological probes to forming the bedrock foundation for what is now the standard model of cosmology. The large advances over the past three years have come from making better and better maps of the cosmos. We review the state of measurements of the anisotropy and outline some of what we have learned since 2000. The recent advancements may be placed roughly into three categories: 1) What we learn from the CMB with minimal input from other cosmic measurements such as the Hubble constant; 2) What we learn from the CMB in combination with other probes of large scale structure; and 3) What we learn by using the CMB as a back light. Future directions are also discussed. It is clear: we have much more to learn from the CMB anisotropy. ",
        "introduction": "It has long been appreciated that the CMB anisotropy could be a powerful probe of cosmology. The foundations of the anisotropy calculations we do today were set out over thirty years ago by Sachs \\& Wolfe (1967), Rees (1968), Silk (1968), Peebles \\& Yu (1970), and Sunyaev \\& Zeldovich (1970). Plots of the acoustic peaks were shown in Doroskevich, Zeldovich, \\& Sunyaev (1978) and Bond \\& Efstathiou (1984) gave the results of detailed numerical calculations.  On the measurement side, the tension between expectations and continuously improving upper limits (e.g., Weiss 1980, Wilkinson 1985, Partridge 1995) was finally alleviated by the discovery of the anisotropy by {\\sl COBE} (Smoot et al. 1992).  At that time, the measured Sachs-Wolfe plateau ($l<20$) was a factor of two {\\it higher} than expectations based on the standard cold dark matter model in which $\\Omega_m\\approx 1.$\\footnote{We use the convention that $\\Omega_m=\\Omega_{cdm}+\\Omega_b+\\Omega_\\nu$ is the cosmic density in all matter components where $cdm$ is cold dark matter, $b$ is for baryons, and $\\nu$ is for neutrinos;  $\\Omega_r$ is the cosmic radiation density (now minuscule); $\\Omega_\\Lambda$ is the corresponding density for a cosmological constant; and $\\Omega_k$ is the corresponding curvature parameter. The Friedmann equation tells us: $1\\equiv\\Omega_\\Lambda+\\Omega_k+\\Omega_m=\\Omega_{tot}+\\Omega_k$. The physical densities are given by, for example, $\\omega_b=\\Omega_bh^2$.} There were many measurements of the anisotropy at the {\\sl COBE} scales and finer between 1992 and 2000 (e.g., see Page 1997 for a table) that culminated in observations of the first acoustic peak (Dodelson \\& Knox 2000, Hu 2000, Pierpaoli, Scott \\& White 2000, Knox \\& Page 2000).  Amidst theories that did not survive observational tests and false clues, a standard cosmological model emerged. Even over a decade ago, the evidence from a  majority of independent tests indicated $\\Omega_m\\approx0.3$ (e.g., Ostriker 1993). It was realized by many that a flat model ($\\Omega_k=0$) with a significant cosmic constituent with negative pressure, such as a cosmological constant, was a good fit to the data. Then, in 1998 measurements of type 1a supernovae (Riess et al. 1998, Perlmutter et al. 1999) directly gave strong indications that the universe was accelerating as would be expected from a cosmological constant. The state of the observations in 1999 is summarized in Figure~1. If the Einstein/Friedmann equations describe our universe, the data were telling us that the universe is spatially flat with matter density $\\Omega_m\\approx 0.3$ and $\\Omega_{\\Lambda}\\approx 0.7$. Independent analyses came to similar conclusions (e.g., Lineweaver 1998, Tegmark \\& Zaldarriaga 2000). The story since the last IAU meeting is that the concordance model does in fact describe virtually all cosmological observations astonishing well. There is now a well agreed upon standard cosmological model (Spergel et al. 2003, Freedman \\& Turner 2003).  \\begin{figure} \\plotfiddle{f1.ps}{3.2truein}{0}{70}{70}{-120}{-20} \\caption{The Cosmic Triangle from Bahcall et al. (1999). This shows the concordance model for cosmological observations at the end of the last millenium. Three different classes of observations, supernovae, clusters, and CMB anisotropy, are consistent if we live in a spatially flat universe with a cosmological constant.} \\label{fig:triangle} \\end{figure}  In the rest of this article we briefly review the CMB observations in \\S2 and summarize what we learn from (almost) just the CMB in \\S3. We then, in \\S4, outline the things we learn by combining the CMB with other maps of the cosmos, in particular the 2dF Galaxy Redshift Survey (Colles et al. 2001). In \\S5 we indicate the sorts of things we hope to learn by using the CMB as a back light for lower redshift phenomena. We conclude in  \\S6. ",
        "conclusions": "\\label{sec:cmbcon}  Over the past few years, and especially with {\\sl WMAP}, the CMB data has become the foundation for the standard cosmological model. Any model that purports to explain the birth and evolution of the universe must be able to predict the results in Figure~2. This is a very stringent requirement. The model elements implicit in the figure---superhorizon fluctuations with cosmic structure seeded by a scale invariant spectrum with Gaussian fluctuations in the metric are at the core of our conception of the universe. They are also at the heart of inflation. Indeed, we have started to directly constrain models of inflation (Peirie et al. 2003). This is not to say that our currently favored model is correct. There are elements of the observations, for example the apparent suppression of fluctuations on the largest angular scales, that may call for something beyond the standard model. However, it is truly astounding that we have a model that naturally explains almost all cosmological observations. The model is eminently testable and precise enough to be experimentally challenged. The model is also young enough to admit new discoveries in such areas as dark energy, dark matter, and the birth and growth of cosmic structure.  We have much more to learn from the CMB. To borrow from Winston Churchill, {\\sl WMAP} marks not the end, not even the beginning of the end, but rather the end of the beginning of what we can learn from the CMB. In IAUs ahead we may hope to hear of how observations of the CMB in combination with other cosmic probes determine the mass of the neutrino or the equation of state of the dark energy. Detection of polarization B-modes (See A. Couray, these proceedings) may be be able to tell us the energy scale of inflation. From the ground, new experiments such as ACT, APEX, and SPT are pushing CMB anisotropy measurements to high $l$ and high sensitivity. New experiments such as BICEP, CAPMAP, Polarbear, QUAD, {\\sl SPORT}, are applying new techniques to measure the polarization in the CMB. There is already talk of {\\sl CMBPOL}, a post-Planck satellite dedicated to polarization measurements. No doubt, precise measurements of the CMB will continue to shed light on fundamental physics, cosmology, and astrophysics for years to come."
    },
    "0402/hep-ph0402187_arXiv.txt": {
        "abstract": "Gravitinos are expected to be produced in any local supersymmetric model. Using their abundance prediction as a function of the reheating energy scale, it is argued that the next generation of {\\it Cosmic Microwave Background} experiments could exclude supergravity or strongly favor \"thermal-like\" inflation models if $B$ mode polarized radiation were detected. Galactic cosmic--ray production by evaporating primordial black holes is also investigated as a way of constraining the Hubble mass at the end of inflation. Subsequent limits on the gravitino mass and on the related grand unification parameters are derived.\\\\  \\begin{center} Phys. Rev. D 69 (2004) 105021 \\end{center} ",
        "introduction": "Although not yet experimentally discovered, supersymmetry (SUSY) is still the best - if not the only - natural extension of the standard model of particle physics. It could provide a general framework to understand the origin of the fundamental difference between fermions and bosons and could help to resolve the difficult problem of mass hierarchies, namely the instability of the electroweak scale with respect to radiative corrections. In global supersymmetry, the generator spinors $\\xi$ are assumed to obey $\\partial_{\\mu}\\xi=0$ \\cite{freedman}. If one wants to deal with local supersymmetry, or supergravity, this condition must be relaxed and $\\xi$ becomes a function of the space coordinates $x$. New terms, proportional to $\\partial_{\\mu}\\xi(x)$, must be canceled by introducing a spin 3/2 particle, called gravitino, as vector bosons are introduced in gauge theories. The gravitino is part of an N=1 multiplet which contains the spin 2 graviton (see \\cite{olive} for an introductive review) and, in the broken phase of supergravity, super-Higgs effects make it massive through the absorption of the Nambu-Goldstone fermion associated with the SUSY breaking sector.\\\\  It has long been known that if the gravitino is unstable some severe constraints on its mass must be considered in order to avoid entropy overproduction \\cite{weinberg}: $m_{3/2}\\gtrsim 10$ TeV. On the other hand, if the gravitino is stable, its mass should satisfy $m_{3/2}\\lesssim 1$~keV \\cite{pagels} to keep the gravitinos density smaller than the full Universe density ($\\Omega_{3/2}<\\Omega_{tot}$).  In spite of the huge dilution, those constraints are not fully evaded by inflation as gravitinos should be reproduced by scattering processes off the thermal radiation after the Universe has reheated \\cite{nano,khlo,ellis,Leigh,Fujisaki,ellis2,Kallosh}. As the number density of such secondary gravitinos is expected to be proportional to the reheating temperature, it is possible to relate the energy scale of inflation with the requirement that they are not overproduced.\\\\  In the first part of this paper, the next generation of cosmic microwave background (CMB) detection experiments is considered as a way of possibly excluding supergravity. It is shown that the energy scale of inflation required to produce an observable amount of tensor mode in the background radiation is not compatible with local supersymmetry in the standard cosmological scenario. In the second part, a new way of constraining the gravitino mass, based on evaporating primordial black holes, is investigated. Taking into account that the black hole masses cannot be much smaller than the Hubble mass at the formation epoch, it is suggested that a detection of cosmic--rays produced by the Hawking mechanism would lead to a lower bound on the reheating scale and, therefore, on the gravitino mass. Links with grand-unified models are given, as an example, in the conclusion. Finally, the basics of the propagation model used to relate the source term to the local spectrum are given in the Appendix A. ",
        "conclusions": "It must be pointed out that such possible constraints on the gravitino mass can be translated into constraints on more fundamental parameters, making them very valuable in the search for the allowed parameter space in {\\it grand unified} models. As an example, in models leading naturally to mass scales in the $10^2$-$10^3$~GeV range through a specific dilaton vacuum configuration in supergravity, the gravitino mass can be related with the GUT parameters \\cite{tkach}: $$ m_{3/2}=\\left( \\frac{5\\pi^{\\frac{1}{2}}\\lambda}{2^{\\frac{3}{2}}}\\right)^{\\sqrt{3}} (\\alpha_{GUT})\\left(\\frac{M_{GUT}}{M_{Pl}}\\right)^{3\\sqrt{3}}M_{Pl}. $$ With $M_{GUT}\\sim 10^{16}$~GeV and a gauge coupling $\\alpha_{GUT}\\sim 1/26$. The superpotential value in the  dilaton direction defines the magnitude of the coupling constant $\\lambda$ of the self-interacting 24 multiplet. Figure~\\ref{lambda} shows how the lower value on $\\lambda$ evolves as a function of the reheating temperature which could be probed by the previously given method, for three different branching ratios. Although not very constraining, this lower limit of the order $1.4\\times10^{-3}$ over the full tested range for $B=1$ could be one of the first experimental constraints on $\\lambda$.  The next generation of CMB experiments will face a new situation. Important efforts are devoted to the search for the polarization $B$ mode \\cite{minneapolis} and the sensitivity should reach scales of inflation of order $10^{15} - 10^{16}$~GeV. This value is slightly higher than the GUT scale if supersymmetry is ignored ({\\it i.e} if gravitinos production is expected not to have occurred), and slightly lower than the GUT scale if supersymmetry is taken into account ({\\it i.e.} in the case gravitinos are expected to be produced by scattering processes). Considering that the grand unified scale is the highest natural value for the reheating temperature, this means that, if a significant amount of entropy was not released after the moduli production, it should not be possible to detect those tensor modes in both scenarios.   On the other hand, cosmic--ray experiments could be sensitive enough to investigate the allowed reheating temperatures if small black holes were formed at the end of inflation. In this case, important limits could be derived on the gravitino mass and on the related GUT parameters.\\\\"
    },
    "0402/astro-ph0402221_arXiv.txt": {
        "abstract": "XMM-Newton observed  SAX J2103.5+4545 on January 6, 2003, while RXTE was monitoring the source. Using RXTE-PCA dataset between December 3, 2002 and January 29, 2003, the spin period and average spin-up rate during the XMM-Newton observations were found to be $354.7940\\pm0.0008$ s and $(7.4\\pm0.9)\\times10^{-13}$Hz s$^{-1}$ respectively. In the power spectrum of the 0.9-11 keV EPIC-PN lightcurve, we found quasi periodic oscillations around 0.044 Hz (22.7 s) with an rms fractional amplitude $\\sim $6.6 \\%. We interpreted this QPO feature as the Keplerian motion of inhomogenuities through the inner disk. In the X-ray spectrum, in addition to the power law component with high energy cutoff and $\\sim6.4$ keV fluorescent iron emission line (Baykal et al., 2002), we discovered a soft component consistent with a blackbody emission with ${\\rm{kT}}\\sim1.9$keV. The pulse phase spectroscopy of the source revealed that the blackbody flux peaked at the peak of the pulse with an emission radius $\\sim 0.3$ km, suggesting the polar cap on the neutron star surface as the source of blackbody emission. The flux of the iron emission line at $\\sim 6.42$ keV was shown to peak at the off-pulse phase, supporting the idea that this feature arises from fluorescent emission of the circumstellar material around the neutron star rather than the hot region in the vicinity of the neutron star polar cap. ",
        "introduction": "The transient X-ray source SAX J2103.5+4545 was discovered by the Wide Field Camera on the {\\it{BeppoSAX}} X-ray observatory during its outburst between 1997 February and September with 358.61s pulsations and a spectrum consistent with an absorbed power law model with the photon index of $\\sim1.27$ and the absorption column density of $\\sim3.1\\times 10^{22}$cm$^{-2}$ (Hulleman, in't Zand, \\& Heise 1998).  After detection of another outburst in November 1999 by the {\\it{all-sky monitor (ASM)}} on the {\\it{Rossi X-ray Timing Explorer (RXTE)}},  the source was found to be active for more than a year, and was continously monitored through regular pointed {\\it{RXTE}} observations. Using pulse arrival times, the orbital period and eccentricity of the orbit were found to be 12.68(25) days and 0.4(2)  (Baykal, Stark, \\& Swank 2000a,b). In the timing analysis, the source was initially found to be spinning up for $\\sim 150$ days, at which point the flux dropped quickly by a factor of $\\simeq 7$, and a weak spin-down began afterwards (Baykal, Stark, \\& Swank 2002). Strong correlation between X-ray flux and spin-up rate was explained by using Ghosh \\& Lamb (1979) accretion disk model. The X-ray spectra well fitted the absorbed power law model with high energy cutoff and a $\\sim $6.4 keV fluorescent emission line (Baykal et al. 2002).  Orbital parameters found by using {\\it{RXTE}} observations of the source (Baykal et al. 2000a,2000b) indicated that the source has a high mass companion. Hullemann et al. (1998) pointed out a B8 type star within the BeppoSAX error box, but its distance ($\\sim 0.7$ kpc) implied a luminosity too low to explain the spin-up that was seen in the RXTE observations. Recently, a possible candidate for the optical companion of SAX J2103.5+4545 with the visual magnitude of 14.2 was discovered (Reig \\& Mavromatakis, 2003).  SAX J2103.5+4545 was also observed with the {\\it{INTEGRAL}} observatory in the 3-200 keV band with significant detection up to $\\sim 100$ keV (Lutovinov, Molkov,\\& Revnivtsev 2003). The spectral parameters found in the {\\it{INTEGRAL}} observations of the source were found to be compatible with those found by Baykal et al. (2002).  Since the beginning of the most recent outburst in June 2002, SAX J2103.5+4545 has been monitored continously by {\\it{RXTE}} through regular pointed observations. It was possible to obtain some simultaneous coverage with the {\\it{XMM-Newton}} observatory on January 6, 2003. The observation of {\\it{XMM-Newton}} revealed a soft spectral component of the source which was well-represented by a blackbody model. This spectral model was verified by simultaneous fitting of January 6, 2003 {\\it{RXTE}}-PCA observation. Using {\\it{XMM-Newton}} dataset, we also discovered $\\sim 22.7$s quasi periodic oscillations (QPO's) of this source for the first time. In this paper, we present our spectral and timing results of the analysis of RXTE and XMM datasets of SAX J2103.5+4545. ",
        "conclusions": "\\subsection{QPO Feature of SAX J2103.5+4545} Quasi-periodic oscillations in the X-ray band having periods in the range of $\\sim 2.5-100$s have been observed in many accretion powered X-ray pulsars: 4U 0115+63 (Soong \\& Swank 1989), EXO 2030+375 (Angelini, Stella,\\& Parmar 1989), 4U 1626-67 (Shinoda et al. 1990), SMC X-1 (Angelini, Stella,\\& White 1991), Cen X-3 (Takeshima et al. 1991), V0332+53 (Takeshima et al. 1994), A0535+262 (Finger, Wilson,\\& Harmon 1996), GRO J1744-28 (Zhang, Morgan,\\& Jahoda 1996; Kommers et al. 1997), X Per (Takeshima 1997), 4U 1907+09 (in't Zand, Baykal, \\& Strohmayer 1998; Mukerjee et al. 2001), XTE J1858+034 (Paul \\& Rao 1998), LMC X-4, and Her X-1 (Moon \\& Eikenberry 2001a,b). The QPO feature that we found in the XMM-Newton EPIC-PN light curve of SAX J2103.5+4545 which has a peak period of $22.7\\mp 0.6$s and fractional rms amplitude of $6.6 \\pm 1.9$ percent is quite typical (e.g. In't Zand et. al. 1998; Paul \\& Rao 1998; Takeshima et al. 1994).  Models that explain the QPO phenomenon in accretion powered X-ray pulsars fall basically into three categories: In the Keplerian frequency model, QPOs are produced due to some inhomogenuities at the inner edge of the Keplerian disk ($r_0$) and modulate the light curve at the Keplerian frequency $\\nu _{QPO}=\\nu_{K}$ (van der Klis et al. 1987). In the beat frequency model, the accretion flow onto the neutron star is modulated at the beat frequency between the Keplerian frequency at the inner edge of the accretion disk and the neutron star spin frequency $\\nu _{QPO}=\\nu_{K}-\\nu _{s}$ (Alpar \\& Shaham 1985). The third model involves accretion flow instabilities (Fronter, Lamb,\\& Miller 1989; Lamb 1988), and applies only to the sources that have luminosities close to Eddington limit, therefore it should not be applicable to our case for which the luminosity is well below the Eddington limit.  In our case, QPO frequency $\\nu _{QPO}= 4.4\\times 10^{-2}$Hz is about one order of magnitude greater than the spin frequency $\\nu _{s}= 2.8185\\times 10^{-3}$ Hz. Therefore, it is difficult to distinguish between a Keplerian model and a beat frequency model.  Assuming that the 22.7 s oscillation in SAX J2103.5+4545 is related to Keplerian orbital motion via either Keplerian frequency model or beat frequency model, and using the QPO and its FWHM values we obtain the radius of inner disk as \\begin{equation} r_{0}= \\bigg(\\frac{GM}{4\\pi^2}\\bigg)^{1/3}\\nu_k^{-2/3} = (1.32^{+0.13}_{-0.11}) \\times 10^{9} {\\rm{cm}}, \\end{equation} where M is 1.4 M$_{\\odot}$ for a neutron star and G is the gravitational constant.  From the strong correlation between pulse frequency derivatives and X-ray flux, Baykal et al. (2002) obtained for the distance to the source $3.2 \\pm 0.8$ kpc and for the magnetic field $(12\\pm 3)\\times 10^{12}$ Gauss. Using the distance and magnetic field values, the inner edge of the Keplerian disk $r_{0}$ can be found as (Ghosh \\& Lamb 1979) \\begin{equation} r_{0}\\simeq 0.52\\mu^{4/7}(2GM)^{-1/7}\\dot M^{-2/7} = (1.67_{-0.25}^{+0.23})\\times 10^9 {\\rm{cm}}, \\end{equation} where $\\mu=BR^{3}$ is the neutron star magnetic moment with B the equatorial magnetic field, R the neutron star radius, and $\\dot M$ the mass accretion rate having the value of $\\simeq 4\\times 10^{15}$g $s^{-1}$ for an accretion luminosity of $\\simeq 7.5\\times 10^{35}$ erg s$^{-1}$ (as estimated in Section 3.3).  The radius of the inner disk inferred from the Keplerian orbital motion of inhomogenuities and the one inferred from the Ghosh Lamb disk accretion model agree each other. This shows that the idea that the QPOs are formed due to the Keplerian motion of inhomogenuities is indeed promising as the explanation of the QPO of SAX J2103.5+4545 and the observed QPO frequency is consistent with the distance and the magnetic field values estimated by Baykal et al. (2002).  \\subsection{Blackbody and Iron Line Features of the Energy Spectrum}  XMM-Newton observations of SAX J2103.5+4545 revealed for the first time that the energy spectrum of the source has a blackbody component peaking at $\\sim 1.90$keV with the emission radius of $\\sim 0.3$km. The blackbody radiation may come from the polar cap of the neutron star as it appears to for the Be/X-ray pulsar system EXO 2030+375 (Reig \\& Coe 1999; Sun et al. 1994) and the millisecond X-ray pulsars SAX J1808.4-3658 (Gierlinski, Done,\\& Didier 2002) and XTE J0920-314 (Juett et al. 2003). Blackbody emission radii on these X-ray pulsars are reported to be greater than $\\sim 1$km. The relatively high surface magnetic field of SAX J2103.5+4545 ($\\sim 10^{13}$ Gauss) is probably the reason for the relatively small blackbody emission radius ($\\sim 0.3$ km) compared to these X-ray pulsars. Although the contribution of blackbody component is relatively more significant for lower energies (i.e. energies smaller than $\\sim 3$ keV), power law flux is $\\sim 3$ times greater than the blackbody flux even at the 1-3 keV energy band.  In our case, it is unlikely that the blackbody emission comes from the reprocessed emission of the surrounding material or the accretion disk as in the case of Her X-1 (Endo et al. 2000), Cen X-3 (Burderi et al. 2000), SMC X-1 and LMC X-4 (Paul et al. 2002), since blackbody component in such cases is expected to be softer ($kT\\sim 0.1$keV). Lower blackbody temperature and smaller blackbody emission radius at the off-pulse phase shown in Figure 4 are also indications of the plausibility of the polar cap emission interpretation, as the regions of the soft polar cap emission must align with the peak of the X-ray pulse of the pulsar.  Using the spin-phase resolved spectroscopy, strength of the iron line feature at $\\sim 6.42$keV was also found to vary significantly with the spin phase as seen in Figure 4. The peak energy of this feature clearly shows that it corresponds to the fluorescent iron K-line complex. This line complex feature is observed in the spectra of most of the X-ray pulsars (White et al. 1983; Nagase 1989) and is generally thought to be produced by the ions less ionized than Fe XVIII in a relatively cool matter around the neutron star (e.g. accretion disk, accretion disk corona) by fluorescent K$\\alpha$ transition.  Variation of the iron line feature with the spin phase can then be interpreted as a sign that it is mainly produced outside the polar cap region of the neutron star, thus should have a peak at the off-pulse parts of the spin phase. From Figure 4, we see that iron line flux and iron line equivalent width vary strongly with the spin phase, peaking at the off-pulse. Similar pulse phase dependence of the iron line feature is also observed in Her X-1 (Choi et al. 1994)."
    },
    "0402/gr-qc0402099_arXiv.txt": {
        "abstract": "We construct solutions of plane symmetric wormholes in the presence of a negative cosmological constant by matching an interior spacetime to the exterior anti-de Sitter vacuum solution. The spatial topology of this plane symmetric wormhole can be planar, cylindrical and toroidal. As usual the null energy condition is necessarily violated at the throat.  At the junction surface, the surface stresses are determined.  By expressing the tangential surface pressure as a function of several parameters, namely, that of the matching radius, the radial derivative of the redshift function and of the surface energy density, the sign of the tangential surface pressure is analyzed. We then study four specific equations of state at the junction: zero surface energy density, constant redshift function, domain wall equation of state, and traceless surface stress-energy tensor. The equation governing the behavior of the radial pressure, in terms of the surface stresses and the extrinsic curvatures, is also displayed. Finally, we construct a model of a plane symmetric traversable wormhole which minimizes the usage of the exotic matter at the throat, i.e., the null energy condition is made arbitrarily small at the wormhole throat, while the surface stresses on the junction surface satisfy the weak energy condition, and consequently the null energy condition.  The construction of these wormholes does not alter the topology of the background spacetime (i.e., spacetime is not multiply-connected), so that these solutions can instead be considered domain walls. Thus, in general, these wormhole solutions do not allow time travel. ",
        "introduction": "An important aspect in black hole physics is that they can be formed through gravitational collapse of matter. This is indeed the case for spherical collapse in an asymptotically flat background.  For other backgrounds, such as asymptotically anti-de Sitter spacetime with plane symmetry, it was found that gravitational collapse of plane symmetric distributions of matter also results in an event horizon \\cite{lemos1}.  The event horizon may have planar \\cite{lemos2}, cylindrical \\cite{lemos3} or toroidal topology \\cite{lemos3,zanchin}. Indeed, the collapse of planar distributions of matter, in a background with a negative cosmological constant, can form a planar black hole (or a black membrane) violating somehow the hoop conjecture.  Upon compactification of one or two coordinates one finds that cylindrical black holes (or black strings), or toroidal black holes can also form from the gravitational collapse of a cylindrical or toroidal distribution of matter, respectively.  In these solutions the mass parameter is a surface mass energy in the planar case, a linear mass density in the cylindrical case, and a mass in the toroidal case \\cite{lemos1}.  A natural extension of these solutions would be to add exotic matter to obtain plane symmetric traversable wormhole solutions, with planar, cylindrical and toroidal topologies. These would add to other non-spherically symmetric wormholes which have already been considered by several authors. For instance, extending the spherically symmetric Morris$\\,$-Thorne wormholes \\cite{Morris}, Visser \\cite{Visser89} motivated by the aim of minimizing the violation of the energy conditions and the possibility of a traveller not encountering regions of exotic matter in a traversal through a wormhole, constructed polyhedral wormholes and, in particular, cubic wormholes. These contained exotic matter concentrated only at the edges and the corners of the geometrical structure, and a traveller could pass through the flat faces without encountering matter, exotic or otherwise.  Gonz\\'alez-D\\'{\\i}az generalized the static spherically symmetric traversable wormhole solution to that of a (non-planar) torus-like topology \\cite{GDiaz}. This geometrical construction was denoted as a ringhole. Gonz\\'alez-D\\'{\\i}az went on to analyze the causal structure of the solution, i.e., the presence of closed timelike curves, and has recently studied the ringhole evolution due to the accelerating expansion of the universe, in the presence of dark energy \\cite{GDiaz2}. Other interesting non-spherically symmetric traversable spacetimes are the stationary solution obtained by Teo \\cite{Teo}, and an axially symmetric traversable wormhole solution obtained by Kuhfittig \\cite{Kuhf}. In this work we study plane symmetric wormholes. We match an interior static and plane wormhole spacetime to a vacuum solution with a negative cosmological constant, i.e., to an anti-de Sitter spacetime. The properties of the junction surface, such as the surface stresses are determined. For the wormhole solutions quoted above, although the throat geometries differ from solution to solution, all these spacetimes are asymptotically flat with trivial topology, at infinity, whereas for the solutions we analyze the infinity carries the same topology as the throat, meaning that these solutions can also be considered as domain walls. Wormholes with this same property are the spherical wormholes joining two Friedmann-Robertson-Walker universes \\cite{visserhoch}.   The plan of this paper is as follows: In Sec. II we present a plane symmetric metric with a negative cosmological constant and analyze the mathematics of embedding in order to obtain a wormhole solution. In Sec. III, we present the Einstein equations for the interior solution, and verify that the null energy condition is necessarily violated at the wormhole throat. In Sec. IV we deduce the exterior plane vacuum solution, through the Einstein equations. In Sec. V, we match an interior plane spacetime to the vacuum solution with a negative cosmological constant, and deduce the surface stresses at the thin-shell. By expressing the tangential surface pressure as a function of several parameters, namely, that of the matching radius, the radial derivative of the redshift function and of the surface energy density, the sign of the tangential surface pressure  is analyzed. We also obtain a general equation governing the behavior of the radial pressure/tension across the junction in terms of the surface stresses. In section VI, we construct a plane traversable model that minimizes the null energy condition violation at the throat, and in which the surface stresses on the thin shell satisfy the energy conditions. Finally, we conclude in Sec. VII. ",
        "conclusions": "We have constructed plane symmetric wormholes (with planar, cylindrical and toroidal topologies) in an anti-de Sitter background. We have determined the surface stresses, analyzed the sign of the tangential surface pressure, displayed an equation relating the radial pressure across the junction boundary, and given a model which minimizes the usage of exotic matter.  We have found that the construction of these wormholes does not involve a topology change, and thus the wormhole solution can be considered a domain wall. As such these wormholes do not allow time travel. We have not considered in this work solutions with zero or positive cosmological constants as, for plane symmetry, they yield solutions with negative total masses.        \\appendix"
    },
    "0402/astro-ph0402198_arXiv.txt": {
        "abstract": "{ The recent detection of gamma-ray lines from radioactive \\al \\ and \\fe \\ in the Milky Way by the RHESSI satellite calls for  a reassessment of the production sites of those nuclides. The observed gamma-ray line flux ratio is in agreement with calculations  of nucleosynthesis in massive stars, exploding as SNII (Woosley and Weaver 1995); in the light of those results, this observation would suggest then that SNII are the major sources of \\al \\ in the Milky Way, since no other conceivable source produces substantial amounts of \\fe. However, more recent theoretical studies find that SNII produce much higher \\feal \\ ratios than previously  thought and, therefore, they cannot be the major \\al \\ sources in the Galaxy (otherwise \\fe \\ would be detected long ago, with a line flux similar to the one of \\al). Wolf-Rayet stars, ejecting \\al \\ (but not \\fe) \\ in their stellar winds, appear  then as a most natural candidate. We point out, however, that this scenario faces also an important difficulty. Forthcoming results of ESA's INTEGRAL satellite, as well as consistent calculations of nucleosynthesis in massive stars (including stars of initial masses as high as 100 \\Ms \\ and metallicities up to 3 \\Zs), are required to settle the issue. ",
        "introduction": "\\al \\ is the first radioactive nucleus ever detected in the Galaxy through its characteristic gamma-ray line signature, at 1.8 MeV (Mahoney et al. 1982). Taking into account its short lifetime ($\\sim$1 Myr), its detection convincingly demonstrates that nucleosynthesis is still active in the Milky Way (Clayton 1984). The detected flux ($\\sim$4 10$^{-4}$ cm$^{-2}$ s$^{-1}$) corresponds to $\\sim$2 \\Ms \\ of \\al \\ currently present in the ISM (and produced per Myr, assuming a steady state situation). The COMPTEL instrument aboard CGRO mapped the 1.8 MeV emission in the Milky Way and found it to be irregular, with prominent \"hot-spots\" probably associated with  the spiral arms (Diehl et al. 1995). The spatial distribution of \\al \\ suggests that massive stars are at its origin (Prantzos 1991, 1993, Prantzos and Diehl 1996). However, it is not yet clear whether the majority of observed \\al \\ originates from the winds of the most massive stars (i.e. above 30 \\Ms, evolving as Wolf-Rayet stars) or from the explosions of less massive stars (i.e. in the 12-30 \\Ms \\ range, exploding as SNII); the uncertainties in the corresponding stellar yields are still quite large (see Sec. 2) and do not allow to conclude yet.  Clayton (1982)  pointed out that SNII explosions produce another relatively short lived radioactivity, \\fe \\ (lifetime $\\sim$2 Myr). Since WR winds do not eject that isotope, the detection of its characteristic gamma-ray lines \\footnote{At 1.117 and 1.332 MeV, resulting from the decay of its daughter nucleus $^{60}$Co}   in the Milky Way would constitute a strong argument  for SNII being at the origin of \\al. Based on detailed nucleosynthesis calculations of SNII (from Woosley and Weaver 1995) Timmes et al. (1995) found that the expected gamma-ray line flux ratio of \\feal \\ (for each of the  two lines of \\fe) is 0.16, if SNII are the only sources of \\al \\ in the Milky Way.  The Reuven Ramaty High Energy Solar Spectroscopic Imager (RHESSI) detected  the galactic \\al \\ emission at a flux level compatible with previous observations (Smith 2003a). Most recently, Smith (2003b) reported the first ever detection of the Galactic \\fe \\ gamma-ray lines with RHESSI; their combined fluxes correspond to a significance level slightly higher than 3 $\\sigma$. The line flux ratio \\feal \\ is found to be 0.16 (for each \\fe \\ line), precisely at the level predicted by Times et al (1995) on the basis of Woosley and Weaver (1995) nucleosynthesis calculations.  This finding of RHESSI appears as an impressive confirmation of a theoretical prediction. However, more recent studies of SNII nucleosynthesis produce different values for the \\feal \\ ratio (see next section), considerably higher than the one of Timmes et al. (1995). Combined with the RHESSI finding, the new theoretical results call for a reassessment of the \\al \\ sources in the Milky Way. In this work we discuss those results and their implications. We argue that none of the proposed sources of \\al \\ satisfies all observational constraints at present. Forthcoming observations by the INTEGRAL satellite, combined with a new generation of stellar nucleosynthesis models (for rotating massive stars up to 100 \\Ms \\ and metallicities up to 3 \\Zs) will probably be required to settle the issue. ",
        "conclusions": "Contrary to a rather widely spread opinion, the recent RHESSI detection of radioactive \\fe \\ in the Milky Way does not imply that \\al \\ is mostly produced by supernova explosions. Recent theoretical results suggest that the \\fe \\ line flux would then be close to the one of \\al \\ (within a factor of two). Assuming that both the RHESSI results and the recent stellar nucleosynthesis results hold, another source of \\al \\ should be found.  Wolf-Rayet stars appear as natural candidates, in view of their absolute \\al \\ yields  (at least in the framework of the Geneva models: either with high mass loss rates and  no rotation - Meynet et al. 1997 - or with mild mass loss rates and rotation - Vuissoz et al. 2003) and  presumably low \\feal \\ ratios. However, the strong dependence of the \\al \\  yields on metallicity suggests that the \\al \\ emissivity should be steeply increasing in the inner Galaxy, while the COMPTEL observations clearly display a milder enhancement at small Galactic longitudes.  Thus, almost twenty years after its discovery (Mahoney et al. 1982), the \\al \\ emission of the Milky Way has not yet found a completely satisfactory explanation. Indeed, the recent observational (COMPTEL, RHESSI) and theoretical (RHHW02, LC03, Vuissoz et al. 2003) results have made the puzzle even more complex than before. The solution will obviously require progress in both directions. From the theory point of view, detailed nucleosynthesis calculations of mass losing and rotating stars up to the final explosion in the mass range 12-100 \\Ms \\ and for metallicities up to 3 \\Zs \\ will be required ; furthermore, the uncertainties still affecting the reaction rates of $^{22}$Ne($\\alpha$,n) (major neutron producer during He burning in massive stars) and $^{59}$Fe(n,$\\gamma$) will have to be substantially reduced. From the observational point of view, the radial distributions of both \\al \\ and \\fe \\ will be needed; such distributions will probably be available if the operation of ESA's INTEGRAL satellite is prolonged for a few years beyond its nominal 2-year operation."
    },
    "0402/gr-qc0402102_arXiv.txt": {
        "abstract": "By means of a highly accurate, multi-domain, pseudo-spectral method, we investigate the solution space of uniformly rotating, homogeneous and axisymmetric relativistic fluid bodies. It turns out that this space can be divided up into classes of solutions. In this paper, we present two new classes including relativistic core-ring and two-ring solutions. Combining our knowledge of the first four classes with post-Newtonian results and the Newtonian portion of the first ten classes, we present the qualitative behaviour of the entire relativistic solution space. The Newtonian disc limit can only be reached by going through infinitely many of the aforementioned classes. Only once this limiting process has been consummated, can one proceed again into the relativistic regime and arrive at the analytically known relativistic disc of dust. ",
        "introduction": "Self-gravitating bodies of constant density have always played a central role in the physics of gravitation. Contributions that have been most significant to the aspects of this subject that will be relevant to this paper can be divided into (1) work done within Newton's theory of gravitation: \\cite{Maclaurin}, \\cite{Poincare}, \\cite{Dyson92,Dyson93}, \\cite{Lichtenstein}, \\cite{Wong74} and \\cite{Eriguchi81} and (2) work done within Einstein's theory of gravitation: \\cite{Schwarzschild}, \\cite{Chand67}, \\cite{Bardeen71}, \\cite{BI76} and \\cite{GonGour02}.  Despite this great investment of effort, it has not yet been possible to explore the whole spectrum of Newtonian, let alone relativistic solutions, even if one restricts oneself to axial symmetry and stationarity. Using sophisticated computer programs, we believe that a great step can be taken toward painting the complete relativistic picture of uniformly rotating, homogeneous and axisymmetric bodies. However, because of the many intricate details in this picture, particularly in the vicinity of limiting configurations, it is necessary to use a computer program robust enough to be able to render such limiting configurations and accurate enough to be able to distinguish between neighbouring solutions of Einstein's equations. For our investigations of homogeneous fluids, we used a code based on multi-domain pseudo-spectral methods \\citep{AKM1,AKM3}, which satisfies these requirements. For a comparison with other codes, see \\cite{Stergioulas}.  It is our intention in this paper to present the relativistic picture in its entirety by describing those parts of it that we have studied explicitly and augmenting them with conjectures as to the rest. With this in mind, we focus our attention neither on the numerical methods nor primarily on properties of individual configurations (as in \\citealt{AKM4,SA}), but instead emphasize the interrelations between various configurations and portray the solution space as a whole. As an interesting example of the newly explored configurations, we present the shape and various physical parameters of a core-ring solution.  The picture that emerges for relativistic homogeneous fluids contains familiar demarcations, which we can use to orient ourselves. It contains, for example, three analytically known solutions: the (inner and outer) Schwarzschild solution \\citep{Schwarzschild}, the relativistic disc of dust solution \\citep{BW71,NM95,NM03} as well as the Maclaurin solution \\citep{Maclaurin} as one of its Newtonian limits. This last solution will be a particularly useful point of departure for describing the corresponding relativistic picture. It represents homogeneous, rotating fluid spheroids and, but for a scaling factor, depends for given mass density on only one parameter, thereby allowing one to refer to the ``Maclaurin sequence''. After introducing some basic equations in Sec.~\\ref{Equations} we thus turn our attention to the Maclaurin solution and its post-Newtonian extension in Sec.~\\ref{Maclaurin}.  It turns out that not all relativistic configurations are connected to each other in a continuous way, and it will be useful to introduce the notion of a class of solutions. A given class will be defined to include all configurations of strictly positive mass that are connected continuously to each other. In Sec.~\\ref{Classes} we review the characteristics of the ``first'' three classes in detail and provide an overview of the remaining classes.  We close with a discussion of the limitations of numerical methods and some thoughts on the completeness of the relativistic picture painted in this paper. ",
        "conclusions": "Numerical evidence suggests that as one proceeds to flatter configurations (higher classes), the departure from the Newtonian sequence grows small. This is reflected in the fact that boundaries associated with highly relativistic attributes such as infinite central pressure or the transition to a Black Hole are presumably absent from all classes upwards of class~II and is a result of excluding many-body configurations. If one were to abandon this restriction and choose some (arbitrary) relationship for the rotation of the various segments, then it would again be possible to reach ``relativistic boundaries'' most likely. With this restriction in place, however, the only path to the disc of dust passes through an infinite number of bifurcation points and sees the configurations growing ever flatter and more and more corrugated until one lands necessarily at the Newtonian Maclaurin disc. Only once this limit has been reached can one proceed again into the relativistic regime, i.e.\\ the relativistic disc of dust \\citep{NM95} and indeed reach the extreme Kerr Black Hole. This picture reinforces Bardeen's speculations \\citep{Bardeen71}, who wrote that the singularities in the post-Newtonian expansion at the bifurcation points ``may forbid the existence of any highly relativistic, highly flattened, uniformly rotating configurations which are simply connected.'' He went on to conclude that the ``only acceptable model of an infinitesimally thin, uniformly rotating disk in general relativity may be one that is made up of an infinite number of disjoint rings.''  The complicated and highly non-linear set of equations describing axially symmetric, stationary, uniformly rotating fluids of constant density can only be solved approximately or numerically. As such, it is difficult to find strict results describing the solution set. It is conceivable, for example, that there exist solutions to these equations that possess no Newtonian limit (and must thus be disconnected from the classes of solutions introduced here). It is also conceivable that Newtonian solutions exist that have not yet been found, but serve as limits to relativistic solutions.  Although our numerical considerations do not preclude the possibility of the existence of undiscovered solutions, we presume to conjecture that the classes presented here cover the entire solution space. Attempts to prove or disprove this conjecture are bound to lead to innovative insights into the nature of Einstein's equations and novel methods for probing their structure. The modification to the picture drawn in this paper resulting from a change in the equation of state will be discussed elsewhere."
    },
    "0402/astro-ph0402145_arXiv.txt": {
        "abstract": "In this paper we analyze the behavior of Galactic, LMC and SMC Cepheids in terms of period-color (PC) and amplitude-color (AC) diagrams at the phases of maximum, mean and minimum light. We find very different behavior between Galactic and Magellanic Cloud Cepheids. Motivated by the recent report of a break in LMC PC relations at 10 days (Tammann et al. 2002), we use the F-statistical test to examine the PC relations at mean light in these three galaxies. The results of the F-test support the existence of the a break in the LMC PC(mean) relation, but not in the Galactic or SMC PC(mean) relations. Furthermore, the LMC Cepheids also show a break at minimum light, which is not seen in the Galactic and SMC Cepheids. We further discuss the effect on the period-luminosity relations in the LMC due to the break in the PC(mean) relation. ",
        "introduction": "\\citet[][hereafter SKM]{sim93} used hydro-dynamical models to explain the observations of \\citet{cod47}: Galactic Cepheids show a spectral type independent of period at maximum light and a spectral type at minimum light that gets later as the period increases. SKM computed radiative hydro-dynamical models of Galactic Cepheids which agreed with Code's observations. SKM interpreted these observational phenomena as being due to the location of the photosphere relative to the hydrogen ionization front. They further used the Stefan-Boltzmann law applied at maximum and minimum light to show that  \\begin{eqnarray} \\log T_{max} - \\log T_{min} = {1\\over{10}}(V_{min} - V_{max}), \\end{eqnarray}  where $T_{max}$ and $T_{min}$ refer to the effective temperature at maximum and minimum light, respectively. Thus, higher optical amplitudes are associated with higher temperature amplitudes, which are in turn related to higher values of $T_{max}$ and/or $T_{min}$. For this study we do not assert a causal relation between higher temperatures and higher amplitudes preferring to refer to these quantities as being ``associated''. If, for some reason, either $T_{max}$ or $T_{min}$ does not increase as the amplitude increases, equation (1) predicts a relationship between amplitude and $T_{min}$ or $T_{max}$ respectively. SKM used data from \\citet{pel76} and \\citet{mof80,mof84} to show that Galactic Cepheids are such that higher amplitude stars are driven to cooler temperatures at minimum light. This, according to equation (1), is because the range of temperatures at maximum light is independent of period for a large range of periods. So the form of the period-color (PC) relation at maximum light is related to the form of the amplitude-color (AC) relation at minimum light, and vice versa.  \\citet{tam03} used Galactic and OGLE LMC/SMC Cepheid data to show that there is a difference in the PC relations in these three galaxies. Furthermore, \\citet{tam02a} and \\citet{tam02} show the existence of two PC relations, one for short ($P<10$ days) and one for long ($P>10$ days) period Cepheids, in the LMC. Motivated by this and the presence of high quality Magellanic Cloud Cepheid data from the OGLE project \\citep{uda99a}, we decided to investigate the properties of Magellanic Cloud Cepheids in terms of their PC and AC relations at maximum, mean and minimum light.  In addition to the two major reasons mentioned above, we list a number of other arguments motivating the present study:  \\begin{enumerate} \\item   Since mean light is just that - the average over a range of values - interesting properties of Cepheids at mean light are due to the average of these properties at all pulsation phases. By investigating the phases of maximum and minimum light, we are studying those phases of stellar pulsation which contribute to the observed properties at mean light. Our interest lies in understanding breaks in the LMC mean light Cepheid PC and period-luminosity (PL) relations at 10 days reported by \\citet{tam02a}. Let $y_{ij}$ be the (absolute) magnitude of Cepheid variable stars, $i=1,...,n_{star}$ at the $j^{th}$ phase ($j=1,...,N$) during a pulsation period $P_i$. Then we can formulate a PL relation at a particular phase as,  \\begin{eqnarray} y_{ij} = a_j + b_j \\log (P_i), \\end{eqnarray}  where $a_j,b_j$ are the unknown coefficients as a function of the phase $j$. If we define $y_i$ as ${\\sum_{j=1}^{j=N}y_{ij}}/N$, the average over the pulsation period, it is easy to show that $y_i = a_m + b_m \\log (P_i),$ where $a_m,b_m$ are the average over phase of the $a_j,b_j$ in equation (2)\\footnote{Note that $b_m$ may not lie in between $b_{max}$ and $b_{min}$, the slopes at maximum and minimum light, respectively. A similar conclusion also holds for the zero-point, $a_m$.}. This will be true if $y$ is measured in intensities and then the intensity mean converted to magnitudes or if $y$ is measured in magnitudes. Of course the magnitude mean and intensity means are in general not equal to each other but the difference is small ($\\sim0.03mag.$) and constant over a wide range of periods (for example, see \\citealt{gie98}). Consequently, some insight into the mean light relation can be gained by studying PL relations at various phases, e.g. at maximum and minimum light. Since the PC relation affects the PL relation (see, e.g.,  \\citealt{mad91} for the basic physics of PC and PL relations), it is of interest to study the PC relation at various phases. Furthermore, the maximum and minimum light are closely associated with the more interesting phases of stellar pulsation: the expansion/contraction velocity is close to its maximum value when the photosphere is passing through the mean radius of the star (see, e.g., \\citealt{mih03} for the details).  \\item   The amplitude is a fundamental observational and theoretical quantity in stellar pulsation. Kanbur and Ngeow (2004, in-preparation) show that the amplitude is a very good descriptor of light curve shape, and the V band amplitudes are correlated to the first Principal Component ($PC1$) of the light curve. In addition, \\citet{kan02} demonstrated that $PC1$ can explain over 90$\\%$ of the variation in light curve shape. Thus the V band amplitude can be taken to be a good descriptor of V band light curve shape. A similar conclusion holds for the I band. Because the optical brightness fluctuations of Cepheids are predominantly due to temperature fluctuations \\citep{cox80}, it is thus instructive to examine AC diagrams at maximum and minimum V band light. Furthermore, since AC relations are related to PC relations through equation (1), their study can serve as a useful complement to strengthen any conclusions reached using PC relations.  \\end{enumerate} ",
        "conclusions": ""
    },
    "0402/astro-ph0402373_arXiv.txt": {
        "abstract": "We present global 2D and 3D simulations of self-gravitating magnetized tori. We used the 2D calculations to demonstrate that the properties of the MRI are not affected by the presence of self-gravity: MHD turbulence and enhanced angular momentum transport follow the linear growth of the instability. In 3D, we have studied the interaction between an $m=2$ gravitational instability and MHD turbulence. We found its strength to be significantly decreased by the presence of the latter, showing that both instabilities strongly interact in their non-linear phases. We discuss the consequences of these results. ",
        "introduction": "In the early phases of star formation, the forming accretion disks are likely to be very massive because of a strong infall from the parent molecular cloud. These massive disks are subject to the development of gravitational instabilities which transport angular momentum outward (Laughlin, Korchagin, \\& Adams 1998). In addition, when sufficiently ionized, the disks are also unstable to the MRI (Balbus \\& Hawley 1991, 1998). The simultaneous development of these instabilities in disks may significantly affect their evolution. We have undertaken a study of self-gravitating magnetized tori by means of numerical simulations.  In section 2, we describe the numerical methods we have used. In section 3, we present the results of the 2D simulations, focusing on the properties of the MRI in a self-gravitating environment. In section 4, we review preliminary results obtained in 3D calculations, and we discuss the implications of our results in section 5. ",
        "conclusions": "We have presented here the first global simulations of massive, magnetized disks.  Using 2D axisymmetric numerical simulations, we have shown that the properties of the MRI are similar in self--gravitating and zero mass disks.  We observe that these disks quickly develop a dual structure composed of an inner thin disk in Keplerian rotation around the central mass, and a thicker outer torus whose rotation profile is close to a Mestel profile.  We have then used 3D simulations to study the angular momentum transport properties in disks when both MHD turbulence and gravitational instabilities are present.  We have found that the gravitational instability is affected by the presence of the turbulence: the gravitational stress tensor is decreased by roughly a factor of $2$ when compared with hydrodynamical simulations.  This results in a smaller mass accretion rate toward the central object. Self--gravitating disks may therefore have a longer lifetime than previously thought.  Note that the simulations presented here use an adiabatic equation of state. In this case, all the energy generated in shocks and compression is locally converted into heat.  This prevents the formation of bound objects by gravitational collapse.  The opposite case would correspond to the use of a locally isothermal equation of state, for which all the energy generated is immediately radiated away.  Several authors (Mayer et al. 2002, Rice et al. 2003, Boss 1997 \\& 1998) have indeed reported gravitational collapse in isothermal calculations of disks, although this issue is still under debate (Pickett et al. 2003).  We believe that the presence of MHD turbulence affects the thermal balance in the disk and therefore needs to be included.  We are currently performing calculations of isothermal massive and magnetized disks."
    },
    "0402/astro-ph0402003_arXiv.txt": {
        "abstract": "The CH star CS 31062-050 ($\\lbrack$Fe/H$\\rbrack=-2.42$) is one of the most useful stars yet discovered for evaluating the s-process in metal-poor stars. It is very abundant in heavy elements (e.g., $\\lbrack$La/Fe$\\rbrack=2.2$), and its relatively cool temperature and low gravity mean that there are many lines of interesting elements present in the spectrum. We measured the abundances of 22 elements with Z$\\geq$29, including the rarely measured Lu and Pd. We derive an upper limit on the Th abundance as well. The abundances in CS 31062-050 show a similar pattern to many other metal-poor CH stars: high $\\lbrack$Pb/Fe$\\rbrack$ and $\\lbrack$Pb/La$\\rbrack$ ratios, low $\\lbrack$Y/La$\\rbrack$ ratios and high $\\lbrack$Eu/La$\\rbrack$ values compared to the solar system s-process. However, the Th limit, with additional assumptions, is not consistent with the idea that the excess Eu in CS 31062-050 is contributed by the r-process. In addition, the observed $\\lbrack$Eu/Tb$\\rbrack$ cannot be explained by any ratio of solar-system s-process and r-process abundances. We therefore argue that the abundance pattern in CS 31062-050 is most likely the result of the s-process, and we discuss possible modifications that could explain the non-solar-system pattern observed. ",
        "introduction": "The sample of identified field stars with [Fe/H]\\footnote{We use the usual notation [A/B]$\\equiv {\\rm log}_{10}(N_A/N_B)_* - {\\rm log}_{10}(N_A/N_B)_{\\odot}$ and log$\\epsilon(\\rm A) \\equiv {\\rm log}_{10}(N_A/N_H)+12.0$. A/B$\\equiv N_A/N_B$.}$<-2.4$ has increased by more than an order of magnitude in the last decade. Because the elements in the atmospheres of these stars have been produced in a small number of nucleosynthetic events, abundance determinations can provide direct tests of model yields from different nuclear processes.  Even in the early work on very metal-poor stars it quickly became clear that there were subclasses of objects with very high [heavy-element/Fe] ratios that could be traced to specific nucleosynthetic origins (e.g., McWilliam \\etal{} 1995).  The elements heavier than the iron peak are made through neutron capture via two principal processes: the r-process and the s-process (Burbidge \\etal{} 1957). The r-process (for rapid process) occurs when neutrons are added much more rapidly than the $\\beta$ decay times of the relevant nuclei.  The site or sites of the r-process are not known, although suggestions include the $\\nu$-driven wind of Type II SNe (e.g., Woosley \\& Hoffman 1992; Woosley \\etal{} 1994) and the mergers of neutron stars (e.g., Lattimer \\& Schramm 1974; Rosswog \\etal{} 2000). The s-process occurs when neutrons are added more slowly and $\\beta$ decays, changing neutrons to protons, keep the nuclei from straying far from the valley of $\\beta$ stability. The He intershell in asymptotic giant branch (AGB) stars is the site of the s-process as well as of C production. C and heavy elements are brought to the surface of the AGB stars during the ``third dredgeup''. Since the r-process produces very neutron-rich nuclei initially, the r-process reaches the neutron magic numbers with considerably fewer protons than the s-process, and, as a result, these two processes produce abundance peaks at different atomic weights. As a result, when the solar-system total abundances (\\tss) are separated into contributions from the s-process (\\sss) and the r-process (\\rss) (e.g., K\\\"appeler, Beer, \\& Wisshak, 1989; Arlandini \\etal{} 1999), some elements are mostly contributed by the r-process, such as Eu, and some by the s-process, such as Ba and La. Therefore Eu is commonly referred to as an ``r-process element'' and Ba and La as ``s-process elements''.  The [Eu/Ba] and the [Eu/La] values are used to estimate the ratio of r-process to s-process contributions to the heavy element abundances in a star, increasing as the r-process fraction increases.  Despite the nomenclature, it is important to remember that all neutron-capture elements lighter than Z=84 are made in both processes (Clayton \\& Rassbach 1967). Th (Z=90) and U (Z=92) can only be made in the r-process.  The surveys of metal-poor stars have uncovered stars rich in C and s-process elements and stars rich in r-process elements. These have been the subject of many follow-up studies, because of the insight they can provide on the production of the heavy elements in the early Galaxy.  \\subsection{The very metal-poor CH stars}  Metal-poor stars with enhanced abundances of C and s-process elements are called CH stars.  In an extensive survey, McClure (1984) and McClure \\& Woodsworth (1990) showed that all of the CH stars in their sample were members of binary systems with orbital parameters consistent with a white dwarf secondary. The explanation for the classic CH stars is that they result from the transfer of C, N, and s-process material produced in an AGB companion which is now a white dwarf.  The abundance patterns in CH stars for the heavy elements are therefore a very accessible means of empirically deriving s-process yields and for inferring the structure of, and physical conditions in, AGB stars. The traditional indicators of s-process material are super-solar [Ba/Fe] (the Pop I version of CH stars are often referred to as barium stars) and subsolar [Eu/Ba].  Theory predicts that s-process nucleosynthesis will depend on the initial metallicity of the AGB star and on its mass (Gallino \\etal{} 1998; Goriely \\& Mowlavi 2000; Busso \\etal{} 2001).  Busso \\etal{}  reported s-process yields as a function of initial [Fe/H] and for AGB stars with 1.5 and 3.0M$_\\odot$ and made comparisons with the available observational data. The basic metallicity dependence is a tilting of the s-process products toward heavier elements with decreasing [Fe/H] of the host AGB star.  The prediction that $^{208}$Pb will have a particularly strong excess has been verified for CS 22183-015 (Johnson \\& Bolte 2002a), HE 0024-2523 (Lucatello \\etal{} 2003) and many of the stars studied by Van Eck \\etal{} (2001, 2003) and Aoki \\etal{} (2002) (A02). On the other hand, Aoki \\etal (2000, 2001) present the analysis of two metal-poor, s-process-rich stars, LP 625-44 and LP 706-7, which have [Pb/Ba] $\\sim$ 0.  \\subsection{Very metal-poor r-process-rich stars} There is a second class of neutron-capture-rich very metal-poor stars in which the abundance pattern of the heavy elements more closely follows that inferred for the r-process elements in the solar system. CS 22892-052 (Sneden \\etal{} 1996, 2000), HD 115444 (Westin \\etal{} 2000) and CS 31082-001 (Cayrel \\etal{} 2001) are the best studied members of this class. Studies have showed a remarkable similarity in the abundance ratios for elements between Ba (Z$=56$) and Hf (Z$=72$), (c.f. Truran \\etal{} 2002) although in the lighter and heavier r-process element peaks there is considerable star-to-star scatter in abundance ratios (e.g., Sneden \\etal{} 2000; Johnson \\& Bolte 2002b; Hill \\etal{} 2002). Some r-process-rich stars, in particular CS22892-052, are also C-rich. The source of that C is unknown. Finally, we note that very metal-poor stars with low [neutron-capture/Fe] ratios have abundance patterns between Ba and Hf that agree with \\rss{} (Sneden \\& Parthsarathy 1983; Gilroy \\etal{} 1988; Johnson \\& Bolte 2001).   \\subsection{Some puzzles and challenges}  As more detailed studies of very metal-poor CH stars became available, some stars and elements did not fit neatly into the picture described above. Despite having lower [Eu/Ba] than the r-process-element-rich  stars, in some CH stars [Eu/Ba] is higher than that \\sss{} or than predicted for the metal-poor s-process. Hill \\etal{} (2000) analyzed spectra of two CH stars, CS 22948-027 and CS 29497-034, and concluded that the observed abundances could not be fit by either a scaled solar-system s-process (\\sss) or scaled solar-system r-process (\\rss), but instead reflected enrichment by both processes.  A02 and Johnson \\& Bolte (2002a) also noted a large spread in [Eu/Ba] for other CH stars but suggested that the abundance pattern seen in these CH stars was due solely to the s-process, albeit one that produces a varying Eu/Ba ratio.  A large percentage of CH stars have Eu/Ba ratios larger than \\sss. With the Hill \\etal{} interpretation, this would suggest that a number of s-process-rich stars are also r-process rich. Some fraction of non-CH field stars are Eu-rich, so it would not be surprising to find some stars that began as r-process-rich and also were polluted by an AGB companion later in their lives. However, while seven of 32 non-CH stars in McWilliam \\etal{} (1995) have [Eu/Fe]\\gtsim 0.5, six of eight CH stars in A02 have such high Eu. So it appears that r-process enrichment cannot be the solution for {\\it all} the high Eu CH stars (but see below).  Recently, Cohen \\etal{} found an extreme example of non-\\sss{} ratios in the CH star HE 2148-1247. The measured Ba/Eu ratio was $\\sim 100$ while theoretical calculations from Arlandini \\etal{} (1999), for example, give Ba/Eu $\\sim$ 640. They favored the addition of r-processed material as well as s-processed material to HE 2148-1247. Qian \\& Wasserburg \\etal{} (2003), in a companion paper, proposed an intriguing theory for the creation of such ``s+r''-process stars. First some s-processed material is accreted from the AGB companion, which turns into a white dwarf. Later in the evolution of the system, the white dwarf accretes matter from the polluted star and suffers an accretion-induced collapse (AIC) to a neutron star. The $\\nu$-driven wind produces an r-process, which also pollutes the companion, but since the white dwarf lacks an H or He envelope, lighter elements, such as Fe, are not manufactured. Therefore, the remaining star is r-process {\\it and} s-process-rich. The AIC could potentially deliver a strong kick to the neutron star, which could explain why some CH stars have recently been shown not to be members of binaries (Preston and Sneden 2001; Hill \\etal{} 2000).  Because the production of the s-process in the AGB star and the production of the r-process in the AIC are connected to the same binary companion, this alleviates some of the concern expressed by A02 and Johnson \\& Bolte (2002) about the high frequency of potential s$+$r stars. However, at least 50\\% of the very metal-poor CH stars would need to be s$+$r stars if that is the explanation for the elevated Eu/Ba ratios (see Figure 13 in Cohen \\etal{} (2003)). As Qian \\& Wasserburg point out, the frequency of AIC events and the parameters necessary to create them are unknown at this time, but the large fraction of very metal-poor CH stars with high Eu implies that the ability to pollute a companion as a AGB star must be tightly correlated with an r-producing AIC event. Another possible problem these authors mention is the still uncertain nucleosynthesis in AIC, which may or may not produce the r-process.  If we hope to understand the origin of neutron-capture elements in CH stars, we need to measure the abundances of as many heavy elements as possible to see if they are consistent with an s$+$r, or s-only or r-only scenario. In this paper,  we present results for another CH star. Because of the combination of large abundance enhancements and atmospheric parameters in this star, we can measure accurate abundances for a number of elements that have been infrequently studied in CH stars. These include Pd, Tb, Ho, Tm and Lu. We have also been able to put interesting limits on the abundance of Th. ",
        "conclusions": "We have determined the abundances of 20 neutron-capture elements in CS 31062-050. This is a C-rich star with several ratios of neutron-rich elements that match models of the metal-poor s-process. However, several well-measured heavy element ratios, most notably Eu/La, are inconsistent with any s-process predictions and the two ratios Eu/Tb and Eu/Dy cannot be explained by any combination of any s-process predictions and r-process predictions. High Eu/La (in the context of the s-process) has been seen before in CH stars leading some authors to suggest CH stars that are both s- and r-process-element rich. For CS 31062-050, our upper limit to the Th abundance argues against the r-process as the complete explanation, at least in this star. The Eu/Tb and Eu/Dy ratios suggests that our understanding of the s- or r-process yields are incomplete for at least some of these elements. An s-process that produces variable amounts of Eu in excess of the current models could eliminate the most glaring differences between observations and models."
    },
    "0402/astro-ph0402235_arXiv.txt": {
        "abstract": "We use redshift observations of two deep 1.4~GHz fields to probe the evolution of the bright end of the radio galaxy luminosity function to $z=1.5$. We show that the number of galaxies with radio power that would correspond locally to an ultraluminous infrared galaxy (ULIG) evolves rapidly over this redshift range. The optical spectra and X-ray luminosities are used to refine this result by separating the sources with clear active galactic nucleus (AGN) signatures from those that may be dominated by star formation. Both populations show extremely rapid evolution over this redshift range. We find that the number of sources with ULIG radio power and no clear AGN signatures evolves as $(1+z)^{7}$. ",
        "introduction": "\\label{secintro}  Star-forming ultraluminous infrared galaxies, or ULIGs ($L_{FIR}>10^{12}~L_\\odot$ from \\markcite{sanders96}Sanders \\& Mirabel 1996; here we mean galaxies without clear signatures of active galactic nucleus [AGN] activity in their spectra, even if there may be some AGN contribution to their luminosities), represent the most luminous tip of the star-forming galaxy population. Locally, ULIGs are very scarce (\\markcite{kim98}Kim \\& Sanders 1998), but if, as is widely believed, submillimeter galaxies are the high redshift analogs, then by $z=1$ they are the dominant contributors to the star formation history (e.g., \\markcite{lilly99}Lilly et al.\\ 1999; \\markcite{bcr00}Barger, Cowie, \\& Richards 2000; \\markcite{gispert00}Gispert, Lagache, \\& Puget 2000; \\markcite{chapman03}Chapman et al.\\ 2003). It is extremely hard to trace the evolution of the ULIG population at low redshifts ($z<1.5$) because of the difficulties of mapping large areas with current submillimeter detectors and of identifying the optical counterparts to low spatial resolution submillimeter or far-infrared (FIR) observations.  Ultradeep decimetric radio surveys can play a major, complementary role. The well-known tight correlation between FIR luminosity and radio power in local star-forming galaxies and radio-quiet AGNs (\\markcite{condon92}Condon 1992) means that we can identify ULIGs out to substantial redshifts based on their radio power. For a 1.4~GHz sensitivity limit of 40~$\\mu$Jy, an $L_{FIR}=10^{12}~L_\\odot$ source will be detected to $z\\sim 1.5$. Thus, sources with radio power that would correspond locally to a ULIG may be seen directly to large redshifts. Since the ultradeep 1.4~GHz samples cover large fields, do not suffer from extinction, and provide subarcsecond positional accuracy, it is easy to develop large samples with highly complete optical counterpart identifications. By contrast, optically selected samples may omit dusty sources at these redshifts, while X-ray samples will select AGN dominated sources.  Current decimetric surveys do not probe the normal star-forming galaxy populations at $z=1$. To the deepest 1.4~GHz sensitivity limits of $\\sim 20$~$\\mu$Jy ($5\\sigma$), a source with radio power corresponding to $L_{FIR}=10^{10}~L_\\odot$ would only be seen to $z\\sim 0.3$. Thus, attempts to provide maps of the entire star formation history over the $z=0-1$ range based on decimetric observations (e.g., \\markcite{cram98}Cram 1998; \\markcite{mobasher99}Mobasher et al.\\ 1999; \\markcite{haarsma00}Haarsma et al.\\ 2000, hereafter H00) have to rely on model luminosity functions to extrapolate the measured data and obtain star formation rates. In these descriptions, much of the star formation lies in sources that are not directly measured by the observations (e.g., H00).  We use ultradeep 1.4~GHz observations of the Hubble Deep Field-North (HDF-N; \\markcite{richards00}Richards 2000) and SSA13 (\\markcite{fomalont04}Fomalont et al.\\ 2004) fields to analyze the evolution of the radio luminosity function (LF) over the range $z=0-1.5$. Our sample of 346 radio sources represents an increase of nearly a factor of five over H00 and enables us to provide a much more accurate determination of the $z=1$ radio LF. We use optical spectral classifications and X-ray characteristics to separate AGN dominated sources from those that may be star formation dominated and to estimate the number density of ULIGs. An analysis of the submillimeter and far infrared properties of the sample will appear in Wang, Barger, and Cowie 2004 in preparation. We assume $\\Omega_M={1\\over 3}$, $\\Omega_\\Lambda={2\\over 3}$, and $H_0=65$~km~s$^{-1}$~Mpc$^{-1}$. ",
        "conclusions": "\\label{secop}  We constructed the LFs for the 1.4~GHz power over the range $z=0.5-1.5$ using the $1/V_a$ method (e.g., \\markcite{felten77}Felten 1977). This is shown for the total spectroscopic sample (AGNs and star formers; {\\it triangles}) in Figure~\\ref{fig3}a, where it is compared with the local measurements of \\markcite{condon89}Condon (1989; {\\it solid line}). To investigate the effects of incompleteness, we computed the maximum total LF ({\\it squares}) for just the HDF-N using the much more complete combined photometric and spectroscopic samples and placing all sources without a photometric or a spectroscopic redshift into the $z=0.5-1.5$ bin. Many of these probably lie at higher redshifts, so this is an upper bound. Finally, we assumed that the HDF-N sample was only 50\\% complete in the $40-80$~$\\mu$Jy range, which approximates the completeness corrections computed by \\markcite{richards00}Richards (2000). The maximum correction for incompleteness is a factor of 2.6 at the lower radio powers but is relatively small at the higher radio powers.  In order to parameterize the evolution, we assumed that the local portion of the LF allocated to star formers in \\markcite{condon89}Condon (1989) undergoes pure luminosity evolution and that allocated to AGNs undergoes pure number density evolution. We used these assumptions to calculate the expected LF in the $z=0.5-1.5$ interval. (The use of the S02 or the MG01 measurements, which were constructed using more similar methodologies to ours for classifying the galaxies, does not change the diagrams significantly.) This simple model provides a reasonable fit to our measured LF, with the AGN number density evolving as $(1+z)^{3}$, and the star-forming luminosities evolving as either $(1+z)^{3}$ ({\\it dotted line}), if we match the observed LF, or as $(1+z)^{3.8}$ ({\\it dashed line}), if we match the maximum total LF. This systematic uncertainty is much larger than the statistical errors and primarily affects the evolution of the star-forming LF because the uncertainties are larger at the lower luminosities where this term dominates. Our results for the evolution of the star-forming LF are consistent with those determined from local samples by MG01 and \\markcite{condon02}Condon et al.\\ (2002), who found a similar evolution of $(1+z)^{3\\pm1}$. Our results are also broadly consistent with those of H00, who found a best pure luminosity evolution model with $(1+z)^{2.9}$.  In Figures~\\ref{fig3}b and \\ref{fig3}c we show the measured LFs in two narrower intervals ($z=0.6-0.9$ and $z=0.9-1.4$, respectively, chosen to match H00) and by spectroscopic class ({\\it diamonds}---star formers; {\\it squares}---AGNs). For the star formers, we also show the results of H00 ({\\it triangles}), which, within the uncertainties, agree with the present measurements. Our results are again broadly consistent with the evolution described above (see Fig.~\\ref{fig3}b caption for description), although the AGN LF is slightly high in the $z=0.6-0.9$ interval. This suggests that the power law evolution model may be too simple and that the evolution of the AGN LF is faster at low redshifts and slower at high redshifts. The star formers begin to show a more extended high luminosity tail in the $z=0.9-1.4$ interval.  \\begin{inlinefigure} \\psfig{figure=f3.eps,width=3.5in} \\vspace{6pt} \\figurenum{3} \\caption{ (a) Radio LF at $z=0.5-1.5$ for all the spectroscopically identified radio sources ({\\it triangles}), compared with the local determination from Condon (1989; {\\it solid line}). Log denotes the base 10 logarithm. The $1\\sigma$ uncertainties are Poissonian. Squares denote the maximum total LF. Vertical dotted line shows the equivalent radio power of an $L_{FIR}=10^{12}~L_\\odot$ source. Dotted curve shows the expected LF if the local star-forming LF undergoes pure luminosity evolution ($(1+z)^3$) and the local AGN LF undergoes pure number density evolution ($(1+z)^3$). Uncertainties are primarily systematic, and luminosities evolving as $(1+z)^{3.8}$ provide a better fit ({\\it dashed line}) to the maximum total LF. (b) Radio LF at $z=0.6-0.9$ for the star formers ({\\it diamonds}) and the AGNs ({\\it squares}), to be compared with the local star former ({\\it steep solid curve}) and AGN ({\\it shallow solid curve}) determinations from Condon (1989). Triangles show the measurements of H00 (with a small correction to make the geometries consistent) for the star formers. Dot-dashed curve shows the expected LF if the local star-forming LF undergoes pure luminosity evolution ($(1+z)^{3.8}$). Dashed curve shows the expected LF if the local AGN LF undergoes pure number density evolution ($(1+z)^{3}$). (c) As for (b), but for the $z=0.9-1.4$ redshift interval. \\label{fig3} } \\addtolength{\\baselineskip}{10pt} \\end{inlinefigure}  The rapid evolution of the radio LF implies that the number density of very luminous sources rises rapidly in the $z=0-1$ interval for both the star former and AGN classes though it may be begining to slow at redshifts $z>1$. In Figure~4 we show the number density of star formers with ULIG radio power ({\\it solid squares}) in three redshift intervals. The value from the  local star-forming radio LF is given by the open square. The evolution is reasonably well described by a $(1+z)^7$ evolution ({\\it dashed line}). This evolution is similar to that inferred by \\markcite{bcr00}Barger et al.\\ (2000) from a comparison of the local number density of ULIGs and near-ULIGs to the $z=1-3$ number density of $>6$~mJy radio/submillimeter sources ({\\it solid circle}). If most of the submillimeter sources are dominated by star formation (\\markcite{bcr00}Barger et al.\\ 2000; \\markcite{chapman03}Chapman et al.\\ 2003), the match to the radio evolution is reasonable, if the radio galaxies classified as star formers are indeed dominated by star formation and the FIR-radio correlation holds to these redshifts as suggested by recent results (Garrett 2002).  \\begin{inlinefigure} \\psfig{figure=f4.eps,angle=90,width=3.5in} \\vspace{6pt} \\figurenum{4} \\caption{ Number density of star formers with ULIG radio power in the redshift intervals $z=0.1-0.6$, $z=0.6-1.0$, and $z=1.0-1.4$ ({\\it solid squares}). Open square denotes the value from the local star-forming radio LF. Solid circles denote the local number density of ULIGs and near-ULIGs and the $z=1-3$ number density of $>6$~mJy radio/submillimeter sources (from Barger et al.\\ 2000). Dashed line shows a $(1+z)^7$ evolution. \\label{fig4} } \\addtolength{\\baselineskip}{10pt} \\end{inlinefigure}    We have given a precise determination of the high power end of the radio LF at high redshifts ($z\\sim 1$). The LF was shown to evolve rapidly with redshift, both for galaxies with AGN spectra and for those with only star-forming signatures. This result is consistent with model expectations based on the local LF and the radio number counts (e.g., \\markcite{condon89}Condon 1989). The number of sources with radio luminosities that would correspond to ULIGs (based on local normalizations) matches the observed evolution from the local ULIG population to the radio/submillimeter sources at $z>1$."
    },
    "0402/astro-ph0402553_arXiv.txt": {
        "abstract": "We analytically study how the behaviour of accretion flows changes when the flow model is varied. We study the transonic properties of the conical flow, a flow of constant height and a flow in vertical equilibrium and show that all these models are basically identical provided the polytropic constant is suitably changed from one model to another. We show that this behaviour is extendible even when standing shocks are produced in the flow. The parameter space where shocks are produced remain roughly identical in all these models when the same transformation among the polytropic indices is used. We present applications of these findings. ",
        "introduction": "Fully self-consistent study of any astrophysical system is generally prohibitive. Very often, for simplicity, it is necessary to construct models which have all the salient features of the original problem. However these models need not be unique. In the present paper we make a pedagogical review of three different models of rotating accretion flows and show that even though they are based on fundamentally different assumptions, they have identical physical properties. What is more, results of one model could be obtained from the other by changing a {\\it physical} parameter, namely, the polytropic constant. In other words, all these models are {\\it identical}.  Accretion disk physics has undergone major changes in the last fifty years. Bondi (1952) studied spherical accretion and found the existence of only one saddle type sonic point in an adiabatic flow. Later, the Keplerian disk model of Shakura and Sunyaev (1973) and thick disk model of Paczy\\'nski and Wiita (1980) the disk solutions became more realistic, though none of them was transonic, i.e., none was passing through any sonic point. Meanwhile, Liang and Thompson (1980) generalized this work for a flow which included angular momentum and discovered that there could be three sonic points. Matsumoto et al. (1984) tried to let the flow pass through the inner sonic point only and found that flow could pass through nodal type sonic points.  Chakrabarti (1989, 1990; hereafter referred to as C89 and C90 respectively) studied transonic properties  of accretion flows which are conical in shape in the meridional plane (`Wedge-shaped Flow') and also flows which are in vertical equilibrium. Subsequently, Chakrabarti \\& Molteni (1993) studied flows of constant height and also verified by time dependent numerical simulations that the flow indeed allows standing shocks in it. In a Bondi (1952) flow, to specify a solution one requires exactly one parameter, namely the specific energy ${\\cal E}$ of the flow. This is in turn determined by the temperature of the flow at a large distance. In an inviscid, rotating axisymmetric accretion flow, one requires two parameters, namely, specific energy ${\\cal E}$ and specific angular momentum $\\lambda$. Once they are specified, all the crucial properties of the flow, namely the locations of the sonic points, the shocks, as well as the complete global solution are determined. C89 numerically studied the properties of the parameter space rather extensively, and divided the parameter space in terms of whether standing shocks can form or not. In the present paper, we compare these models completely analytically and show, very interestingly, that one could easily `map' one model onto another by suitably changing the polytropic index of the flow. In other words, we show that these models are roughly identical to one another as far as the transonic properties go.  In the next Section, we present a set of equations which govern the steady state flow in all the three models. In \\S 3, we present the sonic point analysis and provide the expressions for the energy of the flow in terms of the sonic points. We observe that these expressions are identical provided there is a unique relation among the polytropic indices of these model flows. In \\S 4, we compare shock locations in all the three models. We also compare the parameter space which allows shock formation in these models with the regions obtained using purely numerical methods. In \\S 5, we show that in fact if the relations between the polytropic indices are used, the shock locations in all these models are also roughly identical. Consequently, the apparently disjoint parameter spaces drawn with the same polytropic index overlap almost completely when the aforementioned relations among polytropic indices is used. This remarkable behaviour shows underlying unity in these apparently diverse models. Finally, in \\S 6, we draw our conclusions. ",
        "conclusions": "In this paper, we discovered a unique relation among the polytropic indices of three different models of the axisymmetric accretion flows which ensures identical transonic properties in the sense that if all these models have the same conserved energies and angular momenta, then the sonic points also form exactly at the same place. When we proceeded further to compute the shock locations, we found that even the shocks form roughly at the same places. Apparently, disjoint parameter spaces for shock formation with the same value of polytropic index in three different models exhibits considerable overlap when the same unique relation (Eq. 8) was used. This shows that the models are virtually identical in properties and various disk models belong to one parameter family.  Our finding has given some insight into the relation between the nature of a flow with its equation of state. It seems that the relativistic equation of state in a flow in vertical equilibrium behaves similar to a roughly isothermal flow in a disk of constant height or in a conical flow. It is possible that in the latter models (Model C and H) the geometric compression is smaller and hence it is easier to keep them roughly isothermal while conserving energy as well.  Though our work has been mainly pedagogical, we believe that it could have several applications. For instance, linear and non-linear stability analysis and time dependent calculations (numerical simulations) are easier to perform when the disk is of constant thickness. Our work indicates that once certain properties regarding stability are established in one flow model, they would remain valid in other models as well provided the relation among the polytropic indices is incorporated.  The authors greatly acknowledge financial support from Department of Science and Technology through a Grant (No. SP/S2/K-14/98) with SKC."
    },
    "0402/astro-ph0402079_arXiv.txt": {
        "abstract": "{We deduce the globular cluster formation history of the nearby elliptical galaxy, NGC 5128, by using a chemical enrichment model to accurately reproduce its observed metallicity distribution function (MDF). We derive the observed MDF using recently obtained $U$ and $B$ photometry of the NGC 5128 GC system, with $(U-B)$ used as the metallicity indicator. Our results indicate that the GC system in this galaxy could be the product of two major GC formation episodes. The initial formation episode occured 11-12 Gyrs ago creating 65-75 percent of the mass in the GC system. This was followed by a second late formation episode which peaked 2-4 Gyrs ago and produced the remaining 25-35 percent of GC mass. ",
        "introduction": "Globular cluster (GC) systems have been popular tools in deciphering the star formation histories (SFHs) of galaxies. Their ubiquitous presence in galaxies of all morphological types, coupled with evidence that their creation seems to accompany major star formation episodes \\cite[e.g.][]{LR99} makes them useful tracers of galactic evolution \\cite[e.g.][]{KP98,VDB2000,Yoon2002}. In addition, GC systems are well represented by simple stellar populations (SSPs) with stars of the same age and chemical composition, which makes them easy to study using SSP models \\cite[e.g.][]{Yi2004}.  The exact formation mechanism of GCs has been the subject of much recent debate. Many models for GC formation have been proposed including gaseous mergers \\citep{Ashman92}, in situ formation \\citep[e.g.][]{Harris95}, multiphase collapse \\citep{Forbes97}, dissipationless hierarchical merging \\citep[e.g.][]{Cote98,Cote2000,Cote2002} and hierarchical clustering \\citep{Beasley2003}. While none of these models can be conclusively excluded, the widespread discovery of multimodal metallicity distributions in GC populations effectively rules out an extreme version of the monolithic collapse scenario for their formation \\citep[see e.g.][]{Forbes97}. Furthermore, the correlation between the mean metallicity of GCs and galaxy luminosity indicates that chemical enrichment of the GC system is intimately linked to the evolution of the host galaxy \\citep{Forbes96,Durrell96,Forbes2001,Cote2000,Jordan2004}. An elegant review of the GC formation models mentioned above can be found in \\cite{West2004}.  The object of this study is NGC 5128, the giant elliptical galaxy in the nearby Centaurus Group \\citep[see][for a comprehenive review of NGC 5128]{I98}. Located at a distance of approximately 3.6 Mpc \\citep[e.g.][]{Soria96}, NGC 5128 is the closest giant elliptical system with an estimated GC population of $1550 \\pm 350$ \\citep{Harris84}. It has been widely studied, not only because of its proximity and relative brightness, but also because it displays unusual physical features which suggest that this galaxy is a post-merger remnant. A prominent dust lane containing young stars and HII regions \\citep[e.g.][]{Unger2000,Wild2000} and a series of optical shells \\citep{Malin83}, which have HI \\citep{Schiminovich94} and molecular CO \\citep{Charmandaris2000} gas associated with them, are considered strong evidence that NGC 5128 underwent a recent merger event within the last $10^9$ years. Recently, \\citet{Rejkuba2004} suggested that star formation may have stopped as recently as 2 Myr ago in the north-eastern shell of NGC 5128.  In a series of major works, \\citet{Harris99}, \\citet{Harris2000} and \\citet{Harris2002} performed a comprehensive study of the metallicity distribution of stars in the inner and outer halo of NGC 5128. In this study we deduce the formation history of the GC system of NGC 5128 by accurately reproducing its observed metallicity distribution function (MDF) using a chemical enrichment model. We derive the observed MDF of this elliptical galaxy using recently obtained $U$ and $B$ photometry of 210 clusters in its GC system \\citep{Peng2004a}. The integrated $(U-B)$ colour is used as the metallicity indicator because it is sensitive to metallicity via the opacity effect but relatively insensitive to the effective main sequence turn-off temperature ($T_{eff}$) and therefore to age when $T_{eff} \\sim$ 7000-12000 K \\citep{Yi2004}. Similar techniques using $U$ band colours have been used by \\citet{rejkuba2001} and \\citet{jordan2002} who used $(U-V)$ and \\emph{Hubble Space Telescope} WFPC2 $(F336W-F547M)$, respectively. Although these colours are substantially better metallicity indicators than the previously used $(V-I)$ or $(B-V)$, they still change gradually with age and thus are not as effective as $(U-B)$ in determining metallicity \\citep{Yi2004}. Our work extends previous studies of NGC 5128 in a number of ways. The large number of confirmed GCs with $U$ band photometry makes it possible to derive statistically significant results. In addition, the relative robustness of $(U-B)$ as a metallicity indicator, compared to other optical colours, makes chemical enrichment an effective approach in modelling the formation history of the NGC 5128 GC system. The use of a consistent chemical enrichment code means that we can effectively transform metallicities into ages.  We begin by briefly checking that, as might be expected from previous theoretical and observational results, a single starburst followed by passive evolution (a monolithic scenario) is incapable of reproducing the observed MDF of the NGC 5128 GCs. Performing this check is not a redundant exercise because our modelling essentially yields \\emph{relative likelihoods} for various models to fit the observed MDF of NGC 5128. It is therefore instructive to compare the quality of fit between monolithic and non-monolithic scenarios. The main thrust of the paper, however, is to explore a \\emph{double starburst scenario}, analysing positions, timescales and relative strengths of the two star formation episodes that best explain the MDF of the NGC 5128 GC system. We provide a coherent picture of the formation history of NGC 5128 GCs based on double starburst scenarios that give excellent fits to the observed MDF and show that our results are consistent with the spectroscopic study of \\citet{Peng2004b}, who use the age-sensitive $H_{\\beta}$ index to age-date the GCs in this dataset. ",
        "conclusions": "We studied the formation history of the NGC 5128 GC system by using a chemical enrichment model to accurately reproduce its observed MDF derived from recently obtained $U$ and $B$ photometry. Our results show that the GC system in this galaxy could be the product of two major GC formation episodes, although we do not have adequate resolution in the model or in the observed MDF to make any reliable claim about the possibility of additional formation events.  The double starburst analysis produces high KS probabilities and therefore good fits to the observed MDF, with best-fit models (shaded region in Fig. \\ref{fig:relative_KS_values}) that favour an initial GC formation episode 11-12 Gyrs ago which produced 65-75 percent of mass in the GC system, with a second late formation episode approximately 2-4 Gyrs ago producing the remaining 25-35 percent. The late starburst results in a small fraction ($\\sim$ 5 percent) of the stellar mass in the GC system potentially having ages \\emph{less than 1 Gyr}. If NGC 5128 has over 1500 clusters, then we might expect a handful of those clusters to have ages less than 1 Gyr. We have not made any \\emph{a priori} assumptions about the driving mechanisms behind star formation episodes in our model. While a single starburst followed by passive evolution can clearly be discounted, the nature of our model admits any scenario where multiple episodes of GC formation are possible.  Our results are in general agreement with \\citet{Beasley2003}, who obtain good fits to the observed NGC 5128 MDF using a semi-analytical galaxy formation model \\emph{with metal-poor GC formation halted at $z \\sim 5$}. A quantitative similarity between this study and that of \\citet{Beasley2003} is that the initial GC formation episode (first gaussian starburst) in our model, which gave rise to the metal-poor GC subpopulation, decays rapidly to virtually zero (see Fig. \\ref{fig:best_sfr}) within $\\sim 2$ Gyrs of the formation of the galaxy; i.e. there is \\emph{effective truncation} of metal-poor GC production at high redshift ($z \\sim 3-4$). This closely resembles the truncation employed by \\citet{Beasley2003}, and the similarity is probably due to the fact that our model is comparable to the chemical enrichment prescription employed in the semi-analytical model of \\citet{Beasley2003}. Since the object of both studies is to reproduce \\emph{metallicity distributions} this resemblance is not unexpected.  This study would not be complete without a comparison of our results to those of \\citet{Peng2004b}, who performed a spectroscopic analysis of the GC dataset used in this paper. Based on the age-sensitive $H_{\\beta}$ index, they conclude that metal-poor GCs in NGC 5128 have ages comparable to those in the Milky Way, i.e. $\\sim 12$ Gyrs \\citep[e.g][]{Krauss2003}, and that the metal-rich GCs are consistent with a mean age of $5^{+3}_{-2}$ Gyrs. We thus find that our age estimates, derived using a chemical enrichment approach to exploit metallicity-sensitive photometric colours, are consistent with a study of the same objects using age-sensitive spectroscopic indices.  Our study demonstrates the potential of a chemical enrichment approach in deciphering the formation histories of the GC system in galaxies. As more age and metallicity-sensitive spectro-photometric data become available, methods such as the one used in this study will enable us to set robust constraints on the way galaxies either form or incorporate GCs, crucial to our understanding of galaxy formation."
    },
    "0402/astro-ph0402286_arXiv.txt": {
        "abstract": "{Today, Type~Ia supernovae are essential tools for cosmology, and recognized as major contributors to the chemical evolution of galaxies. The construction of detailed supernova progenitor models, however, was so far prevented by various physical and numerical difficulties in simulating binary systems with an accreting white dwarf component, e.g., unstable helium shell burning which may cause significant expansion and mass loss. Here, we present the first binary evolution calculation which models both stellar components and the binary interaction simultaneously, and where the white dwarf mass grows up to the Chandrasekhar limit by mass accretion. Our model starts with a $1.6$ $\\rm{M_\\odot}$ helium star and a $1.0$  $\\rm{M_\\odot}$ CO white dwarf  in a 0.124~day orbit. Thermally unstable mass transfer starts when the CO core of the helium star reaches $0.53 \\rm{M_\\odot}$, with mass transfer rates of $1\\cdots8 \\times 10^{-6}$ $\\rm{M_\\odot/yr}$. The white dwarf  burns the accreted helium steadily until the white dwarf mass  has reached $\\sim1.3$ $\\rm{M_\\odot}$ and weak thermal pulses follow until carbon ignites in the center when the white dwarf reaches 1.37 $\\rm{M_\\odot}$. Although the supernova production rate through this channel is not well known, and this channel can not be the only one as its progenitor life time is rather short ($\\sim 10^7 - 10^8 $ yr), our results indicate that helium star plus white dwarf systems form a reliable route for producing Type~Ia supernovae. ",
        "introduction": "Type Ia supernovae (SNe~Ia) are of particular importance in astrophysics: They are the  major source for iron group elements in the universe and are thus an essential contributor to the chemical evolution of galaxies (e.g. Renzini~\\cite{Renzini}). And their light curve properties allow it to measure their distances with an excellent accuracy even out to redshifts beyond $z=1$, which makes them a powerful tool to determine the cosmological parameters (e.g. Hamuy et al.~\\cite{Hamuy}; Branch~\\cite{Branch}; Leibundgut~\\cite{Leibundgut}). In particular, the recent suggestion of a non-zero cosmological constant is in part based on SNe~Ia data (Perlmutter et al.~\\cite{Perlmutter}; Riess et al.~\\cite{Riess}). An understanding of the progenitors of these supernovae is clearly required as a basis for these fundamental astrophysical phenomena.  However, even though there seems no doubt that SNe~Ia are produced by the thermonuclear explosion of a white dwarf, it is currently unclear in which kinds of binary systems such an event can occur (Livio~\\cite{Livio}). Among the scenarios which have been put forward as possibilities, the so called single degenerate scenario is currently favored, where a CO white dwarf accretes mass from a non-degenerate companion and thereby grows up to the Chandrasekhar mass (e.g. Hillebrandt \\& Niemeyer~\\cite{Hillebrandt}; Livio~\\cite{Livio}). However, hydrogen as well as helium accretion rates which allow an increase of the CO white dwarf mass due to shell burning are limited to narrow ranges (e.g. Nomoto~\\cite{Nomoto}; Fujimoto~\\cite{Fujimoto}; Iben \\& Tutukov~\\cite{Iben89}). While hydrogen and helium shell sources are prone to degeneracy effects and related thermonuclear instabilities (e.g., nova explosions) at low accretion rates, the helium shell source in accreting white dwarf models has been found to be thermally unstable even for cases where the electron degeneracy is negligible (e.g. Cassisi et al.~\\cite{Cassisi}; Langer et al.~\\cite{Langer02}). This affects in particular the potentially most frequent SN~Ia progenitor systems, where a hydrogen rich star (main sequence star or red giant) is considered as white dwarf companion (e.g. Li \\& van den Heuvel~\\cite{Li}; Hachisu et al.~\\cite{Hachisu}; Langer et al.~\\cite{Langer00}). The hydrogen accretion rates in those systems which may allow steady hydrogen shell burning are about a few $10^{-7}$ \\msyr. With these rates, the subsequent helium shell burning is usually found to be unstable (Iben \\& Tutukov~\\cite{Iben89}; Cassisi et al.~\\cite{Cassisi}; Kato \\& Hachisu~\\cite{Kato99}), and no model sequences which cover the major part of the white dwarf accretion phase could be constructed so far.  In an effort towards overcoming this shortcoming, we consider here the evolution of close helium star plus CO white dwarf systems. In those, the white dwarf develops only a helium shell source, thus avoiding complications involved in double shell source models (e.g. Iben \\& Tutukov~\\cite{Iben89}). Such systems form a predicted binary evolution channel (Iben \\& Tutukov~\\cite{Iben94}), which is confirmed --- even though for lower masses than considered here --- by Maxted et al. (\\cite{Maxted}). Further evidence for the existence of close helium star plus CO white dwarf systems comes from the recent discovery of a helium nova (Ashok \\& Banerjee~\\cite{Ashok}; Kato \\& Hachisu~\\cite{Kato03}).  Simplified binary evolution considerations provided us with an estimate for the optimal initial parameters of our model system. As a result, we embarked on calculating the detailed evolution of a 1.6 \\Msun{} helium star and a 1 \\Msun{} CO white dwarf in a 0.124~d orbit. We introduce our computational method and physical assumptions in Sect.~\\ref{sec:method}. In Sect.~\\ref{sec:results}, the evolution of the considered binary system is presented. We discuss our results in Sect.~\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  The model described above is (to our knowledge) the first self-consistent binary evolution calculation which leads to a Chandrasekhar-mass white dwarf. Its relevance is in part the realistic construction of a supernova progenitor white dwarf model, but even more the fact that it provides for the first time hard evidence of the functioning of a SN~Ia progenitor channel: Our results leave little doubts that some helium star plus white dwarf systems will in fact produce a Type~Ia supernova.  The considered system may have evolved from a wide 8.0 \\Msun{} giant star + 1.0 \\Msun{} white dwarf system though a common envelope phase, which in turn may be the result of two intermediate mass stars in a close orbit. This indicates a system life time of the order of $\\sim 10^7 - 10^8$ yr, which is too short to explain SNe~Ia in elliptical galaxies. Therefore, our results support the idea that at least two SN~Ia progenitor scenarios are realized in nature (e.g., della Valle \\& Livio~\\cite{Valle}). Iben \\& Tutukov (\\cite{Iben94}) estimated the potential SN~Ia production rate through white dwarf + helium giant binary systems to 1.7$\\times10^{-3}$ $\\rm{yr^{-1}}$,  which may constitute a significant fraction of SNe~Ia observed in late type galaxies.  As previously mentioned, SNe~Ia progenitors of the kind considered here will appear as a super-soft X-ray source (SSS). We note that, since a wide range of orbital separations is possible for helium star plus white dwarf systems (Iben \\& Tutukov~\\cite{Iben94}), this may explain SSSs with various orbital periods. For instance, short period systems such as RX J0537.7-7034 (3.5h, Greiner et al.~\\cite{Greiner2}) and 1E0035.4-7230 (4.1h, Schmidtke et al.~\\cite{Schmidtke}) can not be easily explained within the  canonical model which invokes hydrogen rich donor stars (e.g., Rappaport et al.~\\cite{Rappaport}; Kahabka and van den Heuvel~\\cite{Kahabka}), except at low metallicity (Langer et al.~\\cite{Langer00}). Orbital periods of the binary system in the present study in the range $2.65 - 2.97$~h indicate that helium star plus white dwarf systems might be another natural possibility to explain short period SSSs."
    },
    "0402/astro-ph0402415_arXiv.txt": {
        "abstract": "{We compiled a sample of 95 AGNs serendipitously detected in the 3--20~keV band at Galactic latitude $|b|>10^\\circ$ during the RXTE slew survey (XSS, Revnivtsev et al.), and utilize it to study the statistical properties of the local population of AGNs, including the X-ray luminosity function and absorption distribution. We find that among low X-ray luminosity ($\\lx< 10^{43.5}$~erg~s$^{-1}$) AGNs, the ratio of absorbed (characterized by intrinsic absorption in the range $10^{22}$~cm$^{-2}<\\nh<10^{24}$~cm$^{-2}$) and unabsorbed ($\\nh<10^{22}$~cm$^{-2}$) objects is 2:1, while this ratio drops to less than 1:5 for higher luminosity AGNs. The summed X-ray output of AGNs with $\\lx>10^{41}$~erg~s$^{-1}$ estimated here is smaller than the earlier estimated total X-ray volume emissivity in the local Universe, suggesting that a comparable X-ray flux may be produced together by lower luminosity AGNs, non-active galaxies and clusters of galaxies. Finally, we present a sample of 35 AGN candidates, composed of unidentified XSS sources. ",
        "introduction": "\\label{intro}  We have recently \\citep[][ hereafter Paper 1]{revetal04} taken advantage of the excellent calibration, moderate field of view (1~deg radius) and high effective area ($\\sim 6000$~sq. cm) of the PCA spectrometer on board the RXTE observatory to perform  an all-sky survey in the 3--20~keV band from the data accumulated during satellite slews in 1996--2002 -- the RXTE slew survey (XSS). For $90$\\% of the sky at $|b|>10^\\circ$, a flux limit for source detection of $2.5\\times 10^{-11}$ erg s$^{-1}$ cm$^{-2}$ (3--20~keV) or lower was achieved, while a combined area of $7\\times 10^3$~sq.~deg was sampled to record flux levels (for such very large area surveys) below $10^{-11}$ erg s$^{-1}$ cm$^{-2}$.  In Paper~1, a catalog comprising 294 X-ray sources detected at $|b|>10^\\circ$ was presented. 236 of these sources were identified with a single known astronomical object. Of particular interest are 100 identified active galactic nuclei (AGNs) and 35 unidentified sources. The hard spectra of the latter suggest that many of them will probably also prove to be AGNs when follow-up observations are performed. Most of the detected AGNs belong to the local population ($z<0.1$). In addition, the hard X-ray band of the XSS (3--20~keV) as compared to most previous X-ray surveys, performed at photon energies below 10~keV, has made possible the detection of a substantial number of X-ray absorbed AGNs (mostly Seyfert~2 galaxies). These properties make the XSS sample of AGNs a valuable one for the study of the local population of AGNs.  In this paper, we carry out a thorough statistical analysis of the above sample to investigate several key properties of the local population of AGNs, in particular their distribution in intrinsic absorption column density ($\\nh$) and X-ray luminosity function. Knowledge of these characteristics provides important constraints for AGN unification models and synthesis of the cosmic X-ray background, and is further needed to understand the details of the accretion-driven growth of supermassive black holes in the nuclei of galaxies.  In the course of the paper, we compare our results with previously published ones. These include the X-ray luminosity function of local AGNs derived from the HEAO-1/A2 all-sky survey \\citep{picetal82}, the $\\nh$ distribution of optically selected Seyfert~2 galaxies \\citep{risetal99} and the evolving with redshift properties of AGNs inferred largely from medium-sensitivity and deep X-ray surveys \\citep{lafetal02,uedetal03,steetal03}. Finally, we assess the contribution of AGNs with luminosities above $\\sim 10^{41}$~erg~s$^{-1}$ to the total X-ray volume emissivity in the local Universe, as estimated by \\citet{miyetal94}. ",
        "conclusions": ""
    },
    "0402/astro-ph0402623_arXiv.txt": {
        "abstract": "{ Using the sensitive {\\it XMM-Newton} observatory, we have observed the giant \\hii\\ region \\n\\ in the LMC for $\\sim$30~ks. We have detected several large areas of soft diffuse X-ray emission along with 37 point sources. One of the most interesting results is the possible association of a faint X-ray source with BSDL 188, a small extended object of uncertain nature.  The OB associations in the field-of-view (LH9, LH10 and LH13) are all detected with \\x, but they appear very different from one another. The diffuse soft X-ray emission associated with LH9 peaks near HD 32228, a dense cluster of massive stars. The combined emission of all individual massive stars of LH9 and of the superbubble they have created is not sufficient to explain the high level of emission observed: hidden SNRs, colliding-wind binaries and the numerous pre-main sequence stars of the cluster are most likely the cause of this discrepancy. The superbubble may also be leaking some hot gas in the ISM since faint, soft emission can be observed to the south of the cluster. The X-ray emission from LH10 consists of three pointlike sources and a soft extended emission of low intensity. The two brightest point sources are clearly associated with the fastest expanding bubbles blown by hot stars in the SW part of the cluster. The total X-ray emission from LH10 is rather soft, although it presents a higher temperature than the other soft emissions of the field. The discrepancy between the combined emission of the stars and the observed luminosity is here less severe than for LH9 and could be explained in terms of hot gas filling the wind-blown bubbles. On the other hand, the case of LH13 is different: it does not harbour any extended emission and its X-ray emission could most probably be explained by the Sk $-66^{\\circ}$41 cluster alone.  Finally, our \\x\\ observation included simultaneous observations with the OM camera that provide us with unique UV photometry of more than 6000 sources and enable the discovery of the UV emission from the SNR N11L. ",
        "introduction": "Massive stars are known to deeply influence the structure and dynamics of their environment. Their fast stellar winds combined with their huge mass-loss rates and their powerful explosion as supernovae (SNe) can shape the interstellar medium (ISM), creating various structures from small wind-blown bubbles around single stars to large superbubbles around OB associations. This collective action can best be understood by studying nearby \\hii\\ regions, with embedded OB associations containing hundreds of massive stars. In this context, we have chosen to study the \\n\\ complex (Henize \\cite{hen}) in the Large Magellanic Cloud (LMC).  \\n\\ is the second largest \\hii\\ region in the LMC after 30 Doradus, and it may constitute a more evolved version of this latter nebula (Walborn \\& Parker \\cite{wal92}). It harbors several associations of massive stars: LH9, LH10, LH13 and LH14 (Lucke \\& Hodge \\cite{luc}). Its structure is complex and reflects the interactions between the stars and their environment. The central cluster, LH9, is surrounded by a filamentary shell of $\\sim$120~pc in diameter. The combined action of stellar winds and supernova explosions from the members of the cluster has carved a hollow cavity in the surrounding ISM, thereby creating the shell which is also called a superbubble. The hot shocked winds/ejecta that fill this cavity emit X-rays (Mac Low et al. \\cite{mac}) and provide the pressure to drive the superbubble shell expansion. Only supernovae outside the superbubble can produce distinct supernova remnants (SNRs) such as N11L at the western edge of the N11 complex (Williams et al. \\cite{wil}).  \\begin{figure*} \\begin{center} \\end{center} \\caption{Three colour image of the \\n\\ region as seen by the combined EPIC cameras. The red, green and blue colours correspond respectively to the $0.4-1.0$ keV, $1.0-2.0$ keV and $2.0-10.0$ keV energy bands. Note that this image is based on a pn event list that was filtered using \\#XMMEA\\_EP and $0\\le$ pattern $\\le12$. This event list was only used for aesthetic reasons in the aim of creating Figs. \\ref{color} \\& \\ref{totfield}, not for any scientific analysis. \\label{color}} \\end{figure*}  The action of LH9 on its surroundings has also triggered a burst of star formation at the periphery (Rosado et al. \\cite{ros}), leading to the birth of the three other OB associations. Situated to the north of LH9, LH10 is still embedded in its natal cloud but its most massive components have already begun to blow bubbles around them (Naz\\'e et al. \\cite{naz}). The stellar population of LH9 and LH10 has been studied by Parker et al. (\\cite{par}, hereafter PGMW) and Walborn et al. (\\cite{wal99}). To the east of LH9 lies LH13, which appears to contain two tight clusters, Sk $-66^{\\circ}$41 and HNT (Heydari-Malayeri et al. \\cite{hey00}). However, the ages and radial velocities indicate that the $\\le$ 5 Myr old cluster Sk~$-66^{\\circ}$41 is associated with the surrounding \\hii\\ region and the  $\\sim$100 Myr old cluster HNT is unrelated and most probably just a line-of-sight object (Heydari-Malayeri et al. \\cite{hey00}). Situated at the northeast outskirts of \\n, LH14 is the least studied of the four OB  associations; the existence of a few massive stars in LH14 has been  illustrated in a photometric study by Heydari-Malayeri et al. (\\cite{hey87}).  \\begin{figure*} \\begin{center} \\end{center} \\caption{Combined EPIC image of \\n\\ in the energy range $0.4-10$ keV. The detected sources are labelled. The image has been binned by a factor of 50, to obtain a pixel size of 2.5\\arcsec. \\label{totfield}} \\end{figure*}  In summary, the \\n\\ complex harbours a variety of phenomena associated with massive stars (e.g. bubbles and  superbubbles), and even contains a SNR. All these objects should be associated with some hot gas. \\n\\ thus provides an interesting target for deep X-ray observations. A {\\it ROSAT} investigation only permitted to detect some X-ray emission (Mac Low et al. \\cite{mac}, Dunne et al. \\cite{dun}), but the unprecedented sensitivity of {\\it XMM-Newton} enables us to study this region in more details.  In the following sections, we will first describe the observations and the X-ray sources detected in the field. Next, we will focus on the main components of \\n: the SNR, the superbubble, LH10 and LH13. Finally, we will conclude in Section~6. ",
        "conclusions": "We report in this paper the observations by the {\\it XMM-Newton} satellite of the giant \\hii\\ region \\n\\ situated in the LMC. In this field, we have detected large areas of soft diffuse X-ray emission and 37 point sources, one of which is apparently associated with a small and poorly known extended object, BSDL 188.  \\n\\ harbours a wealth of phenomena associated with massive stars: it contains a SNR, N11L, and four OB associations at different stages of evolution and interaction with their surroundings (LH9, LH10, LH13 and LH14). All are detected in the {\\it XMM-Newton} data, except for LH14, which is unfortunately outside of the \\x\\ FOV. The stars from LH9, the largest cluster of \\n, have blown a superbubble which is detected in X-rays as a large soft X-ray emission. It peaks near HD 32228, a dense cluster containing a WR star and several OB stars, and it is rather well confined within the \\ha\\ filaments delineating the superbubble. The combined emission of all individual massive stars of LH9 can not explain the high level of emission observed, nor can the contribution from the hot shocked gas of the superbubble as predicted by the Weaver et al. (\\cite{wea}) model. Hidden SNRs, colliding-wind binaries and the often neglected pre-main sequence stars most probably provide additional X-ray emission. It is probably a combination of all these effects that will resolve the discrepancy, since an estimation of the contribution of the latter alone has been shown insufficient to explain the whole X-ray emission. Moreover, we do not exclude that the superbubble is leaking some hot gas since faint, soft emission is detected to the south of the cluster.  The action of LH9 on its surroundings has probably triggered the formation of a second cluster, LH10. The X-ray emission from this younger cluster consists of three rather pointlike sources, in addition to a soft extended emission of reduced intensity. The two brightest sources are not centered on the most powerful stars of the cluster, i.e. the young O3 stars, but seem associated with the expanding bubbles blown by stars in the SW part of LH10 (Naz\\'e et al. \\cite{naz}). The total X-ray emission of LH10 is four times larger than the total emission from its stellar components. Since the cluster is still very young, the excess emission can probably not be attributed to hidden SNRs or active T Tauri stars, but it probably originates from the hot gas filling the wind-blown bubbles. To completely understand the X-ray emission from these two clusters, disentangling the individual contributions (extended sources vs. a simple accumulation of non-resolved point sources) is of the utmost importance. However, this is not feasible in a reasonable amount of time with the current X-ray observatories, but \\n\\ will be a perfect target for the next generation of X-ray satellites.  In contrast, the X-ray emission from LH13 does not show any extended emission and could be well explained by the stars of the Sk $-66^{\\circ}$41 cluster alone. To the west of the field, the SNR N11L shows hot gas outside the optical bubble, which is associated with an additional radio emission. The spectra of the SNR and of the extended X-ray plume to its north are similar, though with a larger absorbing column for the latter. For the first time, the SNR was also detected in the ultraviolet, and it presents at these wavelengths a morphology very similar to the optical one.  Note that the high excitation blob N11A is not detected in our observation, but this is not completely surprising since at least the theoretically expected X-ray flux from the stellar population embedded in this nebula (Heydari-Malayeri et al. \\cite{hey01}) is very small.  During our X-ray observation, the OM camera onboard {\\it XMM-Newton} has provided unique UV photometry of more than 6000 sources. This photometry is available to the scientific community through CDS."
    },
    "0402/astro-ph0402309_arXiv.txt": {
        "abstract": "We investigate the influence of blending on the Cepheid distance scale using two Local Group galaxies, M31 and M33. Blending leads to systematically low distances to galaxies observed with the HST, and therefore to systematically high estimates of $H_0$. High-resolution HST images are compared to our ground-based data, obtained as part of the DIRECT project, for a sample of 22 Cepheids in M31 and 102 Cepheids in M33. For a sample of 22 Cepheids in M31, the average (median) flux contribution from luminous companions not resolved on the ground-based images in the $V$-band, $S_V$, is about 19\\% (12\\%) of the flux of the Cepheid. For 102 Cepheids in M33 the average (median) values of $S_V$, $S_I$, $S_B$ are 23\\% (13\\%), 28\\% (20\\%), 28\\% (15\\%). For 64 Cepheids in M33 with periods in excess of 10 days the average (median) $S_V$, $S_I$, $S_B$ are 16\\% (7\\%), 23\\% (12\\%), 20\\% (10\\%). ",
        "introduction": "As the number of extragalactic Cepheids discovered with {\\em HST} continues to increase and the value of $H_0$ is sought from distances based on these variables, it becomes even more important to understand various possible systematic errors which could affect the extragalactic distance scale. Currently, the most important systematic is a bias in the distance to the Large Magellanic Cloud, which provides the zero-point calibration for the Cepheid distance scale (e.g. Udalski 2000; Fitzpatrick et al.~2003). Another possible systematic, the metallicity dependence of the Cepheid Period-Luminosity (PL) relation, is also very much an open issue, with empirical determinations ranging from 0 to $-0.4$ mag dex$^{-1}$ (e.g.\\ Sasselov et al.~1997; Udalski et al.\\ 2001).  We define {\\em blending} as the close projected association of a Cepheid with one or more intrinsically luminous stars, which cannot be detected within the observed point-spread function (PSF) by photometric analysis. Blending is thus a phenomenon different from {\\em crowding} or {\\em confusion noise}; the latter occurs in stellar fields with a crowded and complex background due to the random superposition of stars with different luminosities.  We investigate the effects of stellar blending on the Cepheid distance scale by studying two Local Group spiral galaxies, M31 and M33.  We identify some of the Cepheids, discovered by the DIRECT project (Stanek et al.\\ 1999, Mochejska et al.\\ 1999) on archival {\\em HST}-WFPC2 images and compare them to our ground-based data to estimate the impact of blending on our photometry, taking advantage of their superior resolution. ",
        "conclusions": ""
    },
    "0402/astro-ph0402352_arXiv.txt": {
        "abstract": "{We report detection of a Young Stellar Object with an evidence for an outflow in the form of knots  in the molecular hydrogen emission line (2.121$\\mu$m) towards the massive star forming region  IRAS 06061+2151. Near-infrared  images reveal IRAS 06061+2151 to be a cluster of at least five sources, four of which seem to be early B type young stellar objects,   in a region of  12 arcsecs surrounded by a nebulosity. The presence of  the knots that are probably similar to the HH objects in the optical  wavelengths, suggests emerging jets from one of the  cluster members. These jets appear to excite     a pair of knot-like objects (Knot-NW and Knot-SE) and extend over a projected size of 0.5pc. The driving source for the jets is traced back to a member of the cluster  whose position in the H-Ks/J-H color-color diagram indicates that it is a  Class I type pre-mainsequence star.   We also obtained K band spectra of the brightest source in the cluster and of the nearby nebular matter.  The spectra show molecular hydrogen emission lines but do not show  Br $\\gamma$ line (2.167$\\mu$m).  These spectra suggest that  the excitation of the molecular hydrogen lines is probably  due to a mild shock. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0402/astro-ph0402164_arXiv.txt": {
        "abstract": "{ We consider internal shocks as the main dissipation mechanism responsible for the emission in blazars and show that it can satisfactorily account for the properties of all blazars. In particular, we extend previous work (Spada et al. 2001) on powerful objects, to intermediate (BL Lac) and low power sources (Mkn 421), in order to reproduce the whole of the blazar sequence.  The model self-consistently treats the dynamics, spectral emission and its variability.  The key parameters driving the phenomenological sequence are the jet power and the properties of the broad line region, which regulate the cooling efficiency of the emitting particles and thus the shape of the spectral energy distribution. By assuming that the remaining parameters are similar for all objects it has been possible to reproduce the full range of the observed spectral ``states\". A more detailed comparison of the variability properties shows (for Mkn 421) a good agreement in the X--ray band, while in the optical the simulated flux appears to be too variable.  For BL Lac lags ($\\sim$ 10 days) are predicted between the $\\gamma$--rays and the infrared emission. ",
        "introduction": "The discovery that blazars are strong $\\gamma$--ray emitters together with the results of multiwavelength campaigns have allowed to deepen our knowledge on these objects. Their Spectral Energy Distribution (SED) is characterized by two broad peaks (Fossati et al. 1998) strongly variable on different timescales (Wagner \\& Witzel 1995; Ulrich, Maraschi \\& Urry 1997).  Two main radiation processes dominate the emission, namely synchrotron at low frequencies and -- plausibly -- inverse Compton at high energies (see e.g. Sikora 1994 for a review).  The relative luminosity in the two peaks and their peak frequency appear to be functions of the total power (Fossati et al 1998, Ghisellini et al. 1998), resulting into a sequence for the whole of the blazar population properties, ranging from powerful, low frequency peak, through intermediate, to low power, high frequency peak (blue) blazars (see however Padovani et al. 2003).  The blazar emission is variable on energy dependent timescales, typically of weeks--months in the radio and of the order of a day in the $\\gamma$--ray band.  Several studies, mainly based on the modeling of the SED, consistently derived the physical parameters of the emitting region. However key issues in the understanding of relativistic jets remain open, most notably the jet energetics and the particle acceleration process(es). In order to explore these issues and their relationship we have (quantitatively) considered a scenario in which the plasma conditions and their variability are not treated as free parameters, but follow from the jet dynamics, thus relating the observed emission properties with the energy transport along the jet.  Such a scenario assumes that internal shocks are responsible for the dissipation within jets (Rees 1978 and then mostly explored for Gamma--Ray Bursts, Rees \\& M\\'esz\\'aros 1994).  The key assumption of the model is that the energy is channeled into jets in an intermittent way by the central engine, though such a time dependent process cannot be inferred from first principles.  Different parts of the jet (`shells') moving at different speeds can collide giving rise to shocks and dissipation as non--thermal radiation. The mechanism has a limited efficiency (unless the contrast in Lorentz factors between different shells is extremely large, see Beloborodov 2000; Guetta, Spada \\& Waxman 2001, but also Ghisellini 2002), which has to be indeed the case for blazars since most of the energy propagates up to the extended radio lobes.  Beside the low efficiency, the internal shock scenario can naturally account for other blazar properties. It predicts that jets become radiative at $\\gsim 10^{16}- 10^{17}$ cm from the central engine, implying a minimum distance from the accretion disk and a minimum dimension for the $\\gamma$--ray source, as required by observations (in order to avoid copious pair production; see Ghisellini \\& Madau 1996).  Furthermore, successive collisions taking place at larger distances have reduced efficiency -- since the Lorentz factor contrast of the colliding shells decreases -- explaining why the jet luminosity decreases with distance. Finally, the internal shock scenario appears a promising non steady state model to account for the observed large amplitude variability.  A detail study of the predictions of this model via numerical simulations has been carried out by Spada et al. (2001; hereafter S01) for powerful blazars (specifically 3C 279).  As the SED of this object does not represent the whole of the blazar family mentioned above, we focus here on two well studied sources representative of extremely blue and intermediate blazars, namely Mkn 421 and BL Lac itself. Rather than reproducing in detail particular spectra of specific objects, the aim of this work is to determine whether the proposed scenario, including the dynamics and emission properties, i) can account for the different SED along the blazar sequence, ii) under which hypothesis this can occur and iii) whether the internal shock model can reproduce the observed variability behavior.  The outline of the paper is the following.  In \\S 2 we describe the hypothesis on the wind dynamics and on the radiative properties of the shocked plasma.  The results -- specifically referring to the sources Mkn421 and BL Lac -- are presented in \\S 3.  In \\S 4 we draw our conclusions. ",
        "conclusions": "In this work we have considered internal shocks as the dissipation mechanism responsible for the emission in blazars.  As mentioned in the Introduction, this scenario is currently the most accredited to explain the prompt gamma-ray emission in Gamma-Ray Bursts, despite of an efficiency problem (Lazzati, Ghisellini \\& Celotti 1999).  In blazars instead the relatively low efficiency is in fact required as, at least in powerful objects, most of the jet power is carried to the large scales.  In S01 it was found that the internal shock model was successful in reproducing the observed SED and variability properties of a powerful blazar, namely 3C~279. However, the SED of this object does not represent the whole of the blazar class, which covers a wide range of spectral characteristics (luminosity, frequency of the peaks of the emission and their relative intensity). Interestingly these parameters appear to be correlated and the whole class can be seen as a sequence: the frequency and the ratio of the low vs high energy peak intensity increase with decreasing luminosity (Fossati et al. 1998).  We show here that the internal shock model can satisfactorily account also for the properties of the lower power blazars. The key parameters driving the phenomenological sequence are the jet power (proportional to the radiated one) and the intensity of the broad lines.  These parameters in turn regulate the SED shape, as they control the cooling efficiency of the emitting particles. The global radiative efficiency appears instead to be similar for all of the sources examined.  The internal shock scenario determines a characteristic time interval for the injection of relativistic electrons of the order of the dynamical timescale, when the intensity of the spectrum from each collision is maximized.  Two regimes are relevant: {\\it fast} and {\\it slow} cooling, corresponding to whether electrons of energy $\\gamma_{\\rm b}$ can or cannot radiatively cool in the dynamical time. In highly powerful blazars the fast cooling regime dominates in the inner regions (within the BLR) where also most of the power is dissipated.  Consequently the peak frequencies are produced by electrons of energy $\\gamma_{\\rm b} m_{\\rm e} c^2$.  In the weakest blazars instead the slow cooling regime prevails over the whole jet. Consequently only the highest energy electrons can cool over such timescale: the peak frequencies thus shift to high values.  Between these two extremes there are intermediate sources, like BL Lac, with broad lines of intermediate intensity, produced at a distance within which a few shell-shell collisions can occasionally take place. Observationally this corresponds to SED with moderate Compton to synchrotron luminosity ratio, as the scattered seed photons are only the synchrotron ones (in collisions outside the BLR).  The rare collisions within the BLR give rise to dramatic changes in the SED characterized by a large increase of the Compton component (e.g. BL Lac itself, see Fig.~2).  The model considered here, which considers self-consistently the dynamics and spectral emission, predicts also the time dependent spectral properties.  In general, the selected parameters allow to reproduce the full range of spectral 'states' observed in both BL Lac and Mkn 421. A more detailed comparison performed for Mkn 421, and based on the r.m.s. and variability timescales, shows good agreement for the X--ray variability properties, while the simulated optical variations appear to be too large with respect to the considered observed light curve.  An analysis of the predicted cross-correlated variability between the $\\gamma$--ray and other bands reveals that only for BL Lac lags ($\\sim$ 10 days) are expected between the $\\gamma$--rays and the infrared emission.  We conclude that the internal shock scenario can account for the the main properties of blazars and the `blazar sequence'.  While it has been previously pointed out that the key quantity in reproducing the different characteristics of blazars is the ratio between the jet and the disc (i.e. broad lines) luminosities, the internal shock scenario discussed here provides more physical insights, namely i) directly connects the radiated jet luminosity with the jet effective power and ii) accounts for the `preferred' distance where most of the luminosity is dissipated.  While the latter is similar for high and low power blazars, the BLR is instead located at different distances in the different sub--classes of objects (as determined by the ionizing luminosity).  From a more theoretical point of view, the model does not (yet) address the role of a seed magnetic field amplified by the shell-shell collisions. This would probably lead to a faster synchrotron cooling on the large scales and a flatter dependence of the $B$ field from $R$ (Fig.~3). We intend to further explore this issue.  Furthermore, only internal shocks have been considered in this scenario. It is conceivable that at large jet scales some entrainment may occur causing an interaction of the jet with the external medium, possibly leading to external shocks. As a result the radiative efficiency and large scales synchrotron and inverse Compton emission could be enhanced. Clearly, the inverse Compton emission could be also enhanced if a significant photons field is present externally to the jet even on large scales (such as microwave background and/or beamed nuclear radiation and/or dust emission, e.g. Celotti, Ghisellini \\& Chiaberge 2001; Sikora et al. 2002).  Observationally the improved detection sensitivity of the planned TeV Cherenkov telescopes -- such as VERITAS, HESS, MAGIC -- will allow in the near future to measure emission from BL Lacs in low luminosity states, and thus to estimate their flare activity duty cycles.  Since in the internal shock scenario the TeV emission is largely produced by the few powerful internal collisions, it will be then possible to further constrain model parameters such as the range of bulk Lorentz factors and the initial separation of the shells. Analogous results will be likely achieved for the more powerful blazars in the GeV band thanks to the AGILE and GLAST satellites."
    },
    "0402/astro-ph0402487_arXiv.txt": {
        "abstract": " ",
        "introduction": "Over the past ten years the interest in the nature and origin of the highest-energy cosmic rays, those with energy $\\geq 10^{19}$ eV, has grown enormously. Of particular interest are cosmic rays with energy $\\geq10^{20}$ eV. At these energies the cosmic ray particles, be they protons, nuclei, or photons, interact strongly with the cosmic microwave background and should be severely attenuated, except for those whose sources are in our cosmological neighborhood ($\\leq$ 100 Mpc). Also protons of these energies may not be significantly deflected so that some of the cosmic rays may point back to their source. A recent authoritative review by Nagano and Watson \\cite{nw00} presents the experimental and theoretical background. For all but the most recent references I refer the reader this article.  I dedicate this article to the memory of John Linsley (b. March 12, 1925; d. September 15, 2002). He discovered the first cosmic ray with an energy of $10^{20}$ eV \\cite{Linsley1} 42 years ago. John Linsley has contributed so much to our understanding of cosmic ray showers. Many of his contributions are tucked away in old proceedings of cosmic ray conferences. His original development of the concept of elongation rate is to be found in the Proceedings of the 15th ICRC held in Plovdev, Bulgaria \\cite{Linsley2}. (The elongation rate is referred to in section 3 and section 5.3 of this paper.) Up until his death he was active in all aspects of cosmic ray physics. He was the founder of the idea for a satellite-borne fluorescence telescope, which is being realized in the EUSO detector to be placed on the Space Station.  He considered me an upstart when I began with colleagues to argue for a really large surface detector' one which ultimately became the Auger Observatory. However with the passage of time I believe I gained his respect. I only wish that he could witness the progress that is going to be made. I cannot avoid the fantasy that he is now in a position to know all the answers!   \\begin{figure} \\centerline{\\hbox{\\psfig{figure=taup_fig1.ps,height=3.5in,angle=90}}} \\caption{Number of theoretical and speculative papers on the subject of the highest-energy cosmic rays.} \\end{figure}  \\begin{figure} \\centerline{\\hbox{\\psfig{figure=taup_fig2.ps,height=8.5in,angle=0}}} \\caption{Some titles of paper concerning the highest-energy cosmic rays.} \\end{figure}     Much more data will be required to understand the mystery behind the existence of cosmic rays with such extraordinary energies. An article by Michael Hillas \\cite{Hillas1}, published nearly twenty years ago, presented the basic requirements for the acceleration of particles to energies $\\geq10^{19}$ eV by astrophysical objects. The requirements are not easily met, which has stimulated the production of a large number of creative papers.   In Figure 1 I plot the number of theoretical papers, mostly speculative, written on the subject of the highest-energy cosmic rays as a function of time, as found on the Los Alamos server as astro-ph papers. Over the last three years the average has been one paper per week. In Figure 2 I list a random sample of the titles. The authors of these papers deserve a strong response from the experimental community.   \\begin{figure} \\centerline{\\hbox{\\psfig{figure=taup_fig3.ps,height=3.5in,angle=90}}} \\caption{Panorama of the interactions of possible cosmic primaries with the CMB. Curves marked by ``p+$\\gamma_{CMB}$ $\\rightarrow$ e$^{+}$e$^{-}$+p'' and ``Fe+$\\gamma_{CMB}$ $\\rightarrow$ e$^{+}$e$^{-}$+p'' are energy loss lengths (the distance for which the proton or Fe nucleus loses 1/e of its energy due to pair production). The curve marked by ``p+$\\gamma_{CMB}$ $\\rightarrow$ $\\pi^{+}$n or $\\pi^{\\circ}$p'' is the mean free path for photo-pion production of a proton on the CMB. The curve marked ``Fe+$\\gamma_{CMB}$ $\\rightarrow$ nucleus + n or 2n'' is the mean free path for a photo-nuclear reaction where one or two nucleons are chipped off the nucleus. The curve marked ``$\\gamma$ +$\\gamma_{CMB}$ $\\rightarrow$ e$^{+}$e$^{-}$'' is the mean free path for the interaction of a high-energy photon with the CMB. Added for reference is the mean decay length for a neutron indicated by ``n $\\rightarrow$ pe$\\nu$''.} \\end{figure} ",
        "conclusions": "I have attempted in this survey to explain without excessive complication the important aspects for the measurement of the properties of cosmic rays at the very highest energies ($\\geq$ $10^{19}$ eV). There remain major uncertainties in all areas, spectrum, anisotropy, and composition. The AGASA experiment is now complete and one awaits the final catalog of events. It should be possible for the AGASA group to extend their published measurements to larger zenith angles. I eagerly await the additional data and improved analysis from the HiRes group.  I am also looking forward to see the new data which will come from the Pierre Auger Observatory, now under construction in the southern hemisphere. It is a hybrid detector which consists of both surface detectors and fluorescence telescopes. About 10$\\%$ of showers will be observed simultaneously by both techniques. I will not list the advantages of the hybrid detector here; the reader can surely appreciate them. The curious reader is encouraged to visit the many web sites devoted to the Auger Observatory which can all be reached through {\\bf www.auger.org}. As I write these conclusions (February  2004) the Auger Observatory is operating with 220 surface detectors and 6 30$^{\\circ}$ x 30$^{\\circ}$ fluorescence telescopes. By the end of 2004 there will be 700-800 surface detectors and 12-14 fluorescence telescopes. By the end of 2005 it should be complete with 1600 surface detectors covering 3000 km$^2$ and 24 fluorescence telescopes overlooking the array.  The reader can surely appreciate the enormous gain in statistics and the detailed information that will be obtained for each event. However, from long experience in physics I succumb to caution and will refrain from making too many claims of what the Observatory will accomplish. Certainly we will learn a great deal. Pierre Auger made his great discovery 66 years ago. Nature rewarded John Linsley with a 10$^{20}$ eV shower 42 years ago. I hope that the many questions that these great discoveries have raised will be answered in a time much shorter than 42 or 66 years."
    },
    "0402/astro-ph0402214_arXiv.txt": {
        "abstract": "Observations of the PSR\\,B1259$-$63/SS2883 binary system using the CANGAROO-II Cherenkov telescope are reported. This nearby binary consists of a 48\\,msec radio pulsar in a highly eccentric orbit around a Be star, and offers a unique laboratory to investigate the interactions between the outflows of the pulsar and Be star at various distances. It has been  pointed out that the relativistic pulsar wind and the dense mass outflow of the Be star may result in the emission of gamma rays up to TeV energies. We have observed the binary in 2000 and 2001, $\\sim$47 and $\\sim$157 days after the October 2000 periastron. Upper limits at the 0.13--0.54 Crab level are obtained. A new model calculation for high-energy gamma-ray emission from the Be star outflow is introduced and the estimated gamma-ray flux considering Bremsstrahlung, inverse Compton scattering, and the decay of neutral pions produced in proton-proton interactions, is found to be comparable  to the upper limits of these observations. Comparing our results with these model calculations, the mass-outflow parameters of the Be star are constrained. ",
        "introduction": "\\label{sect:intro} PSR B1259$-$63 (($\\alpha,\\delta$)(J2000) $=$ (13$^h$02$^m$47$^s$.68, $-$63\\arcdeg50\\arcmin08\\arcsec.6)) is a 48\\,msec radio pulsar discovered in a 1500\\,MHz radio survey of the southern Galactic plane\\,\\citep{johnston92a} which was subsequently found to be in a highly eccentric orbit with a 10th magnitude main-sequence star, SS\\,2883\\,\\citep{johnston92b,johnston94}. With an orbital eccentricity of 0.87, the separation of the stars varies in the range 0.97$\\sim$14.0$\\times {\\rm 10}^{13}$cm during the orbital period of 1236.72 days. The periastron epoch is MJD~48124.35\\,\\citep{wex98}. SS\\,2883 is of spectral type B2e\\,\\citep{johnston94}, with a mass $M_*$ of $\\sim$10\\,$M_{\\sun}$ and a radius $R_*$ of $\\sim$6\\,$R_{\\sun}$. The luminosity and radius of the B2e star correspond to an effective temperature $T_{\\rm eff}$ of $\\sim$27,000~K at the star surface\\,\\citep{tavani97}. Its characteristic emission disc extends to at least 20\\,$R_*$, similar to the distance between the pulsar and the Be star at periastron. Here we assume a distance of 1.5\\,kpc to the binary system, which has been estimated from optical photometric observations of SS\\,2883\\,\\citep{johnston94}.  The periastron passages have been closely observed at radio frequencies \\,\\citep{johnston99,connors02}. No pulsed emission was detected for about five weeks centered on periastron and the pulsed emission was depolarized for $\\sim$200 days also centered on periastron. Timing measurements have shown that the disc of the Be star is likely to be inclined with respect to the orbital plane \\,\\citep{melatos95,wex98}, which has been suggested in \\citet{kaspi95,tavani97}. In \\citet{connors02}, the unpulsed light curves are discussed with an assumption of two short-time crossings of the pulsar and the disc, before [($\\tau - $18~d) $\\sim$ ($\\tau - $8~d)] and after [($\\tau +$12~d) $\\sim$ ($\\tau +$22~d)]  periastron ($\\tau$).  A weak X-ray signal was first detected by {\\it ROSAT} which observed the system just after apastron in September 1992\\,\\citep{cominsky94}. Through 1994--1996, unpulsed X-ray emission with a single power-law spectrum was detected at the six different orbital phases observed by  {\\it ASCA} \\,\\citep{hirayama96,hirayama99}. The photon index of the X-ray spectrum is about $-$1.6 in the post-periastron to apastron period, steepening towards periastron where the steepest index of $-$1.96 was observed. The 1--10~keV band luminosity varies by about an order of magnitude, from $\\sim$10$^{34}$ ergs s$^{-1}$ around periastron to $\\sim$10$^{33}$ ergs s$^{-1}$ at apastron. The maximum luminosity was detected at $\\tau -$12~d, with the intensity decreasing at periastron, then increasing again. The column density was low and constant ($\\rm{6}\\times\\rm{10}^{21}\\rm{cm}^{-2}$) at all orbital phases. The periastron passage in January 1994 was monitored by a multi-wavelength campaign including observations in the X-ray and gamma-ray bands with {\\it ROSAT}, {\\it ASCA} and {\\it CGRO} \\citep{grove95}. The power-law spectrum (photon index $\\sim -$2.0) extended to the 200\\,keV energy region of {\\it OSSE}, with no pulsations being detected. No emission in the energy range of 1\\,MeV--3\\,GeV was detected down to the observational limits. {\\it OSSE} failed to detect signals at the apastron passage in 1996, however, its upper limit does not conflict with the extrapolation of the {\\it ASCA} spectrum\\,\\citep{hirayama99}. Several TeV observations of the binary system were performed in 1994 and 1997 using the CANGAROO 3.8-m ground-based Cherenkov telescope, resulting in a marginally significant suggestion of gamma-ray signals \\citep{sako97}.  The multi-wavelength spectrum from the 1994 periastron strongly implies that the hard X-ray emission up to 200\\,keV originates from synchrotron radiation of non-thermal electrons\\,\\citep{tavanikaspi94}. Electrons released in the pulsar wind may be accelerated in a shock wave generated in the region where the relativistic pulsar wind interacts with dense mass flow from the Be star. Adjusting the pressure balance between the flows, \\citet{tavani97} have interpreted the measured hard X-ray spectrum on the basis of accelerated particles in the pulsar-side shock, using an approximated approach to the Klein-Nishina effect for emission. They conclude that the energy loss of electrons due to inverse Compton scattering is dominant, and that the Lorentz factor of accelerated electrons is $\\Gamma_{\\rm e} = {\\rm 10}^6-{\\rm 10}^7$. Accretion onto the neutron star is unlikely to be significant, as there is an absence of X-ray/gamma-ray pulsations, an absence of the day--scale fluctuations in X-rays, a relatively low X-ray luminosity, and negligible absorption as a result of the low column density. The consistency of the X-ray luminosities at the same orbital phases in different years supports the idea that the observed time variability is due to binary modulation \\citep{hirayama99,kaspi97}. Recently, another model for the pulsar-side shock considering the Klein-Nishina effect in the emission and cooling process has been proposed. \\citet{shibazaki02} note that the inverse Compton cooling dominated spectrum is flatter, since the Klein-Nishina effect suppresses the cooling of higher-energy electrons. It is argued that synchrotron cooling, instead of the inverse Compton scattering discussed by \\citet{tavani97}, is the dominant process for the energy loss of electrons in the pulsar wind on the account of the steepening of the X-ray spectral index observed around the 1994 periastron.  The light curves of the radio unpulsed emission around periastron have been recently modeled by the adiabatic expansion of synchrotron bubbles formed in the pulsar and the Be star disc interaction \\citep{connors02}. They assume short-time interactions of the pulsar and the disc, as the pulsar should cross the disc twice in the orbital period. When the pulsar enters the disc, electrons are accelerated in the contact surface of the pulsar wind and the disc material, but after the pulsar leaves the disc, the pulsar-wind bubble remains behind, moves in the disc-flow, and decays through synchrotron losses. The model successfully explains the radio data, however, it does not appear to describe the X-ray data well; for example, the weak unpulsed emissions in the X-ray region was observed after these modeled bubbles should have decayed by adiabatic expansion as moving outwards, and, the constant spectral index in the radio region is inconsistent with the steepening observed in the X-ray spectrum.  Electrons accelerated in a pulsar-side shock to Lorentz factors of $\\Gamma_{\\rm e} ~\\gtrsim \\rm{10}^6$ in the radiative environment of the binary system may produce high energy gamma rays. The suggestion that detectable levels of gamma-ray emission may arise in the shocked pulsar wind via inverse Compton scattering \\citep{kirkball99} provided the initial motivation for the observations described here. Subsequently, inverse Compton emission from the un-shocked region of the pulsar wind has been considered \\citep{ballkirk00,balldodd01}. The integrated contribution from the un-shocked pulsar wind may increase the gamma-ray flux around periastron for some conditions. The maximum level of emission in the TeV energy range is estimated to be $\\sim{\\rm 4} \\times {\\rm 10}^{-5}{\\rm MeV cm}^{-2}{\\rm s}^{-1}$ in the integrated energy flux with a wind Lorentz factor of 10$^7$\\,\\citep{ballkirk00} which may raise the TeV gamma-ray flux above our detector's sensitivity of typically ${\\rm 10}^{-11}-{\\rm 10}^{-12}{\\rm TeV cm}^{-2}{\\rm s}^{-1}$. Further studies have considered the effect of the termination of the wind \\citep{balldodd01}, which, depending on its assumped location, may act to decrease the inverse Compton flux compared to previous predictions. On the other hand, \\citet{shibazaki02} predict a maximum energy flux from the pulsar wind of $\\sim {\\rm 10}^{-13} {\\rm erg \\,s}^{-1}{\\rm cm}^{-2}$ at TeV energies around periastron, which would require a very deep observation to detect.  At the contact surface of the pulsar and Be star flows, ions and electrons in the Be star outflow may be accelerated to high energies via the first-order Fermi mechanism. In the dense outflow from the Be star there is a lot of target material for proton-proton interactions and Bremsstrahlung emission. The Be star also provides target photons for upscattering by the inverse Compton mechanism, in addition to the 2.7\\,K microwave background radiation. The densities of these targets increase as the contact surface gets closer to the Be star, and so does the total energy of accelerated particles. In this paper, a new model calculation for gamma-ray emission from the accelerated particles in the Be star outflow, taking into consideration Bremsstrahlung, the inverse Compton mechanism, and proton-proton interactions, is applied to the binary system and discussed along with our observational results. ",
        "conclusions": "The observational results are compared with some model calculations. A new model of gamma-ray emissivity is introduced, considering the particles accelerated in the Be star outflow.  \\subsection{Models of the two flows} Figure\\,\\ref{fig:binary_system_image} schematically illustrates the assumed configuration of the system: the pulsar and its relativistic pulsar wind, the Be star and its polar and disc-like outflows, and the shock composed of three surfaces: pulsar-side shock, contact surface, and Be-star-side shock. Particles are assumed to be accelerated by the shock at the pressure balance between the flows of the two stars. In the figure, the contact discontinuity between the pulsar wind and the equatorial disc of the Be star is illustrated. The alignment of the Be star disc to the orbital plane, and its effect, will be discussed later in calculating light curves over orbital phase.  For the pulsar wind, we adopt the model of \\citet{kennel84} for the synchrotron nebula around the Crab pulsar.  For the Be star mass-flow, the simple model of \\citet{waters86} is used, which represents radiations from Be stars using the IR, optical, and UV observational results. The parameters are chosen so as to be consistent with the observational results of Be stars in general \\citep{cote87} and of the PSR B1259$-$63/SS2883 binary \\citep{johnston94, johnston96, melatos95}. We fully consider the Klein-Nishina effect in the calculations of the emission processes via electrons. Provided the pulsar wind is driven by the spin down luminosity ($\\dot{E}_{\\rm rot}$) of the pulsar, a fraction ($f_{\\rm pw} =$  0.1) of the wind luminosity is assumed to be enhanced in the equatorial plane. Both kinetic and electromagnetic energies are included in $\\dot{E}_{\\rm rot}$ \\,\\citep{kennel84}. The radial distribution of the wind pressure, $P_{\\rm pw}$, is given by \\begin{equation} \\label{eq:p_pw} P_{\\rm pw}(r) = \\frac{\\dot{E}_{\\rm rot}}{f_{\\rm pw}4\\pi r^2 c}~~, \\end{equation} where $r$ is the distance from the pulsar and $c$ is the speed of light.  For the mass-flow of the Be star, we consider a high-density, slow, equatorially orbiting disc-like flow \\citep{waters86}, and a low-density, fast, polar component (stellar wind) \\,\\citep{waters88,dougherty94} as well. The density profile, $\\rho$, is assumed to depend on the distance from the center of the Be star, $R$, as $\\rho(R) = \\rho_0 (R/R_*)^{-n}$ with a power-law index $n$, where $R_*$ is the star radius and $\\rho_0$ is the density of the outflow at the surface of the star. The flow speed $v(R) = v_0(R/R_*)^{n-2}$ is obtained from conservation of mass flux, where $v_0$ is speed of the outflow at the surface of the star. Then the momentum flux of the flow, $P_{\\rm Be}$, is \\begin{equation} \\label{eq:p_be} P_{\\rm Be}(R) = \\rho v^2 = \\rho_0 v_0^2 (\\frac{R}{R_*})^{n-4}. \\end{equation} In our calculation, indices $n$ of 2.5 and 2 are chosen in outflows of disc and polar wind, respectively.  The location of the shock regime is determined by the balance between pressures of the pulsar wind (Eq.\\,\\ref{eq:p_pw}) and of the Be star outflow (Eq.\\,\\ref{eq:p_be}).  We introduce a new parameter, $x$, defined as \\begin{equation} x = \\frac{\\rho_0}{{\\rm 10}^{-12} {\\rm g cm}^{-3}}\\frac{v_0} {{\\rm 10}^6 {\\rm cm s}^{-1}}. \\end{equation} When $x$ is larger, the location of the pressure balance becomes further from the Be star. If we assume that the opening angle $\\theta_{\\rm disc}$ of the disc outflow is 15\\arcdeg\\,\\citep{johnston96}, the parameter $x$ is related to the $\\Upsilon$ of \\citet{tavani97} by $\\Upsilon \\equiv (\\dot{M}/{\\rm 10}^{-8}{M}_{\\sun}\\ {\\rm yr}^{-1})\\ (v_{0}/{\\rm 10}^6 {\\rm cm\\ s}^{-1}) = {\\rm 0.90} \\times x \\times (v_{0}/{\\rm 10}^6 {\\rm cm\\ s}^{-1})$, where $\\dot{M}$ is the mass loss rate. The parameter $x$ depends on $v_0$ and $\\rho_0$, which are obtained directly from UV/optical observations, independent of the disc opening angle. As shown later, the gamma-ray emission is approximately proportional to $x^2$ in our model.  \\subsection{Particle acceleration and gamma-ray spectrum} \\label{sect:discussion-2} First, we deduce the flux $j_{\\rm i}$ (i $=$ e,p) as a function of energy $E_{\\rm i}$  of the particles in the Be star outflow on the basis of Fermi acceleration, where $e$ denotes electrons and $p$ denotes protons. For simplicity, we assume that all ions are protons. Secondly, we calculate the energy flux from the binary system induced from the emission mechanisms of Bremsstrahlung, inverse Compton (electrons) and proton-proton collisions.  In general terms, the momentum spectrum  $dN/dP$ of the shock-accelerated particles is expressed as $N_{0} P^{-\\alpha}$ and 4$\\pi j$ is obtained from  $v dN/dE_{\\rm k}$, where $E_{\\rm k}$ is the kinetic energy, described as $dN/dE_{\\rm k} = dP/dE_{\\rm k} dN/dP$. The constant $N_{\\rm 0,i}$ is evaluated  with the following integration regarding of the energy balance at the shock location in the Be star flow; \\begin{equation} \\label{eq:energy_balance-general} \\int^{E^{\\rm max_{\\rm i}}}_{m_{\\rm i} c^2}{\\frac{dN_{\\rm i}}{dE_{\\rm i}} (E_{\\rm i}) dE_{\\rm i}} =  f_{\\rm acc, i} P_{\\rm Be} (R_{\\rm shock}), ({\\rm i} = {\\rm e,p}), \\end{equation} where $R_{\\rm shock}$ is the distance to the contact surface from the Be star center. We assume $f_{\\rm acc, i} = $ 0.001 and 0.1 is assumed for $i = e$ and $p$, respectively, as the efficiency of the acceleration to be consistent with an e/p ratio in cosmic ray observations (e.g. \\citet{mullerproc,baring99}). The variation of $R_{\\rm shock}$ causes the orbital modulation in the light curve. The orbital inclination to the line of sight has not been included in our calculation, as this effect is less significant in the Be star emission models than in the pulsar wind emission models. Anisotropy of the optical photons from the Be star is neglected for simplicity. Assuming Fermi acceleration, a power-law index of $\\alpha = -$2.0 is taken for the proton momentum spectrum $dN_{\\rm p}/dP_{\\rm p}$ with an assumed compression ratio of 4.0. The spectral index of the electron momentum spectrum at the shock front does not vary much from the canonical $\\alpha = -$2.0 for plausible values of pulsar wind parameters, because inverse Compton cooling in the higher energy electrons are reduced by the Klein-Nishina effect \\citep{shibazaki02}. In addition, synchrotron cooling does not affect the spectral index, since the magnetic-field strength on the Be-star side should be weak. We therefore assume that the electron spectral index has a constant value of $\\alpha = -$2.0.  The integration is performed from the threshold energy or the particle mass, $m_{\\rm e,p} c^2$, to the maximum energy of the accelerated particle $E^{\\rm max}_{\\rm e,p}$, which we assume here to be $\\sim$10$^{15}$~eV. Applying the obtained $j_{\\rm e,p}(E_{\\rm e,p})$, the gamma-ray spectrum from the source at the distance $D$ is calculated. For the proton-proton collision emission mechanisms, the spectrum is calculated as \\begin{equation} \\label{eq:gamma_flux-protons} F_{\\gamma}^{\\rm pp}(E_{\\gamma}) = \\frac{1}{D^2}\\int{n_{\\rm target} dV} \\int\\int{dE_\\pi dE_{\\rm p}} \\frac{2}{p_\\pi} j_{\\rm p}(E_{\\rm p}) \\frac{d\\sigma_{\\rm pp\\rightarrow\\pi}(E_{\\pi},E_{\\rm p})}{dE_{\\pi}} \\end{equation} where $n_{\\rm target}$ stands for $\\rho/m_{\\rm p}$, and $E_{\\rm p,\\pi}$ and  $p_{\\rm p,\\pi}$ denote the energy and momentum of protons or pions, respectively. Full descriptions of the integral limit and $\\sigma_{\\rm pp\\rightarrow\\pi}$ are given in \\citet{naitotakahara94}. The contributions of the inverse Compton (IC) and Bremsstrahlung are calculated from $j_{\\rm e,p}(E_{\\rm e,p})$ as \\begin{equation} \\label{eq:gamma_flux-electrons} F_{\\gamma}^{\\rm IC, Brem}(E_{\\gamma}) = \\frac{1}{D^2}\\int{n_{\\rm target} dV} \\int^{E^{\\rm max}_{\\rm e}}_{m_{\\rm e} c^2} dE_{\\rm e} j_{\\rm e}(E_{\\rm e}) \\frac{d\\sigma}{dE_{\\gamma}}, \\end{equation} where $n_{\\rm target} = n_{\\rm photon}$ and $\\frac{d\\sigma}{dE_{\\gamma}}$ is a cross section which includes the Klein-Nishina effect for inverse Compton emission, and $n_{\\rm target}$ of $\\rho/m_{\\rm p}$ and the cross section $\\frac{d\\sigma}{dE_{\\gamma}}$ of electron-proton and electron-electron interaction are used for Bremsstrahlung emission \\citep{gaisser98,sturner97}. For $n_{\\rm photon}$, we adopt 2.7~K CMB and $T_{\\rm eff} =$ 27,000~K black body radiation from the Be star. In the spatial integration, we assume that the accelerated particles extend into the Be star outflow downstream of the shock. The contributions of different emission mechanisms are calculated with $x_{\\rm disc} =$1500 for the phase of periastron and their differential energy spectra are shown in Fig.\\,\\ref{fig:spectrum_comparison} (the disc and the pulsar wind are assumed to interact at periastron in the calculation). The total gamma-ray flux is deduced as $F_{\\gamma}(E_{\\gamma}) =  F_{\\gamma}^{\\rm Brem}(E_{\\gamma}) + F_{\\gamma}^{\\rm IC}(E_{\\gamma}) +  F_{\\gamma}^{\\rm pp}(E_{\\gamma})$ and the dominant contribution is of $F_{\\gamma}^{\\rm pp}(E_{\\gamma})$. The inverse Compton flux in the sub-TeV energy region, expected from the pulsar-wind side\\,\\citep{shibazaki02}, is comparable to $F_{\\gamma}^{\\rm IC}$ from the Be star outflows, except that the former has a break $\\sim$400 GeV due to the stronger magnetic field in the pulsar wind side. After the spatial integration, the total flux is approximately expressed as \\begin{equation} \\label{eq:gamma_flux-propt} F_{\\gamma}(E_{\\gamma}, x) \\propto x^2 \\frac{1}{n-1} \\frac{1}{R_{\\rm shock}(x,n,v_0)}~~. \\end{equation} $R_{\\rm shock}(x,n,v_0)$ for the same orbital phase does not vary much within the parameter range discussed in the following.  The adopted model parameters are summarized in Table\\,\\ref{table:model_parameters}. Now we discuss the possible ranges of two parameters, $x$ and the density profile index $n$. For the polar component, the value $x_{\\rm polar}$ is set to be proportional to the disc component, $x_{\\rm disc}$. The factor is estimated using the following two equations, $\\dot{M}_{\\rm polar} =4 \\pi R_*^2~\\rho_{\\rm polar, 0}~v_{\\rm polar, 0}~(1-\\sin \\theta_{\\rm disc})$ and $\\dot{M}_{\\rm disc} =4 \\pi R_*^2~\\rho_{\\rm disc, 0}~v_{\\rm disc, 0}~\\sin \\theta_{\\rm disc}$, where the surface density, initial velocity, and mass loss rate of the disc [polar wind] flow are denoted as $\\rho_{\\rm disc[polar], 0}, v_{\\rm disc[polar], 0},$ and $\\dot{M}_{\\rm disc[polar]}$, respectively, and $\\theta_{\\rm disc}$ denotes the  opening angle of disc outflow. The ratio of two $x$ parameters is \\begin{equation} \\frac{x_{\\rm polar}}{x_{\\rm disc}} = 3.49 \\times 10^{-1} \\frac{\\dot{M}_{\\rm polar}}{\\dot{M}_{\\rm disc}}, \\end{equation} assuming $\\theta_{\\rm disc}$ = 15\\arcdeg \\citep{johnston96}. From the observed intensities of the UV line (due to the polar wind) and of the IR radiation (from the disc), \\citet{lamers87} deduce the mass loss ratio of the two flow components as $\\frac{\\dot{M}_{\\rm polar}}{\\dot{M}_{\\rm disc}}$ of 10$^{-1}$--10$^{-4}$. We take $x_{\\rm polar}$ of 10$^{-1} \\times x_{\\rm disc}$ as a rather optimistic value. The thick disc-like flow is an effective site for the production of gamma-ray emission, and makes the dominant contribution to the total intensity.  For $x_{\\rm disc}$, early studies \\citep{waters86, waters88, dougherty94} have estimated possible ranges of ${\\rm 10}^5 < v_{\\rm disc, 0} < {\\rm 10}^7 {\\rm cm s}^{-1}$  and ${\\rm 10}^{-13} < \\rho_{\\rm disc, 0} < {\\rm 10}^{-9} {\\rm g cm}^{-3}$. Thus we investigate 500 $\\le x_{\\rm disc} \\le$ 5000 in this model analysis.  In Eq.\\,\\ref{eq:p_be},  $n = $2 corresponds to a constant speed and $n = $4 corresponds to a constant ram pressure. We can approximate the polar wind with $n_{\\rm polar} =$2, since the polar wind is generally thought to reach a terminal speed within a few stellar radii. In contrast, the disc density falls rapidly with radius as the rotating material gradually accelerates outward. \\citet{ball99} describes the disc of SS\\,2883 with $n_{\\rm disc} \\lesssim$ 4, while \\citet{waters88} gives 2 $< n_{\\rm disc} <$ 3.25 from a general consideration of Be star discs. We adopt $n_{\\rm disc} = $2.5 here. Changing $n_{\\rm disc}$  to 4 and keeping other parameters fixed reduces the emission by a factor of about 2 (Eq. \\,\\ref{eq:gamma_flux-propt}).  There are no fixed limits for the orbital phases in which the Be star disc outflow interacts with the pulsar wind. We consider three possibilities in calculating the light curves; (i) aligned disc to the orbital plane and interaction throughout the orbit, (ii) mis-aligned disc and interaction in the $\\sim$200-day period around periastron ($\\tau$), during which the radio emission is depolarized, or (iii) mis-aligned disc and interaction in two short periods, [($\\tau - $18~d) $\\sim$ ($\\tau - $8~d)] and [($\\tau +$12~d) $\\sim$ ($\\tau +$22~d)],  as discussed in \\,\\citet{connors02}. Eq.\\,\\ref{eq:gamma_flux-propt} with $x_{\\rm polar}$ of 10$^{-1} \\times x_{\\rm disc}$, suggests that the contribution from the polar-wind--pulsar-wind interaction is a factor 1.5$\\times {\\rm 10}^{-2}$ of that from the disc--pulsar-wind interaction. The polar wind is generally assumed to interact with the pulsar wind at all orbital phases. When the disc and pulsar-wind interaction diminishes. the estimated intensity from the system is only of the polar-wind contribution, and is reduced by a factor of $\\sim {\\rm 10}^{-2}$. We take account of (i), implying the maximum effect of the disc-pulsar wind interaction,  though the disc material becomes dilute at larger distances (Eq.\\,\\ref{eq:p_be}). In (ii) and (iii), we consider emissions from the pulsar-wind bubble formed in the disc flow, after the pulsar leaves the disc\\,\\citep{connors02}. The bubble moves at the velocity $v_{\\rm bubble}$ in the outflow and shock acceleration of particles in the flow proceeds in the contact discontinuity between the bubble and the outflow material. Emissions from the moving bubble are calculated along its trace referring to the material and momentum density profiles of the flow, by replacing  $R_{\\rm shock}(x,n,v_0)$ with $R_{\\rm shock}(t=t_0, x,n,v_0) + v_{\\rm bubble}(t-t_0)$  in Eq.\\,\\ref{eq:gamma_flux-propt}, where $t_0$ denotes the time when the pulsar moves out of the disc flow. We assume an initial value of $v_{\\rm bubble} =$ 100${\\rm km}{\\rm s}^{-1}$ which is larger than the value used in \\citet{connors02}, 15${\\rm km}{\\rm s}^{-1}$, but is similar to the model in \\citet{paredes91} as well as to the typical velocity of the disc flow. The adiabatic expansion, which is mainly important for synchrotron emission, does not affect much the emission mechanism mentioned here. The rise time of bubble emission is assumed to be $\\sim$ 1 day.  \\subsection{Comparison with the Results} The observational upper limits are compared with light curves calculated from the model. The energy thresholds of our results have been scaled to 1\\,TeV assuming a $E^{-2.0}$ spectrum. The spectra calculated with the model assumption (i) in disc-pulsar wind interaction, are integrated  (E $\\ge$ 1\\,TeV)  for four different mass outflow parameters, $x_{\\rm disc} =$ 500, 1000, 1500, and 5000 (Fig.\\,\\ref{fig:model_comparison_x}). The outflow parameter is constrained by our results to $x_{\\rm disc} \\le $1500. The light curves with the different model assumptions (i)--(iii) for the fixed mass outflow parameter $x_{\\rm disc}$ of 1500 are shown in Fig.\\,\\ref{fig:model_comparison_interact}. As discussed in the previous subsection, the light curve is reduced by a factor of $\\sim$ 10$^{-2}$ outside the assumed disc--pulsar-wind interaction period since the polar-wind becomes the only counterpart of the pulsar-wind. In addition, contribution from the wind-bubble formed in the disc--pulsar-wind interaction, remains while the bubble is moving in the disc. Thus, for the model assumption (iii) where the disc and the pulsar-wind interact twice in the orbit, the emission peak after periastron consists of the ``second'' disc--pulsar-wind interaction, of the disc-wind-bubble interaction where the bubble is the outcome of the ``first'' interaction, and of the polar-wind--pulsar-wind interaction. For all three assumptions the constraint from the observations, mainly from $Obs.\\,A$, is similar. With this relatively small outflow pressure, the Be star wind may not be able to overwhelm the pulsar wind pressure to produce accretion onto the pulsar, as has been suggested by the X-ray observations.  Besides our emission models based on the Be-star outflows, we discuss the light curve shown in Fig.\\,5 of \\citet{ballkirk00} as the optimum case for TeV emission from the pulsar wind side, using rather ideal model of the inverse Compton scattering on the un-shocked pulsar wind with a wind Lorentz factor of 10$^7$. Our upper limits are modified into units of integral energy flux using approximated spectral indices in Fig.\\,4 of  \\,\\citet{ballkirk00}, but the obtained limit of $\\sim$1$\\times {\\rm 10}^{-5} {\\rm MeV} {\\rm cm}^{-2}{\\rm s}^{-1}$, does not strongly constrain the model since the light curve quickly declines from $\\sim$5 at the periastron epoch to 0.2 in units of ${\\rm 10}^{-5} {\\rm MeV} {\\rm cm}^{-2}{\\rm s}^{-1}$. The integrated flux greater than 1~TeV is obtained from another model calculation of pulsar wind emission using the spectra in Fig.\\,7 of \\citet{shibazaki02}. They argue for the dominance of synchrotron cooling in the energy loss of the pulsar wind electrons. Assuming the distance of 1.5\\,kpc, the  predicted flux of ${\\rm 10}^{-14} {\\rm cm}^{-2} {\\rm s}^{-1}$  is about two orders of magnitude smaller than our limit.  Recently, new projects of ground based Cherenkov telescopes have begun operations\\,\\citep{review-latest}. With the improved sensitivity and the lower energy threshold, they will offer a better opportunity to observe the PSR\\,B1259$-$63 binary system in the high-energy band. For projects such as CANGAROO-III or H.E.S.S., located in the southern hemisphere, a 50-hour observation of the binary system gives a typical sensitivity of $\\sim {\\rm 10}^{-11} {\\rm cm}^{-2} {\\rm s}^{-1}$ with the energy threshold of $\\sim$100\\,GeV\\,\\citep{hess-performance}. The calculated spectra of our models are integrated again, for the energy greater than 100\\,GeV, for comparison with this sensitivity. In Fig.\\,\\ref{fig:future_comparison}, the sensitivity levels of 20-hour, 10-hour and 5-hour (statistically scaled) observations, respectively, are drawn over the calculated light curves. A day-scale light curve might be detectable for the model with $x_{\\rm disc} \\ge \\sim$700 along the periastron passage. \\citet{balldodd01} estimate the $\\sim$100~GeV emission from the pulsar wind with a Lorentz factor of 10$^6$, and their light curves are compared with these expected sensitivities after modification of the unit into the integral energy flux (${\\rm MeV} {\\rm cm}^{-2}{\\rm s}^{-1}$), assuming the spectral shape (Fig.\\,4 of \\,\\citet{balldodd01}). The light curves in Fig.\\,\\ref{fig:future_comparison} $(right)$ are taken from Fig.\\,5 of \\,\\citet{balldodd01} showing terminated (solid line) and un-terminated (dashed line) shock models in the pulsar wind emissions. Both model predictions  are comparable with the detectable flux, at least, around the periastron epoch. From \\citet{shibazaki02}, integrations $E \\ge$ 100~GeV are performed resulting in fluxes of $\\sim$4$\\times {\\rm 10}^{-12} {\\rm cm}^{-2} {\\rm s}^{-1}$ at periastron and $\\sim$1$\\times {\\rm 10}^{-12} {\\rm cm}^{-2} {\\rm s}^{-1}$ at apastron, which is still below the improved sensitivity of ground based detectors."
    },
    "0402/astro-ph0402022_arXiv.txt": {
        "abstract": "The density profiles of dark halos are usually inferred from the rotation curves of disk galaxies based on the assumption that the gas is a good tracer of the gravitational potential of the galaxies. Some authors have suggested that magnetic pinching forces could alter significantly the rotation curves of spiral galaxies. In contrast to other studies which have concentrated in the vertical structure of the disk, here we focus on the problem of magnetic confinement in the radial direction to bound the magnetic effects on the H\\,{\\sc i} rotation curves. It is shown that azimuthal magnetic fields can hardly speed up the H\\,{\\sc i} disk of galaxies as a whole. In fact, based on virial constraints we show that the contribution of galactic magnetic fields to the rotation curves cannot be larger than $\\sim 10$ km s$^{-1}$ at the outermost point of H\\,{\\sc i} detection, if the galaxies did not contain dark matter at all, and up to $20$ km s$^{-1}$ in the conventional dark halo scenario. The procedure to estimate the maximum effect of magnetic fields is general and applicable to any particular galaxy disk. The inclusion of the surface terms, namely the intergalactic (thermal, magnetic or ram) pressure, does not change our conclusions. Other problems related with the magnetic alternative to dark halos are highlighted. The relevance of magnetic fields in the cuspy problem of dark halos is also discussed. ",
        "introduction": "The rotation curves of spiral galaxies, especially dwarf galaxies and low surface brightness galaxies, are used to derive the profiles of the dark halos which are very useful as tests for the standard cold dark matter scenario (e.g., de Blok \\& Bosma 2002). The usual assumption is that the dynamics of the neutral atomic gas is a good tracer of the radial gravitational force. It can be shown that although spiral waves and non-circular motions may contribute to produce substructure in the rotation curves (e.g., de Blok, Bosma \\& McGaugh 2003), they do not alter the rotation curves significantly for the galaxies selected in these investigations. Galactic magnetic fields could also affect the gas dynamics in spiral galaxies \\citep{pid64,per86,bat92,san97,ben00,bec02,bec03}.  The random small-scale component of the galactic magnetic fields acts as a pressure, giving support to the disk and, therefore, leading to a rotation (slightly) slower than the gravitational circular speed. By contrast, some configurations of the large-scale magnetic field can give rise to a faster circular velocity. \\citet{nel88} and \\citet{bat92} proposed the so-called ``magnetic alternative'' to dark matter in which the contribution of the magnetic pinch could be large enough to explain the observed rotation curves of spiral galaxies without the necessity of dark matter. This is an interesting possibility because some universal properties such as the disk-halo conspiracy or the H\\,{\\sc i}-dark matter constant relation, first noticed by \\citet{bos78}, would have a natural explanation in this scenario \\citep{sn96a}, as well as the truncation of stellar disks \\citep{bat02}. However, after the publication of the paper by \\citet{bat92}, a plethora of possible difficulties faced by the model was pointed out by different authors \\citep{cud93,jok93,per93,val94,kat94,pfe94,sn96b,san97}. In an extensive review, \\citet{bat00} argue that all the shortcomings can be overcome and, therefore, magnetism can explain the way in which galaxies rotate, without the help of dark matter.  A certain MHD configuration has the ability of producing a faster rotation of the whole plasma in the system, i.e.~magnetic confinement, when the net contribution of the magnetic fields to the virial theorem is negative (in the conventional notation). In an isolated system, it is well known that the magnetic field contribution to the virial theorem, applied to a large enough volume, is {\\it positive or zero}, but never negative, reflecting the net expansive tendency of magnetic fields (e.g., Shafranov 1966). \\citet{jok93} used this argument to conclude that in the presence of magnetic fields, galaxies would need even more dark matter than in the unmagnetized case. One could argue that, since the virial theorem uses global (volume-integrated) variables, it is possible that magnetism has a ``centrifugal'' action in the vertical direction but centripetal in the radial direction (radial confinement) and, therefore, the expansive tendency of magnetic fields manifests only in the vertical direction. We advance that this claim is very wrong. An alternative way out is to assume that galaxies are not at isolation but embedded in the ubiquitous (magnetized) intergalactic medium which would be responsible for the required radial confinement.  Motivated by these ideas and given the importance of a correct interpretation of the rotation curves for Cosmology, we will try to put bounds on the effects of magnetic fields on the H\\,{\\sc i} rotation curves. However, we do not restrict ourselves to discuss the viability of explaining rotation curves without dark halos; the orbits of Galactic globular clusters, the satellite galaxies and possibly the dynamics of remote stars in the outer Galactic halo, which are not affected by magnetic fields, strongly argue in favor of a dark halo around the Milky Way (e.g.~Freeman 1997). The main issue is to see whether galactic magnetic fields can alter appreciably the H\\,{\\sc i} rotation curves, especially in external galaxies for which the dark matter content is almost entirely derived from the rotation curve of the gas. Only if the answer is negative, we can be confident that the profiles of dark halos, the dark matter content in galaxies as well as other interesting empirical correlations found between, for instance, halo parameters and luminosity are reliable.  Our investigation will be focus on the {\\it radial} equilibrium configuration of magnetized rotating disks. For reasons that should become clear later, the analysis of the radial equation of motion puts stringent bounds on the maximum magnetic effects to the rotation curves; this being free of other complicating assumptions that appear in studies of the vertical equilibrium configuration such as the external pressure, the vertical dependence of the magnetic field within the height of the H\\,{\\sc i} thin disk, or the flattening of the dark halo, which are all very uncertain.  The paper is organized as follows. In \\S \\ref{sec:ingredients} we describe the ingredients of a simplified model in which magnetic fields could give rise to a faster rotation in the outer parts of H\\,{\\sc i} disks. In \\S \\ref{sec:core} the maximum effect on the rotation curves is studied. Other difficulties inherent to the magnetic alternative are discussed in \\S \\ref{sec:caveats}. The case of magnetic fields producing a slower rotation and their role as a possible remedy to the cuspy problem of dark halos is left to \\S \\ref{sec:dwarf} and the final remarks are given in \\S \\ref{sec:final}. ",
        "conclusions": ""
    },
    "0402/astro-ph0402508_arXiv.txt": {
        "abstract": "A combined analysis of the profiles of the main broad quasar emission lines in both {\\it Hubble Space Telescope} and optical spectra shows that while the profiles of the strong UV lines are quite similar, there is frequently a strong increase in the Ly$\\alpha$/H$\\alpha$ ratio in the high-velocity gas. We show that the suggestion that the high velocity gas is optically-thin presents many problems. We show that the relative strengths of the high velocity wings arise naturally in an optically-thick BLR component. An optically-thick model successfully explains the equivalent widths of the lines, the Ly$\\alpha$/H$\\alpha$ ratios and flatter Balmer decrements in the line wings, the strengths of C\\,{\\sc III}] and the $\\lambda$1400 blend, and the strong variability of high-velocity, high-ionization lines (especially He\\,{\\sc II} and He\\,{\\sc I}). ",
        "introduction": "The nature and origin of the broad-line region (BLR) in active galactic nuclei (which we will refer to simply as ``quasars'') has been a long-standing problem. It has been known for some time that the shapes of broad emission lines differ from line to line within the same object depending on the ionization level of the species producing the each line.  For example, He\\,I $\\lambda$5876 is broader than H$\\alpha$ (Osterbrock \\& Shuder, 1982), O\\,I $\\lambda$1304 and C\\,II] $\\lambda$2326 are narrower than C\\,IV (Wilkes 1984), C\\,IV $\\lambda$1549 tends to be broader than Mg\\,II $\\lambda$2798 (Mathews \\& Wampler, 1985) and H$\\alpha$ is narrower than H$\\beta$ (Osterbrock 1977). H$\\alpha$ and H$\\beta$ have weaker wings than Ly$\\alpha$ (Zheng 1992; Netzer et al. 1995)\\footnote{It can be misleading to use FWHM as a measure of how broad emission lines are in the wings. Large contributions to the core can skew the FWHM to lower values. For example, although the line profile ratios reveal that the Ly$\\alpha$ wings are broader, Table 1 of Netzer et al. (1995) shows that FWHM$_{H \\beta} > FWHM_{Ly\\alpha}$}. These differences in line profiles imply differences in physical conditions as a function of velocity and probably as a function of distance from the central ionizing source. One of many unresolved BLR questions is whether these differences can be explained by one BLR with a range of conditions as a function of radius, or whether there are two (or more) fundamentally distinct components. It is common to discuss two separate components, but there are differences in the literature as to how to divide the BLR into two components.  In disk-plus-wind models (e.g., Chiang \\& Murray 1996; Bottorff et al. 1997) the two components are identified with emission from near the disk and emission from a wind above the disk. Some observers and phenomenologists have made distinctions primarily on the basis of the mean degree of ionization, others primarily on the basis of Doppler widths. There is substantial overlap in these classifications. Gaskell (1987) and Collin-Souffrin \\& Lasota (1988) make a division, motivated by profile differences and photoionization modelling issues, into two components, one with a typical nebular spectrum (called ``BLR I'' by Gaskell 1987, and ``HIL'' by Collin-Souffrin \\& Lasota) and one with substantial emission from clouds with large partially-ionized zones (``BLR II'' or the ``LIL'').  The case for two such components is summarized by Gaskell (2000).  Other workers, motivated by analyses of line profiles, have separated the BLR into a ``very broad line region'' (VBLR) and another component. The other component has been called the ``intermediate-line region'' (ILR) by Wills et al. (1993) and Brotherton et al. (1994) or the ``classical broad component'' (BC) by Sulentic, Marziani, \\& Dultzin-Hacyan (2000).  Note that the BC and ILR are not identical and hence the corresponding VBLR components can be quite different. The VBLR, ILR, and BC show both high and low ionization emission. Despite these differing terminologies, there is general agreement that the conditions of the highest velocity gas (which we will call the VBLR) are different from the other components although the VBLR gas might merely be an extension of the other gas. Historically, photoionization models assumed that line-emitting gas was optically-thick. However, several authors have argued for a significant optically-thin component to the VBLR gas, and to date, there has been no resolution to this debate. The main arguments for an optically-thin VBLR have been: \\begin{enumerate} \\item {\\it The general symmetry of Ly$\\alpha$.} Wilkes \\& Carswell (1982) pointed out that BLR clouds could not be both optically-thin and have a net radial motion (e.g., outflow) because Ly$\\alpha$ profiles are symmetric and similar to C\\,IV profiles.  Since the blueshifting of the high-ionization lines requires at least some radial motions (Gaskell 1982), the relative symmetry of Ly$\\alpha$ could mean that the VBLR (highest velocity gas associated with the profile wings) has to be optically-thin. \\item {\\it An apparent lack of variability of the VBLR in Mrk~590 between two epochs.} Ferland, Korista, \\& Peterson (1990) found that the VBLR Balmer emission line flux in Mrk~590 changed little between two epochs about three years apart even though the continuum and line cores changed. Hydrogen recombination lines should change little if the VBLR is optically-thin and hence fully ionized. \\item {\\it The Ly$\\alpha$/H$\\beta$ ratio is higher for the VBLR}. Zheng (1992) showed that for some AGNs the Ly$\\alpha$/H$\\beta$ ratio is much higher in the wings and that it approached the Menzel-Baker case B value that might be expected from optically-thin gas. \\item {\\it Emission line responses might require negative responsivities.} Sparke (1993) pointed out that the shapes of some of the line-continuum cross-correlation functions from 1989 monitoring of NGC\\,5548 seem to require that the emission of the inner BLR declines as the continuum increases. She suggested that optically-thin clouds in the inner BLR could do this. \\end{enumerate} The advent of the {\\it Hubble Space Telescope} (HST) now makes possible the comparison of high-quality line profiles of both the high-ionization UV lines and the Balmer lines in the same objects. In this paper we use such comparisons to investigate the nature of the VBLR and we will argue that the bulk of the VBLR is in fact {\\it not} optically-thin.  In paper II (Snedden \\& Gaskell, in preparation) we discuss the physical conditions in the BLR gas as a function of velocity. ",
        "conclusions": "We conclude that rather than being optically-thin, the high-velocity BLR gas is predominantly optically-thick.  We believe this because optically-thin models fail to explain the following properties of the high-velocity gas: \\begin{enumerate} \\item the equivalent widths, \\item the Ly$\\alpha$/H$\\alpha$ ratio, \\item the Balmer decrement, \\item the similarities of line profiles, \\item the strong variability of the high-ionization lines, \\item the C\\,{\\sc III}]/C\\,{\\sc IV} ratios, and \\item the strength of the $\\lambda$1400 blend. \\end{enumerate} If the high-velocity gas is predominantly optically-thick it is thus not fundamentally different, in terms of optical depth, from the lower-velocity gas."
    },
    "0402/hep-th0402079_arXiv.txt": {
        "abstract": "We consider a generalisation of the DGP model, by adding a second brane with localised curvature, and allowing for a bulk cosmological constant and brane tensions. We study radion and graviton fluctuations in detail, enabling us to check for ghosts and tachyons. By tuning our parameters accordingly,  we find bigravity models that are free from ghosts and tachyons. These models will lead to large distance modifications of gravity that could be observable in the near future. ",
        "introduction": "Where does the force of gravity come from? Anyone unfortunate enough to be near a strongly gravitating object, such as a black hole, would say ``from the curvature of spacetime''. In a weaker gravitational field, we measure fluctuations about some background spacetime~\\cite{Fierz:pauli-fierz}. These fluctuations correspond to spin-2 particles called gravitons. Traditionally, we now say that ``gravity comes from the exchange of {\\it massless} gravitons''. We believe  this because it reproduces Newton's Law of gravitation, and that is well tested experimentally. But how well is Newton's Law really tested? The truth is that our experimental knowledge of gravity only covers distances between $0.2$mm and $10^{26}$cm. The lower bound is as a result of Cavendish experiments, and the upper bound corresponds to 1\\%  of the current Hubble length. In units of $\\bar h=c=1$, this is equivalent to an energy scale \\be \\label{range} 10^{-31}~\\textrm{eV} < p < 10^{-3} ~\\textrm{eV} \\ee Perhaps, therefore, we should alter our previous statement. Gravity at the experimental scale (\\ref{range}) could be due to the exchange of massless and/or massive particles, so long as the masses are less than $10^{-31}~\\textrm{eV}$.  Indeed, massive gravitons appear automatically in some higher derivative gravity theories~\\cite{Hindawi:spin2}. The simplest non-trivial scenario would contain a single massive graviton~\\cite{Fierz:pauli-fierz}. If we have a combination of gravitons of many different masses,  we have {\\it multigravity}~\\cite{Kogan:2000, Kogan:branemulti, Kogan:modification, Kogan:6dmulti,Kogan:review, Papazoglou:multithesis}. In this paper, we will encounter {\\it bigravity}~\\cite{Kogan:adsbranes, Damour:univclass}. This is the simplest example of multigravity,  in which gravity is mediated by both a massless graviton and a single {\\it ultralight} graviton.  Multigravity is clearly of interest from a purely theoretical point of view. However, it is also of interest to phenomenologists because it predicts new gravitational  physics at very large distances. At distances beyond the Compton wavelength of the ultralight mode, the ultralight mode is turned off, and gravity is mediated by the massless mode alone. Large distance modifications of gravity have been {\\it in vogue} recently, as they could offer an explanation to the current acceleration of the universe~\\cite{Perlmutter:acc, Riess:acc, Deffayet:DGPcosmo, Damour:nonlin, Lue:dark,Lue:cosmic}.  Unfortunately, there are number of problems with many existing models of modified gravity. One very serious problem is the presence of ghosts (see, for example, ~\\cite{Kogan:2000, Gregory:GRS, Pilo:ghost, Dubovsky:ghost, Dubovsky:DGP}). Although the ``AdS brane'' model in~\\cite{Kogan:adsbranes} is ghost-free, modifications of gravity are hidden behind the AdS horizon. The six dimensional model of Kogan {\\it et al}~\\cite{Kogan:6dmulti} is also thought to be ghost-free, although this has not been confirmed as the model is too difficult to work with. The DGP model~\\cite{Dvali:DGPmodel}, meanwhile, is simple, ghost-free, and even predicts new infra-red physics, that could one day be observable. However, it is {\\it not} a bigravity model; gravity is due to a resonance of continuum modes.  Let us describe the DGP model in more detail. It is given by the following action \\be \\label{DGPaction} S_\\textrm{DGP}=M^3 \\int_\\textrm{bulk} \\sqrt{g} R(g)+m_{pl}^2 \\int_\\textrm{brane} \\sqrt{\\gamma}R(\\gamma) \\ee In other words, we have a four-dimensional Minkowski brane embedded in a five-dimensional Minkowski bulk. The key ingredient is the brane localised curvature, $R(\\ga)$. This could be generated by quantum corrections, if matter were present on the brane. Localised curvature can also appear in string theory~\\cite{Corley:EH,Antoniadis:CY}. In the DGP model, Newton's Law is reproduced up to a distance $r=m_{pl}^2/2M^3$.  Beyond this scale the behaviour of gravity is five-dimensional.  In this paper, we will consider a generalisation of the DGP model. We will add a second brane with localised curvature. We will also allow for a bulk cosmological constant, and for the branes to have tension. Since the extra dimension is finite, we will get a discrete graviton mass spectrum. Our aim is to ask the following question: is it possible to obtain an interesting modified theory of gravity {\\it without} introducing ghosts and tachyons? The answer will be ``yes''. For certain parameter regions, we will discover bigravity models that are tachyon and ghost-free, leading to potentially  observable new physics in the infra-red. It is interesting to note that a single DGP-like brane in a compact extra dimension also has a discrete mass spectrum, but does {\\it not} exhibit bigravity~\\cite{Dvali:power}.  The rest of this paper will be organised as follows: in section~\\ref{sec:setup}, we will describe our set up in more detail, giving the bulk and boundary equations of motion. In section~\\ref{sec:background}, we  will derive solutions for the background spacetime. We will perturb about this background in section~\\ref{sec:pert}, arriving at the linearised equations of motion. In section~\\ref{sec:radion} we will focus on the {\\it radion} mode. This corresponds to fluctuations in the brane separation. We will calculate its effective action to quadratic order, in order to determine whether or not it is a ghost. For de Sitter branes we will find that the radion is tachyonic. In section~\\ref{sec:spectrum}, we  will turn our attention to the gravitons. We will derive the mass spectrum, and show that we can have {\\it observable} bigravity. We will also check for ghosts by calculating the graviton effective action. In section~\\ref{sec:analysis}, we will play around with our parameters until we find a bigravity model that is ghost-free, and might one day lead to observable, new, infra-red physics. Section~\\ref{sec:conclusions} will contain some final remarks. In particular, we will discuss some important issues, such as the famous van Dam-Veltman-Zakharov (vDVZ) discontinuity~\\cite{vanDam:VDVZ, Zakharov:VDVZ}, and the recently discovered strong coupling problems in massive gravity~\\cite{Arkani-Hamed:massive, Schwartz:decon}. Finally, the appendix contains a detailed calculation of the mass spectrum for AdS branes. ",
        "conclusions": "\\label{sec:conclusions} In this paper we have studied a class of braneworld models, with localised curvature on the branes. Some of these models give rise to bigravity, leading to large distance modifications of gravity for a four-dimensional observer. After checking for ghosts and tachyons, we rejected most, but not all of the models we had considered.  One of these well behaved models consists of a flat bulk ($\\La=0$) sandwiched in between two flat branes ($\\ka=0$) of zero tension. The branes are close together, and the curvature terms on the brane are large and positive ($r_i \\gg l$). The mass spectrum is precisely that of bigravity. We have a massless graviton, an ultralight graviton, and a tower of heavy KK modes. The model is completely free of ghosts and tachyons.  To add some substance, let us put some numbers in. If we take $l \\sim 10^3 ~(\\textrm{eV})^{-1}$, and $r_i \\sim 10^{59} ~(\\textrm{eV})^{-1}$, we find that \\be m_\\textrm{light} \\sim 10^{-31} ~\\textrm{eV}, \\qquad m_\\textrm{heavy} \\gtrsim 10^{-3} ~\\textrm{eV}. \\ee These masses lie outside of the range for which gravity is well tested (\\ref{range}), so we have no contradiction with experiment. Since the four-dimensional Planck mass, $m_{pl} \\sim 10^{18}~\\textrm{GeV}$, we conclude from equation (\\ref{Planck}) that $M \\sim 10^{-2}-10^{-1}~\\textrm{eV}$. While the fundamental Planck scale is very low, it does not violate the validity bound (\\ref{eqn:quantumbound}), and exceeds the scale probed by Cavendish experiments ($10^{-3}~\\textrm{eV}$). In principle there are cosmological and astrophysical bounds that one should consider, such as the effect on star cooling due to graviton emmision into the bulk. This will be left for future research.  Note that $m_\\textrm{light} \\sim 10^{-31} ~\\textrm{eV}$ is just about heavy enough for us to expect interesting observations today. This could be important in trying to explain cosmic acceleration. If we wanted to reduce this mass even further, we would have to increase $r$. Even a small increase would push the Planck mass below the Cavendish scale. Curiously, we seem to be in just the right place to start seeing  interesting new physics, either in the infra-red, or the ultra-violet\\footnote{I would like to thank John March-Russell for this observation}.  We can go beyond this model by switching on a small AdS curvature in the bulk, and still get ghost-free bigravity. Although a tachyon appears for de Sitter branes, this is not the case for flat and anti-de Sitter branes. When the bulk AdS length is of the same order as the brane separation, we find that ghost-free bigravity becomes impossible.  By checking for ghosts and tachyons in our models, we have carried out the first important tests of viability. However, we should be aware that there are other possible problems. As with all models of massive gravity, we should be concerned with the famous vDVZ discontinuity~\\cite{vanDam:VDVZ, Zakharov:VDVZ}. Our linearised equations of motion (\\ref{eqn:linbc}) suggest that this may well be an issue. This is especially true for the case $\\La=\\ka=0$, as there is no ``brane-bending'' effect that could cancel off any unwanted degrees of freedom (see, for example, ~\\cite{Giddings:lingrav, Csaki:props}).  Even our AdS brane model will suffer this problem. This is a surprise, as it is often said that there is no vDVZ discontinuity in AdS space~\\cite{Kogan:VDVZ, Porrati:VDVZ}. However, this result relies on the graviton mass going to zero faster than the inverse horizon size. In order to ensure that we had {\\it observable} bigravity, we demanded the opposite of this.  Crucially, there may be a way around this problem. Vainshtein {\\it et al}~\\cite{Vainshtein:nonVDVZ,Deffayet:noVDVZ}, have argued that the perturbative expansion in Newton's constant in~\\cite{vanDam:VDVZ, Zakharov:VDVZ}, is inconsistent, as the graviton mass goes to zero. Specifically, the standard linearised analysis near a heavy source is only valid at distances \\be r \\geqslant \\left(\\frac{r_M}{m^4}\\right)^{\\frac{1}{5}} \\ee where $r_M$ is the Schwarzschild radius of the source, and $m $ is the graviton mass. To see if a mass discontinuity really is present, we need to calculate the Schwarzschild solution on a brane, and compare it to the standard four-dimensional massless result. This is a highly non-trivial exercise that is beyond the scope of this paper.  For Pauli-Fierz theory~\\cite{Fierz:pauli-fierz}, the breakdown of the linearised analysis has been linked to strong coupling phenomena at the following scale~\\cite{Arkani-Hamed:massive} \\be E_\\textrm{strong} \\sim \\left(\\frac{l_{pl}}{m^4}\\right)^{-\\frac{1}{5}} \\ee where $l_{pl}$ is the Planck length\\footnote{For certain non-linear extensions of Pauli-Fierz, $E_\\textrm{strong} \\sim \\left(l_{pl}/m^2\\right)^{-\\frac{1}{3}}$~\\cite{Schwartz:decon}. This higher scale also appears in the DGP model~\\cite{Luty:strong}.}.  Large terms in the full propagator tend to signal this problem.  Unfortunately, we might expect our theory to suffer the same fate. We can see this by looking at $h_{zz}$ in a fixed wall gauge (see equation (\\ref{fixed})). By the mean value theorem, there exists $z_0 \\in [ 0, l]$ such that $B^\\prime (z_0)=1/l$. For small $l$, it is clear that $h_{zz}$ can be very large.  A similar strong coupling scale also  exists in the DGP model~\\cite{Luty:strong,Rubakov:strong}. However, it has recently been argued that this scale could be unphysical, and is just a result of the naive perturbative expansion~\\cite{Dvali:IR}. Clearly, both the mass discontiniuty and strong coupling problems are highly contentious issues at the moment. For this reason we have focused on the possible existence of ghosts in our models. There is no contention there: ghosts are undesirable, but can be avoided.  There is still much to do. It would be very interesting to study the phenomenology of these models in more detail, particularly in the context of cosmic acceleration. Coincidentally, we have been forced to choose a brane separation $l \\sim 10^3 ~(\\textrm{eV})^{-1}$, which agrees with the brane separation chosen in~\\cite{Cognola:multi}. In~\\cite{Cognola:multi}, the mass spectrum is motivated by a discretized Randall-Sundrum model, with $l$ chosen so that there is a small effective cosmological constant. Would the presence of an ultralight mode in the spectrum affect these results? We would probably expect the answer to be ``no'', because the dominant contribution to the vacuum energy would still come from the heavier modes. Nevertheless, it is worthy of further investigation.  Our models could certainly be generalised in a number of ways. We could  consider more branes, abandon $\\mathbb{Z}_2$ symmetry, or even introduce higher derivative terms in the bulk and on the brane. Branes embedded in solutions to Gauss-Bonnet gravity have been the subject of much research recently (see, for example~\\cite{Davis:junction, Padilla:gbholog}), motivated by the link to string theory. Indeed, it would also be nice if our models, or at least some generalisation, could be derived from a more fundamental theory~\\cite{Corley:EH,Antoniadis:CY}. However, the high degree of fine tuning could prove an obstacle in this respect.  Let us end by summarising our main result: we have discovered braneworld models that exhibit bigravity, without introducing ghosts. Recall that bigravity naturally gives rise to new gravitational physics in the infra-red. The new physics occurs when the massive graviton ``switches off'', at distances beyond its Compton wavelength. Our models are an improvement on the ghost-free model given in~\\cite{Kogan:adsbranes}, because they lead to potentially observable modifications of gravity. In~\\cite{Kogan:adsbranes}, all modifications are hidden behind the AdS horizon.       \\vskip .5in \\centerline{\\bf Acknowledgements} \\medskip I would like to thank John March-Russell, Valery Rubakov, Graham Ross, Syksy R\\\"as\\\"anen, and Ben Gripaios for  helpful dicussions. In particular, I would like to thank John and Graham for proof reading this article. Thanks must also go to Ben, Ro, Ash and Perks for being a constant source of inspiration, Leppos and Beyonc\\'e. And to Bruno Cheyrou for being the new Zidane. AP was funded by PPARC."
    },
    "0402/astro-ph0402058_arXiv.txt": {
        "abstract": "s{The discovery of a supernova emerging at late times in the afterglow of GRB~030329 has apparently settled the issue on the nature of the progenitor of gamma-ray bursts. We now know that at least a fraction of cosmological GRBs are associated with the death of massive stars, and that the two explosions are most likely simultaneous. Even though the association was already suggested for GRB~980425, the peculiarity of that burst did not allow to extend the association to all GRBs. The issue is now to understand whether GRB~030329 is a ``standard burst'' or not. I will discuss some peculiarities of GRB~030329 and its afterglow lightcurve showing how, rather than a classical cosmological GRB, it looks more like a transition object linking weak events like GRB~980425 to the classical long duration GRBs. I will also discuss the problems faced by the Hypernova scenario to account for the X-ray features detected in several GRBs and their afterglows.} ",
        "introduction": "The progenitor of Gamma-Ray Bursts (GRBs) - i.e. the astronomical object that is associated to the energy release powering the GRB emission, and that probably disappears in the process - has been a major unknown since their discovery in the late sixties\\cite{kle}. Since the discovery of afterglow emission\\cite{cos,van}, that made possible their precise localisation in the sky, several circumstantial pieces of evidence have been collected, linking the burst emission with massive star formation phenomena\\cite{kul,fru,hol,blo,laz1}. Even though such studies pointed toward an association of GRB explosions with the death of massive stars, the issue was far from being solved, the main worry being that multi-wavelength modelling of afterglows yielded typically a uniform ambient medium\\cite{pan}, contrary to the stratified wind expectations for the massive star association\\cite{che}.  The exact association of the burst with the star death was therefore put under debate. The simplest scenario (Hypernova or Collapsar) would call for a single explosion, in which the GRB would be produced by a relativistic jet propagating along the rotational axis of a fast spinning star, which explodes as a more normal supernova along the equator\\cite{mac}. Alternatively, the two explosions (SN and GRB) could be separated by a relatively short interval of time (Supranova), during which a meta-stable compact object is left, whose eventual collapse cause the GRB explosion\\cite{vie}. Finally, new population synthesis studies showed that neutron star binary systems may be short lived, allowing for the GRB explosion within the host galaxy\\cite{per} even in the classical binary merger scenario\\cite{eic}.  The association of the peculiar GRB~980425 with SN1998bw\\cite{gal} and, more recently, that of GRB~030329 with SN2003dh\\cite{sta,hjo} and of GRB021211 with SN2002lt\\cite{del} has given strong support to the idea of a single explosion, which simultaneously generates a GRB and a particularly energetic SN explosion. There is however a number of observations that seem to point to a more complex association. The issue is therefore to understand whether GRB~980425 and GRB~030329 are ``classical GRBs'' and to which extent all the observations can be incorporated in a single coherent picture.  In this review I will critically discuss some aspects of this problem, underlying possible peculiarities of the prompt and afterglow emission of the two bursts robustly associated to SN explosions. I will then discuss progenitor indications from X-ray spectroscopy of the prompt and afterglow emission, discussing the problems that arise when we attempt to include these observations in a simple Hypernova scenario ",
        "conclusions": "In this paper I have tried to address two questions: is GRB~030329 a typical GRB so that we can safely claim that all GRBs are associated to SNe and that the time delay between the two explosions is negligible? And the second: is there any evidence that is inconsistent with the above conclusion from independent observations?  The answer to the first question is that indeed GRB~030329 is much more similar to a classical cosmological GRB that GRB~980425, the first to be associated to the explosion of a massive star. The similarity of GRB~030329 with ``classical GRBs'' is however not complete. First there is evidence that the energy released by GRB~030329 in gamma-rays is smaller than usual by an order of magnitude, even though taking into account the energy released by the inner engine in the form of less relativistic material brings back the total energy budget to the''standard'' $E=10^{51}$~erg observed in cosmological GRB explosions\\cite{fra,pan}. Secondly, and possibly related to the delayed energy input of energy, the afterglow lightcurve is much more complex than any previously observed GRB, and we showed that this is not due to the more complete sampling, but to an intrinsic variability that is unprecedented in cosmological GRBs.  An intriguing possibility is that there is a standard inner engine, possibly a black hole surrounded by a dense hot accretion disk, which releases a fairly standard energy in the form of a relativistic jet. The jet has however to open its way into the star and this creates the diversity in the observed properties of GRBs. Different relevant properties of the star may be related to its pre-explosion radius and/or rotation. A compact fast spinning star should offer less resistance to the jet propagation along its polar axis, giving origin to a cosmological GRB, in which most of the jet energy can escape untouched and produce $\\gamma$-ray radiation. A more extended or less rotating star may offer more resistance. In this case a sizable fraction of the jet energy should be used to open up a funnel in the star, so that the resulting GRB would look under-energetic. Part, if not all, this energy may be recycled in a delayed slower fireball component that would catch up with the relativistic jet and re-energise it at later times.  This unification picture shall however be taken with care, at least until a final word is said about the reality of X-ray absorption and emission features. These cannot be easily incorporated in a single SN/GRB explosion scenario. Even though none of the claimed features is incontrovertible in terms of statistical significance, they form a consistent set of observations all naturally predicted in the two explosion Supranova scenario\\cite{vie}. Instead of having a variable stellar radius and/or rotation, a unification picture may call for a variable delay between the two explosions. Some delays are very short, and produce under-energetic GRBs, since part of the jet energy is wasted to reach the star surface. Short ($<$~several weeks) delays do not produce detectable $\\gamma$-rays, since the jet is completely choked inside the optically thick SN remnant\\cite{gue}. Longer delays would instead produce a ``classical GRB'', since the jet has no baryonic material to cross.  Also this simple scenario faces however some problems, since there are many ``classical GRBs'' that show sign of red bumps at late times, usually identified as the emergence of the SN lightcurve on top of the power-law afterglow decay\\cite{blo2,rei,gal2,laz7,blo3}. These SN explosion should be simultaneous to the GRB explosion. A unification scenario that would take into account all the observations with their most probable interpretations would therefore need to take into account the possibility of both a variability in the progenitor star properties and of the explosion delay.  \\bigskip  {\\it I wish to thank S. Covino, F. Frontera, A. K\\\"onigl, G. Ghisellini, P. Mazzali, R. Perna, E. Pian, M. J. Rees, E. Rossi, L. Stella and M. Vietri for the fruitful collaboration and discussions that led to the development of many of the ideas presented in this paper.}"
    },
    "0402/astro-ph0402572_arXiv.txt": {
        "abstract": "Gemini Observatory's northern and southern telescopes are both presently being outfitted with facility mid-infrared imagers/spectrometers. This will allow observers the unique opportunity to apply to one observatory for all-sky spectroscopic access in the mid-infrared with the light gathering power of 8-meter telescopes. Gemini South has recently commissioned the Thermal-Region Camera and Spectrograph (T-ReCS) and is now available to perform queue observations for the community. T-ReCS is capable of low-resolution long-slit spectroscopy of R$\\sim$100 near 10 and 20 $\\mu$m, and medium-resolution long-slit spectroscopy of R$\\sim$1000 near 10 $\\mu$m. Gemini North is presently commissioning Michelle, which will be capable of low-, medium-, and high-resolution long-slit spectroscopy of R$\\sim$200, 1000, and 3000, respectively, near 10 $\\mu$m, as well as low-resolution long-slit spectroscopy of R$\\sim$200 near 20 $\\mu$m. Michelle can also perform echelle spectroscopy of R$\\sim$30000 at 10 and 20 $\\mu$m. The low-, medium-, and high-resolution spectroscopic modes of Michelle will be available to the public for the fall semester of 2004, and the echelle mode is expected to be available in 2005. ",
        "introduction": "The twin 8-m telescopes of Gemini Observatory are located in Hawaii in the northern hemisphere, and Chile in the southern hemisphere. Both telescopes are being outfitted with mid-infrared cameras/spectrometers, thus allowing observers the unique opportunity to apply for all-sky observations of mid-infrared targets. Another way in which Gemini sets itself apart from other 8-m class telescope facilities is its optimization for infrared observations.  Ideally, a mirror coating should reflect as close to 100\\% of the light that strikes it. In the thermal infrared region of the spectrum, however, a mirror coating emits a great deal of infrared radiation, and the statistical noisiness of this emission reduces an infrared instrument's sensitivity to astronomical sources. The emissivity of an optical surface is defined as the ratio between its level of emission and that of a perfect blackbody emitter, and it is roughly the inverse of the reflectivity. Therefore, in order to maximize infrared sensitivity, the mirror emissivity must be as low as possible.  The Gemini telescopes employ single monolithic primary mirrors, rather than segmented ones whose intersegment gaps increase overall telescope emissivity. Furthermore, Gemini plans on coating primary and secondary mirror surfaces of both telescopes with silver (rather than the usual aluminum). Silver has a much lower emissivity than aluminum in the infrared, translating to unprecedented sensitivities for Gemini's infrared instruments.  Gemini South's secondary mirror is already coated in silver, yielding a total system emissivity of 3.0\\% at 9 $\\mu$m. In May of 2004, the primary mirror of Gemini South will also be coated in silver, bringing the total emissivity down to an estimated 2.2\\% or below. Gemini North also plans to be fully silver coated by the end of 2005. ",
        "conclusions": "With their versatile mid-infrared observing capabilities, T-ReCS and Michelle will empower astronomers with the tools needed to explore the nature of a broad range of astronomical objects and environments. The large infrared-optimized collecting area provided by the state-of-the-art Gemini telescopes, coupled with the large throughputs and excellent optics of T-ReCS and Michelle, are a powerful combination that will surely advance the field of mid-infrared astronomy."
    },
    "0402/astro-ph0402091_arXiv.txt": {
        "abstract": "We examine the selection effects that determine how the population of inspiraling binary compact objects (BCOs) is reflected by those potentially observed with ground-based interferometers like LIGO.  We lay the ground-work for the interpretation of future observations in terms of constraints on the real population and, correspondingly, binary star evolution models.  To determine the extra-galactic population of inspiraling binaries we combine data on distance and blue luminosity from galaxy catalogs with current models of the galactic BCO mass distribution to simulate the physical distribution of binaries in the nearby universe. We use Monte Carlo methods to determine the fraction of binaries observable by the LIGO detectors from each galaxy as a function of the BCO chirp mass. We examine separately the role of source distance, sky position, time of detection, and binary system chirp mass on detection efficiency and selection effects relevant to the three LIGO detectors.  Finally, we discuss the implications of the nearby geography of space on anticipated GW detection and compare our results to previous studies, which have assumed uniform galaxy volume density and fixed chirp mass for binary compact objects.  From these considerations, actual BCO inspiral observations or significant upper limits on the coalescence rate anticipated in the near future by ground-based interferometers can be used to improve our knowledge of the galactic binary inspiral rate and to constrain models of radio pulsar characteristics and binary star evolution channels leading to neutron star or black hole binaries. ",
        "introduction": "\\label{sec:intro}  Binary compact objects (BCO) with neutron stars or black holes hold a special place among gravitational wave (GW) sources. The discovery of PSR~B1913+16~\\citep{ht}, the first binary pulsar system, inspired the detailed study of inspiraling compact binaries and provided the first observational evidence for the existence of gravitational radiation~\\citep{TW89}. Binary systems like PSR~B1913+16 are driven to coalescence by a GW emission catastrophe: in the last approximately 20\\,s before they coalesce they radiate their remaining binding energy (approximately $2\\times10^{52}$~ergs) as gravitational waves in a band accessible to the large ground-based detectors like the Laser Interferometer Gravitational Wave Observatory~\\citep[LIGO;][]{ligoR} and VIRGO~\\citep{virgo}\\footnote{Somewhat less energy will be radiated in the bands accessible GEO \\citep{ligoR} and TAMA \\citep{TAMA}, and less still in the bands accessible to the bar detectors AURIGA \\citep{AURIGA}, ALLEGRO \\citep{ALLEGRO}, EXPLORER \\citep{EXPLORER1,EXPLORER2} and NAUTILUS \\citep{NAUTILUS}.}.  Current observational constraints on the population of neutron star or black hole binary systems depend on radio pulsar observations of just a handful of Galactic binary systems \\citep[e.g.,][]{burgay,kal03}. In contrast, the LIGO and VIRGO detectors will observe stellar mass inspiraling BCOs at extra-galactic distances. They will also be sensitive to black hole binaries, which are not observable electromagnetically. Correspondingly, observations by this new generation of detectors can help constrain the binary coalescence rate density in the nearby Universe, and binary evolution models for the formation of such sources, in ways not possible with electromagnetic observations alone. In this work we begin laying the ground-work for the astrophysical interpretation of future GW observations of BCO inspiral by focusing attention on the selection effects associated with GW observations: in particular, effects associated with binary component masses, the GW antenna beam and the local geography of the Universe.  Over the past decade there have been many predictions of the detection rate in LIGO of BCO inspirals \\citep[e.g., ][]{belc02,kal03}.  These calculations all begin by estimating the Galactic coalescence rates and extrapolating it to other galaxies. The observed rate is then calculated assuming that galaxies (and, thus, binaries) are distributed homogeneously and isotropically in the local universe, and that the LIGO detector observes all coalescing binaries inside a fixed distance, which is the radius of a sphere that would have the same volume as is effectively surveyed by the detector. These approximations are inadequate when we wish to go beyond order-of-magnitude predictions and actually interpret observed events as constraints on the actual compact object binary population, as is our goal.  To improve on past models for the physical population of inspiraling compact binary systems, we use galaxy catalogs to model the actual distribution of galaxies in the local universe and we use stellar synthesis calculations \\cite[specifically those of ][]{belc02} to model the mass distribution of binaries within each galaxy. From the constructed population models, we determine the compact binary coalescence rate and distribution with binary system mass that we expect the LIGO detector system to observe, taking full account of each galaxy's distance and declination, the LIGO detector system's noise spectrum, and its position and orientation on Earth.  Our principal goal in relating a physical population model to the distribution that we expect modern GW detectors to observe is to enable observations by those detectors to constrain the population model (see \\citealt{bulik} and \\citealt{bulik2} for recent studies with similar goals). Through the calculations described here, comparison of future observed rates or rate upper limits constrain stellar synthesis models and the overall binary compact object population.  While our principal interest is in preparing for this kind of interpretation of forthcoming observations, as a by-product of our investigations we have improved detection rate predictions as well.  In \\S\\ref{sec:background} we present an overview of the various approaches used so far for the extrapolation of Galactic detection rates to extragalactic distances and introduce our novel galaxy-by-galaxy approach, whereby we calculate the detectability of BCO inspiral for each galaxy in our catalog.  In \\S\\ref{sec:method} we describe how we calculate, from the detailed extra-galactic population model described in \\S\\ref{sec:background}, the observed distribution of BCOs.  In \\S\\ref{sec:results} we discuss our results, including the LIGO detector system's efficiency for detecting binaries from different galaxies in the nearby universe, the expected observed coalescing binary mass distribution, new detection rate predictions, and the implications of the geography of the nearby univese for detection of binary compact object systems. We end in \\S\\ref{sec:conclusions} with a summary of our main conclusions. ",
        "conclusions": "\\label{sec:conclusions}   In anticipation of the development of GW astrophysics in the next several years, we consider the effects of observational selection effects on the detectability of BCO inspiral events. Our primary goal is to develop a realistic framework for the astrophysical interpretation of rate constraints (from upper limits or inspiral detections) anticipated in the next few years. This interpretation should account for the main selection effects associated with ground-based GW observations and properly constrain models of radio pulsar and BCO populations. As a result of our calculations we also make realistic estimates for the extrapolation of Galactic inspiral detection rates based on the known spatial distribution of galaxies in the nearby universe and the expected mass distributions of binary compact objects.  Our results are summarized as follows: \\begin{itemize}  \\item The local distribution of galaxies mostly relevant to NS-NS detections with initial LIGO is in fact very different from isotropic in sky direction and volume density. Most importantly the Virgo cluster represents a significant step in the cumulative blue luminosity (or the cumulative number of MWEG) all concentrated at a given (rather unfortunate) sky position. Failure to properly account for this local distribution of BCO sources would lead us to underestimating the importance of an upper limit on the inspiral rate derived from GW observations.  \\item Until this study, because of the assumption of isotropic distribution of galaxies, detection rates of NS-NS inspirals have been {\\em under}estimated by factors of $2-4$ and BH-BH inspirals have been {\\it over}estimated by nearly a factor of $2$. These factors include the systematic uncertainties due to the chirp-mass distributions that are not very well constrained.  \\item Detections of inspiral events and measurements of compact object masses are expected to provide us with tighter constraints on the BCO mass distributions and thus on the physics of BCO formation. However, our analysis shows that mass distributions of detected BCOs are strongly skewed towards higher masses (because of their stronger signals) compared to the parent mass distribution (see Figures \\ref{fig:chirpdist} and \\ref{fig:chirp}). For our reference population model, we find that about half of detected inspirals correspond to binaries with high chirp masses ($5-9$\\,M$_{\\odot}$). Since event rate limits in this range will be most constraining to BCO models, it is evident that there is a need for the development of efficient search methods for such massive systems. Understanding the systematics of this bias will be crucial for the astrophysical interpretation of such detections.  \\item Inspiral detection efficiency depends strongly on the host galaxy sky position and the binary orbit orientation with respect to the detectors; as a result the true maximum distance for an optimally oriented binary can exceed the average detection distance by more than a factor of two.  \\item Using the current most favorable (at peak probability) estimates of NS-NS inspiral rates for the Galaxy \\citep{kal03} and our results on $N_{G}$ values, we find expected initial LIGO detection rates in the range of one event per 200 -- 3 years (for the reference pulsar population model at 95\\% confidence level the range is one event per 3 - 50 years).   \\end{itemize}  From the various galaxy physical properties we have considered the blue luminosity (corrected for reddening), but we have ignored galaxy metallicity and star formation history. Both of these factors affect the expected mass distribution of compact object binaries as well as their birth rate, for a given luminosity. For example, metallicity affects massive stellar winds and the final compact object masses. This effect has already been taken into account in \\cite{inspiral} where the Magellanic clouds have been reached by LIGO. In principle we would like to include these effects in our calculations (to the extent of our current understanding of binary evolution and how it is affected by these factors); however at present it does not seem possible since this information is not available in detail for every galaxy in the catalogs, and we chose to ignore these factors instead of include them for only a very small subset of sample.  With the work presented here we also advance a paradigm for using initial LIGO binary inspiral to constrain models of binary evolution and BCO formation and of pulsar population properties. Using the current estimates of NS-NS inspiral rates in MWEG \\citep{kal03} and scaling to the BH-BH population, it is clear that LIGO should eventually provide an astrophysically significant bound on the rate of BH-BH inspirals in the nearby universe. In the context of a particular binary synthesis model, such a bound can be translated to a bound on the MWEG BH-BH inspiral rate as well as on the rate fir the NS-NS and NS-BH sub-populations. The derived bound on the NS-NS sub-population can be compared to the estimates that arise from binary pulsar observations \\citep{kal03}. Note that GW observations may also directly bound the coalescence rate for the NS-BH and NS-NS sub-populations at a significant level. All these bounds will be consistent only for certain binary formation and synthesis models. In this way, GW observations will contribute to our understanding of compact binary formation and evolution."
    },
    "0402/astro-ph0402328_arXiv.txt": {
        "abstract": "{\\bf \\noindent Recent observations $^{1-6}$ have revealed an unexpectedly high binary fraction among the Trans-Neptunian Objects (TNOs) that populate the Kuiper-belt.  The discovered binaries have four characteristics they comprise a few percent of the TNOs, the mass ratio of their components is close to unity, their internal orbits are highly eccentric, and the orbits are more than 100 times wider than the primary's radius.  In contrast, theories of binary asteroid formation tend to produce close, circular binaries.  Therefore, a new approach is required to explain the unique characteristics of the TNO binaries. Two models have been proposed\\cite{Weidenschilling2002,Goldreich2002}. Both, however, require extreme assumptions on the size distribution of TNOs.  Here we show a mechanism which is guaranteed to produces binaries of the required type during the early TNO growth phase, based on only one plausible assumption, namely that initially TNOs were formed through gravitational instabilities\\cite{GoldreichWard1973} of the protoplanetary dust layer.} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0402/astro-ph0402602_arXiv.txt": {
        "abstract": "We study the statistical properties of the 1st year WMAP data on different scales using the spherical mexican hat wavelet transform. Consistent with the results of Vielva et al. (2003) we find a deviation from Gaussianity in the form of kurtosis of wavelet coefficients on $3-4^\\circ$ scales in the southern Galactic hemisphere. This paper extends the work of Vielva et al. as follows. We find that the non-Gaussian signal shows up more strongly in the form of a larger than expected number of cold pixels and also in the form of scale-scale correlations amongst wavelet coefficients. We establish the robustness of the non-Gaussian signal under more wide-ranging assumptions regarding the Galactic mask applied to the data and the noise statistics. This signal is unlikely to be due to the usual quadratic term parametrized by the non-linearity parameter $f_{NL}$. We use the skewness of the spherical mexican hat wavelet coefficients to constrain $f_{NL}$ with the 1st year WMAP data. Our results constrain $f_{NL}$ to be $50\\pm 80$ at 68\\% confidence, and less than 280 at 99\\% confidence. ",
        "introduction": "The current cosmological model assumes Gaussian initial conditions, created by inflation. This assumption regarding the nature of primordial density perturbations can be verified by studying the distribution of temperature fluctuations in the cosmic microwave background (CMB). While the simplest inflationary models predict Gaussian primordial perturbations, there are other models of inflation, such as those involving multiple scalar fields, features in the inflaton potential or phase transitions, that could give rise to detectable non-Gaussianity. Hence studies of Gaussianity help distinguish between different early universe scenarios. Gaussianity is also a key underlying assumption of CMB data analysis wherein the angular power spectrum fully specifies its statistical properties, and must be tested. Non-Gaussianity can also be associated with secondary anisotropies in the CMB, or with foreground contamination and systematic effects.  Prior to the release of WMAP data there was no clear evidence of cosmological non-Gaussianity. Since the release of 1st year WMAP data, a number of tests of non-Gaussianity have been performed, with somewhat differing results. Each statistic is sensitive to a different kind of non-Gaussianity, hence there is need for a wide variety of tests. Komatsu et al. (2003) use an optimized test based on the bispectrum, as well as Minkowski functionals, while Colley \\& Gott (2003) study the genus, and both groups report consistency with Gaussianity. Gaztanaga \\& Wagg (2003) do a 3-pt angular correlation function analysis and find consistency with Gaussianity as well. Chiang et al. (2003) perform a study of the phases of spherical harmonics and find some evidence for non-Gaussianity at high multipoles. Copi et al. (2003) find some evidence for low $l$ correlations and deviation from isotropy. Park (2003) find a large difference between the genus amplitudes of the northern and southern hemispheres and a positive genus asymmetry in the southern hemisphere. Eriksen et al. (2003a) compute the 2 and 3-pt correlations and report a significant north-south asymmetry; Eriksen et al. (2003b) use Minkowski functionals and find a significant genus in the northern hemisphere and again indications of north-south asymmetry. Hansen et al. (2004) use local curvature and find non-Gaussianity/asymmetry in the data on scales of a few degrees. Gurzadyan et al. (2004) find ellipticity in the temperature anisotropy features in the data, consistent with what was found previously in BOOMERang data. Vielva et al. (2003; hereafter V03) report a non-Gaussian signal in the southern hemisphere at high significance in the form of kurtosis on $\\sim 4^\\circ$ scales using the spherical mexican hat wavelet transform on WMAP data. Some of the detections of non-Gaussianity and/or asymmetry thus far reported in the WMAP data are at the level of 99\\% or greater.  Wavelet transforms are useful tools in non-Gaussianity studies because they enable the signal on the sky to be studied on different scales, with simultaneous position localization, so that the obscuring effects of the central limit theorem, that can exist in both real and Fourier spaces, are reduced. With wavelets any non-Gaussian detection can be localized on the sky in scale and position, so that its nature and source can be better determined. Planar wavelets have been used in Gaussianity studies of the CMB by Pando et al. (1998), Hobson et al. (1999) and Mukherjee et al. (2000), while Barreiro et al. (2000) use the spherical Haar wavelet, and Cay\\'on et al. (2001,2003), Mart\\'inez-Gonz\\'alez et al. (2002), and V03 use the spherical mexican hat wavelet (SMHW). Wavelet methods have been compared with other pixel or Fourier based methods in Hobson et al. (1999), Aghanim et al. (2003), Cabella et al. (2004), and the performance of isotropic as well as highly anisotropic multi-scale bases in distinguishing between different sources of non-Gaussianity in the CMB has been studied in Starck et al. (2003).  In this paper we use the spherical mexican hat wavelet transform to probe non-Gaussianity in the WMAP data. We extend the work of V03 by performing new multiple tests of the robustness of the non-Gaussian signal. We also look at the non-Gaussianity in terms of an excess in the number of cold pixels, and in terms of scale-scale correlations amongst wavelet coefficients. Further, we place constraints on a popular form of non-Gaussianity (a quadratic term in the curvature perturbations parametrized by the non-linearity parameter $f_{NL}$). This paper is organized as follows. In \\S 2, we present results from using the SMHW transform on WMAP data. Whilst confirming the results of V03, we perform new multiple tests of the robustness of the non-Gaussianity signal in the kurtosis spectrum (a) through the use of different (extended) masks, and (b) relaxing the assumption of a simplified noise model. We find that the signal shows up even more significantly in the form of the number of cold pixels (or coefficients). In \\S 3, we examine scale-scale correlations amongst the wavelet coefficients. We find significant deviations from Gaussianity, a corroboration of the signal detected and described in \\S 2. In \\S 4 we obtain constraints on the non-linearity parameter $f_{NL}$. We conclude in \\S 5. ",
        "conclusions": "We have analysed the first year WMAP data using a spherical mexican hat wavelet. We detect non-Gaussianity at $\\sim$ 99\\% significance, consistent with that reported by Vielva et al. (2003). This detection corresponds to a positive kurtosis and to the presence of a larger than expected number of cold pixels (wavelet coefficients) in the southern Galactic hemisphere on scales $3-5^\\circ$.  We have tested for changes in the significance of the signal with the type of mask used. The signal is found to be robust, and is found in the ILC map as well. We have also compared confidence contours obtained for the kurtosis spectra using the full noise simulation maps provided by the WMAP team, containing $1/f$ noise and other effects from data processing, to those obtained from using simulations that contain just white noise. We find very good agreement.  We have also applied another test statistic, the scale-scale correlations between wavelet coefficients. Significant scale-scale correlations are seen amongst the coefficients over the range of scales that indicate the above non-Gaussianity.  We then use the skewness statistic on the different scales to place constraints on the non-linear coupling parameter $f_{NL}$, the motivation being to see how much non-Gaussianity of this particular form is allowed by current data. It is also a way to compare the sensitivity of different test statistics to this parameter. Constraints obtained are closely consistent with those obtained by Komatsu et al. (2003) using the cubic statistic and Minkowski functionals on the same data. The constraints on $f_{NL}$ derived here could possibly be made more stringent if we used spherical wavelets on the Wiener filtered map of primordial perturbations as discussed in Komatsu, Spergel \\& Wandelt (2003). We will explore this in a future paper. The kurtosis statistic is not sensitive to this form of non-Gaussianity. We will present constraints on other forms of non-Gaussianity implied by the kurtosis statistic of the WMAP data elsewhere."
    },
    "0402/astro-ph0402434_arXiv.txt": {
        "abstract": "We examine the effect of self-gravity in a rotating thick-disk equilibrium in the presence of a dipolar magnetic field. In the first part, we find a self-similar solution for non-self-gravitating disks. The solution that we have found shows that the pressure and density equilibrium profiles are strongly modified by a self-consistent toroidal magnetic field. We introduce 3 dimensionless variables $C_B$, $C_c$, $C_t$ that indicate the relative importance of toroidal component of magnetic field ($C_B$), centrifugal ($C_c$) and thermal ($C_t$) energy with respect to the gravitational potential energy of the central object. We study the effect of each of them on the structure of the disk. In the second part, we investigate the effect of self-gravity on the these disks; thus we introduce another dimensionless variable ($C_g$) that shows the importance of self-gravity. We find a self-similar solution for the equations of the system. Our solution shows that the structure of the disk is modified by the self-gravitation of the disk, the magnetic field of the central object, and the azimuthal velocity of the gas disk. We find that self-gravity and magnetism from the central object can change the thickness and the shape of the disk. We show that as the effect of self-gravity increases the disk becomes thinner. We also show that for different values of the star's magnetic field and of the disk's azimuthal velocity, the disk's shape and its density and pressure profiles are strongly modified. ",
        "introduction": "The theory of accretion disks, motivated in a large degree by their occurrence in some binary systems, particularly cataclysmic variables, has been most fully developed for the thin Keplerian disks \\cite {pringle1}. Based on their geometric shapes, accretion disks are generally divided into two distinct classes, thin disks and thick disks. The theory of thin accretion disks (Shakura \\& Sanyev 1973) is well understood whereas there is no universally accepted model for thick accretion disks. The current interest in this theory is due to the possibility that thick disks may be relevant to the understanding of central power sources in radio galaxies and quasars. Observational evidence suggests that in the center of many galaxies, matter is somehow ejected to large distances and gives off high energy radiation as it interacts with the external medium. There are other theoretical reasons for pursuing the study of thick disks. In the theory of thin disks, radial pressure gradients are neglected and vertical pressure balance is solved separately.It was shown that this approximation is valid as long as the disk is geometrically thin. This condition may be violated in the innermost region of accretion disks around stellar black holes and neutron stars. The study of thick disks provides a better theoretical understanding of thin disks as a limiting case and enables us to deal with intermediate situations \\cite{Frank}  When considering the formation process of astrophysical objects, such as galaxies and stars, the most crucial factor is self-gravity. In the standard thin accretion disk model, the effect of self-gravity is neglected, and only pressure supports the vertical structure. By contrast, the theory of self-gravitating accretion disks, is less developed. Early numerical work of self gravitating accretion disks began with N-body modelling (Cassen \\& Moosman 1991 ; Tomley, Cassen \\& Steinman-Cameron 1991). The time evolution of non-self-gravitating viscose disk has long been studied, and now we have a good theory describing its steady structure and its basic time-dependent behavior \\cite {shakura}; \\cite {pringle}. But with the added assumption of self-gravity in the disk, it is not easy to follow its dynamical evolution, mainly because the basic equations for the disks are highly nonlinear \\cite {paz}; \\cite {fuk}. To solve the nonlinear equations of self-gravitating disks, the technique of self-similar analysis is sometimes useful. Several classes of self-similar solutions were known previously, but all of them considered a disk in a fixed, external potential. Self-similar behavior provides an important class of solutions to the self-gravitating fluid equations. On the one hand, many physical problems often attained self-similar limits for a wide range of initial conditions. On the other hand, the self-similar properties allows us to investigate properties of solutions in arbitrary details, without any of the associated difficulties of numerical hydrodynamics.  Pen (1994) presented a general classification of self-similar self-gravitating fluids. Fukue \\& Sakamoto (1992) also analyzed the vertical structure of self-gravitating disks, but it is impossible to compare their models with realistic disks, because they computed the vertical structure using the thin-disk approximations for polytropes. Finally, using the numerically method proposed by Hachisu (1986) and Hachisu et al. (1987), Woodward et al. (1992) developed a code to investigate the interaction between a disk and its central object. However, they only considered polytropic disks. Hashimoto, Eriguchi \\& Muller (1995) presented a two-dimensional equilibrium model for self-gravitating Keplerian disks. They showed that the shape of the disk (or disk thickness) in flounced by the rotational law and the ratio of the disk mass to the mass of the central star. Bodo \\& Curir (1992)computed the equilibrium structure of a self-gravitating thick accretion disk by an iterative procedure which produced a final density distribution in equilibrium with the potential coming from it. They showed that the geometrical size and shape the disks influenced by self-gravity of the disk.  Accretion disks, containing magnetic fields, have been the subject of intense study in recent years. The role of the magnetic field in the equilibrium of accretion disks has been investigated by some authors \\cite{blandford2}; \\cite{pringle2} for the thin disk models , an ideal magneto hydrodynamics (MHD) equilibrium with azimuthal velocity and poloidal magnetic field has been analyzed \\cite {Lovelace}; \\cite {mobarry}. Using numerical methods, these authors found that the magnetic field may change the shape and angular momentum distribution in the disk. Thick disk configurations with a poloidal magnetic field has been studied by \\cite {tri}, in the MHD framework \\cite {das}. They investigated the equilibrium structure of thick disks and their stability in the presence of a dipolar magnetic field due to a non-rotating central object. Their solution shows that the pressure and the density equilibrium profiles are strongly modified by a toroidal magnetic field, resulting from the interaction between the permanent dipolar magnetic field and the inertia of the gas disk. In a magnetized disk, the inertia of the gas is expected to bend the magnetic field lines backwards, creating a toroidal component, which in turn may collimate a hydrodynamic outflow over long distances, forming jets. They assumed that the disk is non-accreting, stationary, axisymmetric, non-viscous, magnetized and that it is in equilibrium around a compact object, with only an azimuthal motion $V_{\\phi}$.  We are interested in analyzing the role of self-gravity in thick disk equilibrium in the presence of the dipolar magnetic field of a central star. The outline of this paper is as follows: the general formalism of the problem is discussed in section 2, a self-similar solution of the equilibrium of non-self-gravitating accretion disks in the presence of a dipolar magnetic field is constructed in section 3, a self-similar solution of self-gravitating magnetized accretion disks is constructed in section 4 and a summary of the main ideas is given in section5. ",
        "conclusions": ""
    },
    "0402/astro-ph0402616_arXiv.txt": {
        "abstract": "% The advent of wide-area multicolour synoptic sky surveys is leading to data sets unprecedented in size, complexity and data throughput. VO technology offers a way to exploit these to the full but requires changes in design philosophy. The Palomar-QUEST survey is a major new survey being undertaken by Caltech, Yale, JPL and Indiana University to repeatedly observe $\\frac{1}{3}$ of the sky ($\\sim 15000$ sq. deg. between $-27^\\circ \\le \\delta \\le 27^\\circ$) in seven passbands. Utilising the 48-inch Oschin Schmidt Telescope at the Palomar Observatory with the 112-CCD QUEST camera covering the full 4$^\\circ$ x 4$^\\circ$ field of view, it will generate $\\sim 1$TB of data per month. In this paper, we review the design of QUEST as a VO resource, a federated data set and an exemplar of VO standards. ",
        "introduction": "The new availability of wide-field images from Schmidt telescopes in the 1940's meant that astronomers no longer had to make educated guesses about where to look to find new and interesting phenomena but were now spoilt for choice. The advent of synoptic surveys presents more extreme opportunities; as an illustration, consider the SDSS which over the course of 5 years represents a factor of a million increase in information over previous surveys; however, the \\htmladdnormallinkfoot{LSST}{http://www.lssto.org} (Large Sky Synoptic Telescope, Tyson (2002)) will amass a SDSS every 3 nights.  Although overviews of synoptic surveys are riddled with cliches concerning undiscovered countries and uncharted waters, the exploration of the temporal domain results in data sets that are not just more voluminous than before, but far richer and more complex (Paczynski 2001; Djorgovski et al. 2000). This presents challenges to all aspects of astronomy: data gathering, distribution, reduction, analysis, storage, archiving, dissemination and mining. VO technologies are being designed precisely to meet these types of challenges, but to use them requires changes in survey design philosophies. ",
        "conclusions": ""
    },
    "0402/astro-ph0402085_arXiv.txt": {
        "abstract": "We present Keck/NIRSPEC near-IR images and Magellan/IMACS optical spectroscopy of the host galaxy of GRB~031203.  The host is an actively star-forming galaxy at $z=0.1055 \\pm 0.0001$. This is the lowest redshift GRB to-date, aside from GRB~980425. From the hydrogen Balmer lines, we infer an extinction of $A_V = 3.62 \\pm 0.25$ or a total reddening $E_T(B-V) = 1.17 \\pm 0.1$ toward the sightline to the nebular regions. After correcting for reddening, we perform an emission-line analysis and derive an ISM temperature of $T=13400 \\pm 2000$\\,K and electron density of $n_e = 300 \\cm{-3}$. These imply a metallicity $\\lbrack$O/H$\\rbrack$~$= -0.72 \\pm 0.15$\\,dex and a roughly solar abundance pattern for N, Ne, S, and Ar. Integrating \\halph, we infer a dust-corrected star formation rate (SFR) of $> 11 {\\rm M_\\odot \\, yr^{-1}}$. These observations have the following implications: (1) the galaxy has a low $K'$-band luminosity $L \\approx L_K^*/5$, typical of GRB host galaxies; (2) the low redshift indicates GRB~031203 had an isotropic-equivalent $\\gamma$-ray energy release smaller than all previous confirmed GRB events.  The burst discovery, near the detection limit of INTEGRAL, raises the likelihood of identifying many additional low $z$, low flux events with {\\it Swift}; (3) the large SFR, low metallicity, and the inferred hard radiation field is suggestive of massive star formation, supporting the collapsar model; (4) several lines of evidence argue against the identification of GRB~031203 as an X-ray flash event. ",
        "introduction": "\\label{sec-intro}  The study of the host galaxies of gamma-ray bursts (GRBs) plays a central role for studies of progenitor theory and afterglow observations.  First, host redshifts are required to derive the burst energy and afterglow timescale and to determine the GRB rate density evolution (e.g. \\citealt{djk03}).  Second, the host photometric properties \\citep{sokolov01,cba02} together with burst locations \\citep{bloom02} within hosts support the notion of a progenitor population intimately connected with star formation (e.g. \\citealt{pac98}). It is now widely accepted that long-duration GRBs originate from the deaths of massive stars, as as suggested by \\cite{woosley93} and confirmed by recent observations \\citep[e.g.][]{bloom99,stanek03,hjorth03}.  Third, limits on time-variable hydrogen- and metal-absorption in the hosts provide constraints on the physical state and chemical composition of the interstellar medium (ISM) in the vicinity of the GRBs \\citep{perna98,draine02,mirabal03}.   At the same time, long-duration GRBs present an alternative means to address a number of open questions in different sub-fields of extragalactic research related to the formation of massive stars.  For example, early-time spectroscopy of GRB optical afterglows allows us to measure the metallicity and dust content of the progenitor environment through absorption line studies \\citep[e.g.][]{vreeswijk01}. Late-time galaxy spectroscopy of the GRB hosts allows us to study the ionization state and metallicity of the general ISM properties of host galaxies using various spectral diagnostics such as the [Ne\\,III]/[O\\,II] ratio \\citep{bloom01}. Together these results could impose strong constraints on the stellar initial mass function (IMF) and chemical enrichment history of the galaxy population that hosts GRBs.  Previous work presents some evidence that GRBs arise preferentially in low-metallicity environments with a top-heavy IMF \\citep{savaglio03}, which is supported by the agreement between the luminosity distribution of GRB hosts and faint blue galaxies found in the field.  In addition, galaxies selected by association with GRBs are less affected by dust than those found in optical or sub-millimeter surveys \\citep{berger03} and therefore may be adopted to constrain the amount of obscured star-formation in the universe \\citep{blain00, djk03}.  Finally, afterglow studies of GRBs discovered in the early universe ($z \\gtrsim 8$) will undoubtedly advance our understanding of the first generations of stars as well as chemical evolution in the early universe.   In this paper we report the spectroscopic identification of the host of GRB~031203 at $z=0.1055$ and present a case study of imaging and spectral properties of the host population.  The long-duration (20\\,sec) GRB~031203 triggered the IBIS instrument of INTEGRAL on 3 December 2003 at 22:01:28 UT with an initial localization of 2.5\\arcmin\\  \\citep[radius;][]{Gotz03}. A 56k sec observation with XMM-Newton beginning at 03:52 UT on December 4, 2003 led to the discovery of an X-ray source (hereafter S1), not present in the ROSAT point source catalog \\citep{Campana03}, near the center of the INTEGRAL error circle \\citep[$\\alpha=08^h 02^m 30^s.19$, $\\delta=-39^\\circ 51' 04''.0$, J2000;][]{Santos-Lleo03}.  Source S1 was later reported to have faded over the first XMM pointing \\citep{Rodriguez-Pascual03} and the refined position of S1 \\citep{vaughan04} is consistent with a fading radio source \\citep{Frail03,Soderberg03}. Moreover, the expanding X-ray dust echo around the afterglow is consistent with a bright explosive event (in X-rays) at the position of the galaxy and coincident in time with GRB~031203 \\citep{vaughan04}. The simplest assumption is that both the X-ray and radio transients are afterglow emission from GRB~031203.  Before the radio transient was found, several attempts were made to identify the optical afterglow \\citep[see][]{cobb04}. The identification of the optical transient was first reported by \\citet{Hsia03} to be within the error circle of S1, but this identification was invalidated later based on the detection of sources in the J- and F-plates of the DSS \\citep{Bloom03}.  Since the radio transient was later found to be coincident with this source and because the source appeared extended in an $I$-band image, \\citet{Bloom03a} suggested that the source was a galaxy---either in the foreground or the host of GRB~031203. Subsequent radio astrometry improved the coincidence of the radio and optical source \\citep{Soderberg03}, confirming that this galaxy was indeed the host of the GRB \\citep{pro03a}.  A detailed overview of the optical astrometry of the host is presented in \\cite{cobb04}.  We acquired optical spectroscopy of the galaxy on the Magellan telescopes and reported a preliminary redshift of $z= 0.105$ \\citep{pro03b}.  Although the host galaxy lies at low Galactic latitude and is subject to significant extinction, its proximity affords a careful examination of its properties.  We present detailed IR imaging and optical spectroscopy of the galaxy and analyze these observations to determine its luminosity, metallicity, and star formation rate.  This paper is organized as follows: $\\S$~2 reviews the observations and data analysis; we perform an emission line analysis in $\\S$~3 to determine the reddening, metallicity, and relative abundances of the nebular region; and $\\S$~4 discusses the star formation rate and implications of the galaxy properties for the GRB phenomenon.  Throughout this paper we will adopt a standard $\\Lambda$~cosmology ($H_0 = 70 \\mkms$ Mpc$^{-1}$, $\\Omega_\\Lambda = 0.7$ and $\\Omega_m = 0.3$).   \\begin{figure}[ht] \\begin{center} \\includegraphics[width=3.6in]{f1.eps} \\figcaption{SCAM images obtained with the NIRSPEC instrument on the Keck~II telescope on 2003 December 05.6. The arrow designates the host galaxy studied in this paper. The field of view is $\\approx 46''\\times 46''$ with orientation N up and E left and the pixel size is 0.183$''$ on a side. } \\label{fig:img} \\end{center} \\end{figure}  \\begin{figure*} \\begin{center} \\includegraphics[width=6.8in]{f2.eps} \\figcaption{Magellan/IMACS spectrum of the host galaxy of GRB~031203. The spectrum has a FWHM~$\\approx 5$\\AA\\ resolution.  The lower two panels have been smoothed by 3~pix ($\\approx$ 1/2 resolution element) for presentation purposes only. The strong emission line features are centered at $z=0.1055 \\pm 0.0001$ and provide measurements of the extinction and the nebular temperature, density, metallicity and SFR.  Note the detections of [NeIII] $\\lambda 3868$ and [OIII] $\\lambda 4363$ as well as the low [OIII]/\\hbeta\\ ratio. Together these observations imply a relatively hard radiation field dominated by massive stars and a metal-poor gas.  The strong \\halph\\ emission corresponds to a SFR of $> 10 \\msol {\\rm yr^{-1}}$. } \\label{fig:spec} \\end{center} \\end{figure*} ",
        "conclusions": "\\subsection{Luminosity, Metallicity, and Star Formation Rate}  Even with nearly one magnitude of foreground extinction at 1$\\mu$m, \\grbnm\\ offers a rare opportunity to study the large-scale environment of a nearby GRB event. Our analysis of the nebular region of \\grbnm\\ reveals a metal-poor, star-forming galaxy. We can combine our extinction analysis with the apparent near-IR magnitudes and the redshift of \\grbnm\\ to assess the luminosity of the old stellar population within the galaxy. Adopting a Galactic reddening of $E_G(B-V) =  0.78$ (see above), we derive extinction corrections of 0.67, 0.45, and 0.28 magnitudes for the $J, H$ and $K'$ bands respectively. Adopting the standard $\\Lambda$~cosmology, $z_{gal} = 0.1055$ and a k-correction of $-0.2$~mag, we calculate absolute magnitudes of $-21.20$, $-21.47$, and $-22.35$~mag in $JHK'$ respectively. The resulting $K'$ luminosity is +1.9~mag fainter than $M(K)_*$ at $z=0$ \\citep{cole01}, i.e., the galaxy has a low luminosity $(L \\approx L_K^*/5)$ and presumably a low total mass. In these respects, \\grbnm\\ is a relatively ordinary, faint galaxy. The $K'$-band luminosity is consistent with the median absolute K magnitude compiled by \\cite{lefoch03} for GRB host galaxies.  Apparently, these galaxies are, as a class, underluminous in terms of their old stellar population.  The galaxy is notable, however, for a large star formation rate (SFR). With our spectroscopic observations, we infer the SFR from the integrated, extinction-corrected \\halph\\ luminosity (Table~\\ref{tab:emlin}). Adopting the \\halph\\ relation given by \\cite{kennicutt98},we find SFR(\\halph)~$= 11 \\, {\\rm M_\\odot yr^{-1}}$. We believe the uncertainty in SFR(\\halph) is dominated by uncertainties in the SFR calibration which \\cite{kennicutt98} estimates to be $\\approx 30\\%$. We also attempt to infer the SFR from the dust-corrected [OII] luminosity using the comprehensive prescription described by \\cite{kewley04}. Unfortunately, their analysis is only applicable to emission-line regions with log(O/H)~+12~$> 8.2$ (Kewley, priv.\\ comm.);  the SFR corrections diverge at the nominal metallicity of \\grbnm\\ (see their Figure~9)\\footnote{These points not withstanding, we caution the reader that SFR values, irrespective of reddening corrections, derived for other GRB hosts based on the [OII] lines alone may have uncertainties of $>100\\%$, especially if the metallicity is poorly constrained.}. In addition, the [OII] lines have greater sensitivity to the dust corrections and much poorer SNR than the \\halph\\ emission. Therefore, we consider only the SFR value implied by \\halph. We stress that this SFR value should be considered a lower limit to the total SFR because (i) the galaxy is larger than the slit used and (ii) the galaxy may contain regions of star formation which are enshrouded in dust. The latter effect is presumably small for this relatively metal-poor galaxy while the former effect may increase the SFR by up to 50$\\%$. Indeed it will be of interest to measure the late-time radio and sub-mm flux from this galaxy to assess the level of obscured star formation and compare long-wavelength emission of this host with GRB hosts at higher redshifts \\citep[following][]{berger03}.   To qualitatively assess the nature of star formation for \\grbnm, we would like to contrast its characteristics with local samples. One simple comparison is the rest-frame equivalent width of the \\halph\\ line.  For \\grbnm, we measure EW$_{\\rm rest}$(\\halph)~$= 550.7  \\pm 2.3$\\AA\\ which is significantly larger than normal, star-forming galaxies \\citep{jansen00,nakamura04}. This large value suggests a system undergoing a short, very intense burst of star formation.  Another valuable diagnostic is the ratio of SFR to total luminosity. Unfortunately, there is no large, single survey to date which has compared SFR and near-IR luminosity.  We therefore estimate the B-band luminosity from the measured continuum flux at $\\lambda_{obs} \\approx 4500$\\AA.  We observe $F_{4500} = 8.0 \\sci{-18}$~ergs/s/cm$^{-2}$/\\AA\\ which translates to a Vega magnitude $B=18.1$\\,mag, consistent with the detection of the galaxy in the DSS-J plate \\citep{Bloom03}. Adopting $E_G(B-V)=0.78$, we impose a dust correction of 3.1~mag and derive an absolute magnitude $M_B = -19.3$. We stress that this luminosity is likely an underestimate of the total B-band light because the galaxy overfilled the 0.75$''$ slit. Furthermore, this $E_G(B-V)$ value may be an underestimate (see the discussion in $\\S$ref{sec-emiss}). We compare the measured \\halph/B-band luminosity of \\grbnm\\ against the KPNO International Spectroscopic Survey, an emission-line survey of galaxies with $z<0.095$ selected in low-dispersion objective-prism spectra \\citep{salzer00}. Restricting our comparison to those galaxies with accurate measurements of $L({\\rm H\\alpha})$ and $M_B$ \\citep{gronwall04}, we note \\grbnm\\ falls at the upper end of the distribution, i.e. its SFR is $\\approx 5 \\times$ higher than galaxies with similar B-band luminosity. A portion of this offset could be explained by correspondingly higher slit losses for the B-band light than \\halph\\ and/or a higher $E_G(B-V)$ value. These corrections not-withstanding, we suspect the galaxy has a higher than average SFR per unit B-band luminosity than other local galaxies.   We can also place \\grbnm\\ on the metallicity/luminosity locus of KISS galaxies \\citep{melbourne02}.  \\grbnm\\ lies below the entire distribution.  Furthermore, it falls $\\Delta{\\rm (O/H)} \\approx -0.9$~dex or $\\Delta M_B \\approx -2$~mag off their fit to the KISS sample even if we adopt $M_B = -19.3$\\,mag. We speculate this offset is characteristic of a very young star forming region.  Perhaps we have observed the galaxy prior to the production and/or distribution of significant metals into the nebular regions.  Indeed, a similar effect is observed for the star-bursting Lyman break galaxies at high redshift \\citep{pettini01,shapley04}.  \\subsection{The case for identifying \\grbnm\\ as the host galaxy of GRB~031203}  The conclusions drawn in the following subsections hinge on the identification of \\grbnm\\ as the host galaxy of GRB~031203. In particular, this sets the redshift of GRB~031203 based on the observed emission lines of \\grbnm. Before proceeding, therefore, we review the evidence for this allegation.  First, \\cite{soderberg04} and \\cite{Rodriguez-Pascual03} respectively have localized fading radio and x-ray sources within the half-light radius of \\grbnm\\ \\citep[see][]{cobb04}. Integrating the K-band number density function of \\cite{chen04} to $K'=16.5$, the number density of galaxies is $7.7\\sci{-5} / \\square''$.  Therefore, the likelihood of laying within $2''$ of a galaxy at least as bright as \\grbnm\\ is $<0.1\\%$. Second, \\grbnm\\ has a $K'$-band luminosity typical of the GRB host galaxy distribution \\citep{lefoch03}.   Third, this galaxy exhibits a SFR $(>10 \\msol {\\rm yr^{-1}})$ which is $> 98\\%$ of all galaxies at $z \\lesssim 0.1$ \\citep{nakamura04}. This argument is independent of the apparent magnitude and, furthermore, we note that the SFR is comparable to that of other GRB hosts \\citep[e.g.][]{berger03}.  Fourth, the galaxy is metal-poor, a trait frequenly attributed to GRB host galaxies. Fourth, \\grbnm\\ exhibits a lower metallicity than $>99\\%$ of all galaxies with $M_B < -19.3$\\,mag \\citep{melbourne02,lamareille04}. Again, GRB hosts appear to be metal-poor in general. Finally, there is recent evidence that a supernova was associated with GRB~031203 with a location coincident with \\grbnm\\ \\citep{thomsen04,cobb04,galyam04}. Although none of these points can be considered a definitive argument for associating \\grbnm\\ with GRB~031203, together they present a very strong case for a physical connection between GRB~031203 and \\grbnm.  \\subsection{The Energetics of GRB~031203 and Arguments against an X-ray Flash} \\label{sec:ibis}   The fluence in the prompt gamma-ray emission in GRB~031203 has not yet been reported. To estimate this quantity, we fit the INTEGRAL SPI-ACS light curve\\footnotemark\\footnotetext{{\\tt http://isdc.unige.ch/cgi-bin/cgiwrap/$\\sim$beck/ibas/spiacs/ibas\\_acs\\_web.cgi}} with a double exponential (single pulse). Assuming Poisson weighting, scaling the reported peak flux to 1.3 $\\times 10^{-7}$ erg cm$^{-2}$ s$^{-1}$ (20--200 keV), and assuming that the IBIS {\\it count} rate is dominated by the photons in this low energy range, we estimate the fluence to be $(4 \\pm 1) \\times 10^{-7}$ erg cm$^{-1}$ (20--200 keV). This assumes a combined uncertainty of 25\\% in the peak flux and integrated (model) light curve. We can estimate the prompt isotropic-equivalent energy release $E_{iso}$(20--2000 keV) by assuming a $k$-correction to rest-frame (20--2000 keV) from an ensemble of bright bursts: $k = 2.6 \\pm 0.9$. We find $E_{\\rm iso}$(20--2000 keV) $= (2.6 \\pm 1.1) \\times 10^{49} h_{70}^{-2}$ erg with a plausible maximum of $\\approx 5 \\times 10^{49}$ erg if the burst was spectrally hard ($E_0 = 1000$ keV). This is nearly identical to an independent estimation by Watson et al., who found (under a differing set of assumptions) $E_{\\rm iso}$(20--2000 keV) $= 2.6 \\times 10^{49} h_{70}^{-2}$ erg (scaling to our chosen value of Hubble's constant).  As noted in Watson et al. (2004), this value of $E_{\\rm iso}$(20 -- 2000 keV) is about 30 times fainter than that inferred in the (geometry corrected) prompt emission of other cosmological GRBs \\citep{bfk03}. If this burst was collimated, then the true energy release in the $\\gamma$-ray bandpass was even lower. Soderberg et al. (2003) also pointed out that the kinetic energy in the blastwave, as proxied by the X-ray afterglow emission, was $10^{3}$ times lower than other GRBs at comparable times (Berger al. 2003). These two results seem to suggest that the total energy in the relativistic component was substantially lower than the other cosmological GRBs.  Watson et al.\\ (2004) have suggested that GRB~031203 may be an X-ray flash based upon analysis of the dust echo. However, since the occurrence of a bright soft X-ray component to the prompt burst has not been firmly established, we consider the XRF possibility less likely than Watson et al.\\ (2004). We note that the Watson et al.\\ argument rests on two points.  First, that the ratio of the 0.2-10 keV {\\it fluence} (estimated from the dust echo) to the 20-200 keV fluence (estimated from the {\\it peak flux} of the GRB) yields a reliable power law index.  And second, that the detection of this burst at energies $\\gtrsim 100$ keV by the INTEGRAL SPI-ACS is consistent with the spectral slope estimate.  We caution however that the first point mixes fluences and peak fluxes and cannot take any possible spectral evolution into account; until a reliable time-integrated spectrum above 20 keV is given, this procedure is subject to considerable uncertainty.  And contrary to claim of Watson et al.\\ we believe that the spectrum derived from this procedure cannot be said to be consistent with the SPI-ACS response.  Indeed, the ACS threshold is not well-defined, as it varies along the collimator, and blocking by INTEGRAL instruments for various angles complicates the response considerably.  There is at present no accurate calibration of the ACS.  Thus we believe that it is impossible to confirm or disprove the X-ray rich burst hypothesis based on either argument.   \\cite{vaughan04} have reported a dust scattered X-Ray halo from {\\it XMM-Newton} observations taken six hours after the burst. Comparing the spectral shape of the halo with the afterglow, they inferred a hydrogen column density $N_H = 8.8 \\pm 0.5 \\sci{21} \\cm{-2}$ consistent with the Parkes 21cm observations along this sightline \\citep{mcclure01}. \\cite{vaughan04} then estimated a time-integrated X-ray flux by assuming $A_V = 2$\\,mag, a value 4.4$\\times$ smaller than out total derived extinction. We caution, therefore, that they may have overestimated the X-ray flux in their analysis. Thus, we suggest GRB~031203 is most like GRB~980425, i.e.\\ a low-redshift underluminous burst associated with a supernova.  In this respect, it is interesting to note the single pulse nature of the burst \\citep{bloom98}, which may be related to the emission mechanism.   \\subsection{Implications for the GRB event}  Collapsars \\citep{woosley93}, the leading scenario for long-duration GRBs, requires a connection between the instantaneous ($\\lesssim 10^7$~yr) star-formation and the probability of bursting \\citep{bloom98,fryer99}. This model describes the GRB progenitor as a massive star ($> 30 \\msol$ at zero-age main sequence) which loses all of its hydrogen envelope and most of the helium. Importantly, \\cite{macfadyen99} noted that low metallicity favors a collapsar event because low metallicity decreases mass loss leading to a more massive, more rapidly rotating star, i.e.\\ characteristics necessary to make a disk and black hole. Examining \\grbnm\\, in this context, we note: (i) the presence of [NeIII] emission; (ii) O$^{++}$ is the dominant oxygen ion for the full range of extinction and temperatures considered; and (iii) the emission line regions have a significantly sub-solar metallicity. Points (i),(ii) indicate the presence of a relatively hard radiation field.  That is, massive stars $(> 30 \\msol$) contribute significantly to the spectral energy distribution of the host. Furthermore, point (iii) follows the assertion of \\cite{macfadyen99} that the collapsar event is more easily achieved in lower metallicity systems.  This point is accentuated by the fact that the system has a lower than average metallicity given its B-band luminosity. Together, these observations of \\grbnm\\ lend support to the collapsar model for GRB's.  Aside from the anomalous GRB~980425, the only other long duration GRB with a known low redshift $(z < 0.2)$ is GRB~030329 \\citep[see][for a review]{lipkin03}. The discovery of GRB~031203 at such a low redshift has several important implications.  First, it indicates the likelihood of many additional low redshift GRB events at low flux/fluence levels. Since so few low redshift bursts had been discovered in the past 6 years (1 out of $\\sim$100 well-localized GRBs) and the peak flux of GRB~030329 was more than two orders of magnitude over the detection threshold, on probabilistic grounds, \\citet{price03} suggested that bursts with redshifts as low as GRB 030329 would be rare even in the {\\it Swift} era (see also Schmidt 1999). The low redshift discovery of GRB\\,031203, however, requires a redress of this conclusion. Second, the peak flux of 1.3 $\\times$ 10$^{-7}$ erg cm$^{-2}$ s$^{-1}$ \\citep[20--200 keV;]{Mereghetti03} implies that GRB~031203 was one of the faintest rapid and well-localized bursts to date. While there is still some uncertainty in the $k$-corrections, it appears that the {\\it prompt} isotropic-equivalent energy of GRB~031203 appears to bridge the gap between the anomalous 980425 and the remainder of the bursts (see also Soderberg et al.\\ 2003). Again, we point out that 031203 is distinguished from other low-energy GRBs (030329, 980326, etc.; Bloom, Frail, \\& Kulkarni 2002\\nocite{bfk03}) in that the energy of 031203 in soft $\\gamma$-rays (20--2000 keV) appears low even before (the unknown) geometry correction. Given the uncertainties related to estimations of the energy release in prompt X-rays, the total energy budget of the burst is rather uncertain. Analysis of the radio afterglow \\citep{soderberg04} may eventually yield a tighter constraint on the total kinetic energy of the ejecta. Lastly, given the low volume associated with the universe at $z \\lesssim 0.1$, one surmises that the frequency distribution of GRB energies likely increases to lower energy. This may be naturally explained as the effect of viewing angle to the GRB event \\citep[e.g.][]{woosley99}, but could also be the result of an intrinsically broad distribution of GRB luminosities \\citep[e.g.][]{kouveliotou04}. In either case, the discovery of GRB~031203 significantly increases the likelihood of detecting many additional low $z$, low energy GRB events with the launch of {\\it Swift}.  {\\it Note Added After Original Submission:} Newly surfaced $\\gamma$-ray satellite data in Sazonov, Lutovinov, \\& Sunyaev (2004) (recently relayed in Gal-Yam et al.~2004), appears to exclude the XRF hypothesis. This confirms our prior inferences --- that this event was a low-luminosity GRB and not an XRF --- presented in section 4.3."
    },
    "0402/gr-qc0402071_arXiv.txt": {
        "abstract": "s{ In the framework of a flat FLRW model we derive an inflationary regime in which the scalar field, laying on the plateau of its potential, admits a linear time dependence and remains close to a constant value. \\\\ The behaviour of inhomogeneous perturbations is determined on the background metric in agreement to the ``slow-rolling''  approximation. We show that the inhomogeneous scales which before inflation were not much greater then the physical horizon, conserve their spectrum (almost) unaltered after the de Sitter phase. } ",
        "introduction": "The Standard Cosmological Model (SCM) finds many confirmations in the picture of the actual Universe \\cite{KT90}, but its shortcomings to describe very early stages of evolution appear as soon as the so-called  horizon and flatness paradoxes are taken into account \\cite{G81,L82,L83}. \\\\ The Inflationary Paradigm (IP) has acquired progressively an increasing interest, because it provides a natural explanation for such paradoxes \\cite{G81,L82,L83}; the IP success relies overall on the consistent and simultaneous treatment of many different aspects of the cosmological puzzle. The capability to generate a Harrison--Zeldovich Spectrum (HZS) for the density perturbations outstands among these, and in fact the IP predicts it from the quantum fluctuations of the scalar field during the de Sitter phase. \\\\ This picture is well-grounded, but has to face the delicate point regarding the mechanism by which the quantum inhomogeneities approach a classical limit \\cite{Sta79}. The quantum origin of the perturbations is also supported by the exponential suppression of the ultra-relativistic inhomogeneities during the de Sitter phase \\cite{IM03}.  In this work we show the existence of an inflationary regime allowing a classical origin for the HZS. In fact we deal with perturbations of the scalar field $\\phi$  which, if described by a HZS before inflation, survive to the Universe exponential expansion; in our solution the inhomogeneities become super-horizon--sized  and become seed for the structure formation when they re-enter the horizon after the IP. ",
        "conclusions": ""
    },
    "0402/astro-ph0402036_arXiv.txt": {
        "abstract": "{Very recently, the discovery of the {\\it largest known planetary nebula on the sky} surrounding the DO white dwarf \\object{PG\\,1034+001} with an apparent diameter of about 2\\degr, corresponding to a linear diameter of 3.5 - 7.0 pc at the likely distance of 100 - 200 pc, has been reported by Hewett \\ea (2003). A careful inspection of available sky survey data has now shown that this planetary nebula, \\object{Hewett\\,1}, is surrounded by an elliptical emission shell with an apparent diameter of 6\\degr\\,$\\times$\\,9\\degr\\ ($16.2^{+6.1}_{-4.5} \\times 24.3^{+9.1}_{-6.8}\\, \\mathrm{pc}$ at $d = 155^{+58}_{-43}\\ \\mathrm{pc}$). A further emission structure, detected northeast of the central star may indicate another shell with a size of 10\\degr\\,$\\times$\\,16\\degr. >From presently available observational data we do not have indications whether the emission arises from material which was ejected from \\object{PG\\,1034+001} or from ionized ambient ISM.  Improved proper motion data combined with radial velocity and distance from the literature have enabled us to derive a Galactic orbit for the central star \\object{PG\\,1034+001}. Its thin disk orbit and the morphology of the first halo suggest that the nebula is in an advanced stage of interaction with the interstellar medium. ",
        "introduction": "\\label{int} \\object{PG\\,1034+001} ($\\alpha_\\mathrm{2000} = 10^\\mathrm{h} 37^\\mathrm{m} 04\\fs0$, $\\delta_\\mathrm{2000} = -0\\degr 8\\arcmin 20\\arcsec$) is a very hot, hydrogen-deficient DO white dwarf. Werner \\ea (1995) determined \\Teffw{100^{+15}_{-10}}, \\loggw{7.5\\pm 0.3} (cgs) by means of NLTE model atmosphere techniques. From its position in the $\\log g$ - $\\log \\Teff$ plane \\sA{evo}, they concluded that \\object{PG\\,1034+001} is a successor of the hydrogen-deficient PG\\,1159 stars.  Very recently, Hewett \\ea (2003) found by investigation of spectra from the Sloan Digital Sky Survey (SDSS, York \\ea 2000) that, at a distance of $155^{+58}_{-43}\\ \\mathrm{pc}$ (Werner \\ea 1995), a planetary nebula (PN) with a linear diameter of about 3.5 - 7.0\\,pc surrounds \\object{PG\\,1034+001}. This is larger than the PN \\object{Sh\\,2-216} with a size of 3.5\\,pc at distance of $130^{+9}_{-8}\\ \\mathrm{pc}$ (Harris \\ea 1997). Only one PN, namely \\object{PN\\,G080.3-10.4} (\\object{MWP\\,1}) surrounding the PG\\,1159 star \\object{RX\\,J2117.1+3412} (Appleton \\ea 1993), may be physically larger. At a distance of $1.4^{+0.7}_{-0.5}\\ \\mathrm{kpc}$ (Motch \\ea 1993), its apparent diameter of 14\\arcmin 45\\arcsec\\ (Rauch 1997) corresponds to a physical size of $6^{+3}_{-2}\\ \\mathrm{pc}$.  However, whatever an improvement of the distance determinations for the PN around \\object{PG\\,1034+001} or \\object{RX\\,J2117.1+3412} may yield --- our inspection of the Southern H-Alpha Sky Survey Atlas (SHASSA, Gaustad \\ea 2001) has shown two much more extended emission structures surrounding \\object{PG\\,1034+001} which might make this object the largest PN known to date. In the following, we will describe our discovery in more detail. ",
        "conclusions": "\\label{con}  We have discovered a huge elliptical halo with a linear diameter of $16.2^{+6.1}_{-4.5} \\times 24.3^{+9.1}_{-6.8}\\, \\mathrm{pc}$ at $d = 155^{+58}_{-43}\\ \\mathrm{pc}$ surrounding the PN \\object{Hewett\\,1} and its exciting star \\object{PG\\,1034+001} on SHASSA images. An even larger emission structure is also visible but much fainter. >From presently available observational data we do not have indications whether the emission arises from material which was ejected from \\object{PG\\,1034+001} or from ionized ambient ISM. Clearly better spectra and images as well as a better distance determination are needed in order to understand the nature of the PN surrounding \\object{PG\\,1034+001}.  It is worthwhile to note that the largest PNe known so far, \\object{Sh 2-68} (Kerber \\ea 2002b), \\object{Hewett\\,1}, and \\object{MWP\\,1}, have been found around so-called born-again post-AGB type central stars (Iben \\ea 1983). Since their exciting stars have quite different masses (0.50\\,M$_\\odot$, 0.62\\,M$_\\odot$, and 0.83\\,M$_\\odot$, respectively, see \\ab{evo}), the size of the PN does not appear to be related to the final mass of the ejecting star.  We have derived the Galactic orbit of \\object{PG\\,1034+001} and determined that it belongs to the thin disk population. The kinematical parameters and the morphology of the halo indicate that the nebula is in an advanced stage of interaction with the ISM. Since this is one of the closest and most evolved PNe, it is ideally suited for detailed study of this process which governs the return of processed matter to the ISM."
    },
    "0402/astro-ph0402200_arXiv.txt": {
        "abstract": "Our current knowledge of neutron star formation, progenitors, and natal masses, spins, magnetic fields, and space velocities is briefly reviewed from a theorist's perspective. More observational information is badly needed to constrain theoretical possibilities. ",
        "introduction": "Although only a wink in its life, the moment of birth of a neutron star marks a spectacular astrophysical event with far-reaching consequences. Neutron stars originate from the apocalyptic death of massive stars in supernova (SN) explosions. While Baade and Zwicky (1934) first came up with this visionary suggestion, the link is now unambiguously established by associations of pulsars and compact X-ray sources with young gaseous SN remnants, e.g.\\ in the famous cases of the Crab pulsar and Crab nebula, Vela pulsar and nebula, or the Cassiopeia A remnant with the compact central object that was pinpointed with high resolution by the {\\em Chandra} X-ray Observatory.  Neutron stars certainly belong to the most exotic known objects. With the size of roughly three gravitational radii they contain more than a solar mass of matter at a density exceeding that in atomic nuclei. The gravitationally bound object is kept in mechanical equilibrium by repulsive interactions and degeneracy pressure of nucleons that balance the enormous pull of gravity. The extreme compactness allows pulsars to rotate with periods as low as a millisecond and to possess surface magnetic fields up to 15 orders of magnitude higher than that of the Earth. Extraordinary conditions like these make them unique astrophysical laboratories for nuclear physics, particle physics, and gravitational physics.   \\begin{figure}[htp!] \\centerline{\\psfig{file=jankah_1.ps,width=0.49\\textwidth} \\psfig{file=jankah_2.ps,width=0.49\\textwidth}} \\vspace{7pt} \\centerline{\\psfig{file=jankah_3.ps,width=0.49\\textwidth} \\psfig{file=jankah_4.ps,width=0.49\\textwidth}} \\caption{\\small Meridional cuts through a differentially rotating, convective nascent neutron star 0.75 seconds after its formation. The rotation axis coincides with the ordinates of the plots. In a quasi-steady state, convection is strongest at intermediate radii in a region of essentially constant specific angular momentum near the equator. It is only weakly developed closer to the rotation axis where elongated, convective cells occur that are aligned parallel to the axis. A steep gradient of the specific angular momentum suppresses convective motion perpendicular to the axis in this region. In contrast, in a non-rotating star convective activity takes place in a spherical shell. The hydrodynamic simulation was carried out with neutrino diffusion taken into account (Keil 1997, Janka \\& Keil 1998, Janka et al.\\ 2001). {\\em Top left:} Contours of constant density (solid lines) between $3.67\\times 10^{10}\\,$g$\\,$cm$^{-3}$ and $2.68\\times 10^{14}\\,$g$\\,$cm$^{-3}$, increasing with a factor of 1.37, and of constant temperature (dotted lines) between 4$\\,$MeV and 30$\\,$MeV with steps of 1$\\,$MeV. {\\em Top right:} Lepton fraction $Y_{\\mathrm{l}}$ (density of electrons plus electron neutrinos minus their antiparticles relative to the number density of nucleons). {\\em Bottom left:} Specific angular momentum $j_z$; the rotation period is $\\sim$1$\\,$ms at 3$\\,$km distance from the rotation ($z$) axis, $\\sim$2.5$\\,$ms at 10$\\,$km and $\\sim$6$\\,$ms at 20$\\,$km. {\\em Bottom right:} Total velocity in radial and lateral directions. The arrows indicate the flow direction in a meridional plane.} \\label{janka_fig1} \\end{figure}   The birth of a neutron star constitutes the transition of matter on a macroscopic scale to the most extreme state realized after the big bang. The throes are signaled by the conversion of huge amounts of gravitational binding energy mostly to neutrinos (up to $\\sim$99\\% or several 10$^{53}\\,$erg), some to kinetic energy of the explosion ejecta or wind loss ($\\sim$1\\%), and minor parts to electromagnetic radiation ($\\sim$10$^{49}\\,$erg) and gravitational waves ($\\sim$10$^{46}\\,$erg, possibly more). These forms of energy release suggest potential observability, but the rate of nearby SNe is rather low and available empirical data are correspondingly sparse. Much of our knowledge of neutron star formation and birth properties like mass distribution, spins, magnetic fields, proper motions, is therefore based on theoretical work, which, however, is hampered by the enormous complexity of the problem and barely constrained degrees of freedom, e.g.\\ in the initial conditions or input physics for models. A brief moment in evolution therefore poses a big challenge for exploration. ",
        "conclusions": ""
    },
    "0402/astro-ph0402493_arXiv.txt": {
        "abstract": "{We have carried out calculations of ionization equilibrium and deuterium fractionation for conditions appropriate to a completely depleted,  low mass pre--protostellar core, where heavy elements such as C, N, and O have vanished from the gas phase and are incorporated in ice mantles frozen on dust grain surfaces.  We put particular emphasis on the interpretation of recent observations of \\HTDP \\ towards the centre of the prestellar core L~1544 (Caselli et al. 2003) and also compute the ambipolar diffusion timescale.  We consider explicitly the ortho and para forms of \\MOLH, \\HTHP, and \\HTDP. Our results show that the ionization degree under such conditions depends sensitively on the grain size distribution or, more precisely, on the mean grain surface area per hydrogen nucleus.  Depending upon this parameter and upon density, the major ion may be H$^{+}$, \\HTHP, or \\DTHP.  We show that the abundance of ortho-\\HTDP \\ observed towards L~1544 can be explained satisfactorily in terms of a complete depletion model and that this species is, as a consequence, an important tracer of the kinematics of prestellar cores. ",
        "introduction": "Determining the physical structure of pre--protostellar cores is one of the keys to understanding the development of protostars. In the simplest view of the development of such cores, they contract slowly towards some ``pivotal state'', subsequent to which dynamical collapse sets in. Clearly, one should attempt to determine the parameters of this pivotal state, since these parameters dictate the subsequent evolution. It seems probable that the pivotal state is marked by high column density and low temperature. Likely candidates can be identified using measurements of dust emission and absorption (e.g. Andr\\'{e} et al. 2000), which indicate temperatures (of both gas and dust) of about 10~K or below, molecular hydrogen column densities in the range $10^{22}$ to $10^{23}$ cm$^{-2}$, and sizes of several thousand AU.   An additional result from recent observational studies of prestellar cores is that depletion of molecular species on to dust grain surfaces is also a marker of relatively evolved cores (i.e. close to the pivotal state) with high central densities and column densities (e.g. Caselli et al. 2002a, Tafalla et al. 2002). Tafalla et al. show that CO becomes depleted by at least a factor of 10, relative to its canonical abundance, above densities of $5\\, 10^4$ \\percc . CS shows  similar behaviour,  whereas nitrogen containing species such as \\AMM \\ and \\NTHP \\ have abundances which are either constant or (in the case of ammonia) increase somewhat in the high density gas surrounding the dust emission peak. Their interpretation of these observations is that molecular nitrogen (\\MOLN), which is the  source  of the  observed \\AMM \\ and \\NTHP , is sufficiently volatile to remain in the gas phase at densities of order $10^5$ \\percc \\ (essentially because of spot heating of the grain surfaces by cosmic rays), whereas CO (the source of carbon for most C-containing species) is not.  What happens at still higher densities? It seems likely that, at sufficiently high densities, all species containing heavy elements will condense out. The difference in sublimation energies for CO and \\MOLN \\ is not great (e.g. Bergin and Langer 1997) and thus one might expect \\MOLN \\ and other N--containing species to disappear from the gas phase at densities only somewhat higher (or dust temperature slightly lower) than found for CO. Whilst it is true that the laboratory experiments may not accurately simulate interstellar ice surfaces, it seems probable that \\MOLN \\ disappears from the gas phase at densities above $10^6$ \\percc .  In fact, the densities inferred from mm dust continuum emission in several objects are of order $10^6$ \\percc \\ and there  is some evidence for ``holes'' in the \\NTHP \\ distribution at densities above $3\\, 10^5$ \\percc .  For example, Bergin et al. (2002) find a flattening of the \\NTHP \\ column density distribution in B68 which could be interpreted as being due to an absence of \\MOLN \\ towards the dust peak. Belloche (2002) finds a ``\\NTHP \\ hole'' around the peak of the Class 0 source IRAM04191, which is most easily interpreted in terms of \\MOLN \\ condensing out at the highest densities in the prestellar core. Last, but not least, Caselli et al. (2003, in what follows CvTCB) have detected compact emission (2000 AU in size) in the $1_{10}-1_{11}$ \\HTDP \\ line  towards the dust emission peak of L~1544. This latter discovery is significant in view of the fact that CvTCB infer an abundance of order $10^{-9}$ relative to \\MOLH . Estimates of the ionization degree in L~1544  (Caselli et al. 2002a) are of order $2\\, 10^{-9}$. Taken together, these results imply that \\HTDP (and hence probably \\HTHP ) is a major ion. It seems likely that this can be the case only if ions containing heavy elements, such as \\NTHP \\ and \\HCOP \\, are absent, which in turn can be true only if the heavy element content of the gas has condensed on to grain surfaces.  If this is the case, one must rely on species lacking heavy elements (such as \\HTDP ) to trace the kinematics of the high density nucleus of the pre--protostellar core.  Also, it becomes relevant to establish the ionization degree and other physical parameters of a region where species containing heavy elements have condensed out.  In this paper, we study some of the consequences of complete depletion in  an isolated, low mass prestellar core. In particular, we compute the ionization degree and its dependence upon physical parameters such as the gas density and the grain size. We consider in detail deuterium fractionation in these conditions and try to explain the observed \\HTDP \\ (or, more precisely, the ortho-\\HTDP ) abundance in objects such as L~1544. In section 2, we discuss the model assumptions and, in section 3, we summarize the results. Then, in section 4, we consider some of the implications of our calculations and in section 5, we briefly summarize our conclusions. Some preliminary results of our study have been presented elsewhere (Walmsley et al. 2004). ",
        "conclusions": "We summarize here our most important conclusions. \\begin{itemize} \\item We believe that our data strongly suggest that the detection of \\HTDP \\ by CvTCB is due to complete freeze--out of the heavy elements in the core of L~1544. If so, {\\it only} species such as \\HTHP \\ and its deuterated counterparts are useful tracers of the kinematics in the dense centre of L~1544 and other similar objects. \\item Knowledge of the mean grain surface area is needed in order to understand phenomena such as deuterium fractionation; it seems likely that this remains true even if some of the heavy species are in the gas phase. Our results imply that, in L~1544, the mean grain surface area per H--atom is above $10^{-21}$ cm$^{2}$, which is comparable to values found in the diffuse interstellar medium. \\item The degree of ionization and the timescale for ambipolar diffusion, determined under conditions of complete heavy element depletion, are not significantly different from values reported elsewhere in the literature, evaluated assuming degrees of depletion more typical of dense interstellar clouds. However, we estimate that, under the conditions considered here, the effect of the coupling between the neutral gas and charged grains is to increase the ambipolar diffusion timescale by between one and two orders of magnitude; this finding will be the subject of further investigation. \\end{itemize}  \\appendix"
    },
    "0402/astro-ph0402346_arXiv.txt": {
        "abstract": "We have developed an algorithm to find haloes in an $N$-body dark matter simulation, called {\\scshape voboz} ({\\scshape vo}ronoi {\\scshape bo}und {\\scshape z}ones), which has as little dependence on free parameters as we can manage.  By using the Voronoi diagram, we achieve nonparametric, `natural' measurements of each particle's density and set of neighbors. We then eliminate much of the ambiguity in merging sets of particles together by identifying every possible density peak, and measuring the probability that each does not arise from Poisson noise.  The main halo in a cluster tends to have a high probability, while its subhaloes tend to have lower probabilities.  The first parameter in {\\scshape voboz} controls the subtlety of particle unbinding, and may be eliminated if one is cavalier with processor time; even if one is not, the results saturate to the parameter-free answer when the parameter is sufficiently small. The only parameter which remains, an outer density cut-off, does not influence whether or not haloes are identified, nor does it have any effect on subhaloes; it only affects the masses returned for supercluster haloes. ",
        "introduction": "A crucial step in comparing $N$-body simulations to the observed galaxy distribution is to identify the possible sites of galaxy formation, called dark-matter haloes, in the simulations.  Unfortunately, the concept of a dark-matter halo is not precisely defined.  There is no firm observational definition of dark-matter haloes, since they can only be observed indirectly; for example, through gravitational lensing.  There are a couple of possible theoretical definitions.  One of them is a region exceeding a certain overdensity, such as the canonical overdensity of virialization, 200.  This is often used when seeking Halo Occupation Distributions (e.g.\\ Berlind \\& Weinberg 2002), which statistically characterize the number of galaxies inside haloes (implicitly dark-matter hereafter) as a function of halo mass.  However, if we want to look beyond the statistical placement of galaxies inside haloes, we should use another definition of a halo (or subhalo): a density peak to which some mass is gravitationally bound.  In the language of an $N$-body simulation, a particle is the core of a halo if it is a local density maximum, and there exists at least one other particle bound to it.  One of the first halo-finding algorithms (HFAs), still in wide use because it is so fast and conceptually simple, is the Friends-of-Friends algorithm (Davis et al.\\ 1985).  This HFA groups together all particles within a specified linking length, a free parameter which is usually set by the canonical overdensity of virialization.  Friends-of-Friends is useful if one is looking for large structures exceeding this overdensity, but it is incapable of finding subhaloes within these structures, and sometimes structures are unduly linked if there happens to be a stream of particles connecting them.  Most HFAs developed since Friends-of-Friends begin with an explicit measurement of the density, which is not uniquely or obviously defined given a set of particles.  In one Eulerian method ({\\scshape denmax}, Bertschinger \\& Gelb 1991), each particle is smoothed with a Gaussian of a fixed spatial resolution.  As with the Friends-of-Friends algorithm, the free parameter is set roughly by the critical overdensity of virialization. While this value of the parameter tends to give virialized objects, it smears out subhaloes (Neyrinck, Hamilton \\& Gnedin 2004, hereafter NHG); one runs the risk of missing structures smaller than any fixed smoothing length.  On the other hand, using a smoothing length that is too small misses the less-dense outskirts of haloes.  Another HFA, called {\\scshape bdm} (Klypin \\& Holtzman 1997), finds density maxima by placing spheres randomly in the simulation, and then moving them at each iteration to the center of mass of particles within them.  Maxima are then joined if they lie within a specified radius.  Another way to find the density, called {\\scshape skid} (Weinberg, Hernquist \\& Katz 1997, Jang-Condell \\& Hernquist 2001) uses a Lagrangian, `smoothed particle hydrodynamics' (SPH) density estimate based on the distances to the nearest $N_{dens}$ particles.  This density estimate is arguably an improvement over {\\scshape denmax}'s because there is no fixed spatial resolution, but in its place there is an arbitrary, fixed mass resolution.  This is undesirable because haloes can exist with only two particles, which a fixed mass resolution is likely to miss.  The next step in halo finding is to group the particles together.  In {\\scshape denmax} and {\\scshape skid}, particles slide along density gradients until they reach density maxima.  In {\\scshape hop} (Eisenstein \\& Hut 1998), which uses a Lagrangian density estimator similar to that in {\\scshape skid}, each particle `hops' to the densest particle among its neighbors, and continues in this manner until it reaches a local density maximum.  Then, groups of particles are joined together if saddles linking them exceed a specified density, another free parameter.  Recently, Kim \\& Park (2004) have developed another HFA, called {\\scshape psb}.  Around each density maximum (calculated with a small spatial smoothing length), {\\scshape psb} finds the largest isodensity contour enclosing only that peak, using a couple of parameters to do so.  It then assigns neighboring particles to density peaks using considerations such as whether they are energetically bound, and whether they lie outside the tidal radius.  It seems that {\\scshape psb} reliably uncovers subhaloes, but it does require several free parameters. ",
        "conclusions": "We have developed a halo-finding algorithm, called {\\scshape voboz}, which is nearly parameter-free, and which has a resolution limited only by the discreteness of particles in the simulation it is analysing.  The Voronoi diagram allows us to fix the densities and the sets of neighbors for all particles in a `natural,' parameter-independent way, on arguably the finest possible scale that contains meaningful structure.  Further degrees of freedom are eliminated by assigning to each halo (or subhalo) a probability that it exists, i.e.\\ that it did not arise from Poisson noise.  Of the two parameters in {\\scshape voboz}, one of them exists only to save processor time, and needs not be used if one has no processor time constraints.  This parameter controls the delicacy with which particles are unbound, and the results saturate at the `right' answer (the parameter-free situation in which one particle is unbound at a time) when the unbinding becomes sufficiently delicate.  Additionally, most of the extra haloes uncovered at extreme delicacy have meager probabilities.  The remaining parameter is a density cut-off, necessary because haloes do not extend to low densities in the real universe.  However, it only affects the masses of the largest haloes in the simulation, and not the masses or detection of subhaloes, which is where reliability in halo-finding algorithms is most needed.  An ideal halo-finding algorithm would find groups of particles in both position and velocity space, or even consider neighboring time slices in a simulation.  In {\\scshape voboz}, the velocities are used only to decide if particles are energetically bound to haloes found in position space.  This is an approximate criterion, but it seems to work acceptably well.  The probability {\\scshape voboz} returns for each halo is certainly tied to the finite mass resolution of the simulation it is analysing.  However, as we decrease a cut-off in halo probability, an interesting signal emerges in the halo correlation function which resembles the 1-halo term in the halo model of large-scale structure.  This suggests an interpretation of halo probability as a measure of a halo's `sub-haloness,' which may increase {\\scshape voboz}'s appeal for researchers of the halo model.  The {\\scshape voboz} code is publically available, at \\url{http://casa.colorado.edu/~neyrinck/voboz/}."
    },
    "0402/astro-ph0402583_arXiv.txt": {
        "abstract": "We develop a new approach toward a high resolution non-parametric reconstruction of the primordial power spectrum using WMAP cosmic microwave background temperature anisotropies that we confront with SDSS large-scale structure data in the range $k\\sim0.01-0.1\\,h\\mathrm{Mpc^{-1}}$. We utilise the standard $\\Lambda$CDM cosmological model but we allow the baryon fraction to vary. In particular, for the concordance baryon fraction, we compare indications of a possible feature at $k\\sim0.05\\,h\\mathrm{Mpc^{-1}}$ in WMAP data with suggestions of similar features in large scale structure surveys. ",
        "introduction": "The Wilkinson Microwave Anisotropy Probe (WMAP) experiment has recently provided a high resolution full sky cosmic microwave background (CMB) map that provides an unprecedented opportunity to probe the structure and the contents of the  universe (Spergel et al. 2003). The corresponding angular power spectrum has been used to set  constraints on cosmological parameters. In the framework of $\\Lambda$CDM cosmologies, such constraints are strong if the initial conditions are parametrized by primordial power spectra in the form of power laws with or without running spectral index. However parameter estimation can be weakened and biased if more complex forms of initial conditions are adopted, cf. \\cite{dunkley} and \\cite{bla}. We explore another such possibility here.  Recent papers from the Sloan Digital Sky Survey (SDSS) team report measurements of the matter power spectrum and the cosmological parameters that give a best fit to its shape (Tegmark et al. 2003a,b; Pope et al. 2004). In particular Pope et al. (2004) finds that for the for a value of the matter density, $\\Omega_{m}h=0.264\\pm0.043$ and baryon fraction, $f_{b}=\\Omega_{b}/\\Omega_{m}=0.286\\pm0.065$ at 1 $\\sigma$. These determinations, in particular that of the baryon density, are in marginal conflict with the data from both WMAP (Spergel et al. 2003) and primordial nucleosynthesis (BBN) (Cyburt et al. 2003; Cuoco et al. 2003), which, together with the Hubble constant (Freedman et al. 2001) determination, prefer a lower baryon fraction, $f_{b}\\simeq0.2$ when a power law primordial power spectrum in a $\\Lambda$CDM cosmology is adopted.  Such a discrepancy {\\it might} be due to the presence of intrinsic features in the power spectrum that affect the SDSS (and other survey) data, such as for example a dip and a bump at $k\\sim0.035 \\,h \\mathrm{Mpc}^{-1}$ and $k\\sim0.05 \\,h \\mathrm{Mpc}^{-1}$ (see also Atrio-Barandela et al. 2001; Barriga et al. 2001). It is useful to explore the implications of such a possibility, even though the data at present is not compelling, because very similar features in the matter power spectrum have been detected in independent large-scale structure surveys, notably the 2dF, Abell/ACO cluster and PSCz surveys (Einasto et al. 1997; Percival et al. 2001; Miller \\& Batusky 2001; Hamilton et al. 2000). In this work, we show that, if the power law hypothesis for the shape of the initial power spectrum is relaxed, a still good or even better agreement between the different data sets can be obtained. Our strategy is principally comprised of two steps: a high definition reconstruction of the primordial power spectrum from WMAP data (for given sets of cosmological parameters) and a convolution of the initial power spectrum with the SDSS window functions to determine the statistical agreement between the two data sets.  We choose to study only these two data sets in a limited range of $k-$space in order to make a preliminary exploration of the possible influence of primordial features.  We find that for reasonable choices of the cosmological parameters, the WMAP-reconstructed matter power spectrum shows oscillations that are very similar to those of SDSS once the convolution is performed. This results in an improved concordance for the inferred baryon fraction and at the same time is suggestive of a deviation from a power-law primordial power spectrum.   \\section[]{Method}  The temperature CMB angular power spectrum is related to the primordial power spectrum, $P^{0}$, by a convolution with a window function that depends on the cosmological parameters: \\begin{equation} C_{\\ell}=4\\pi\\int\\Delta^{2}_{\\ell}(k)\\, P^0(k) \\, \\frac{\\mathrm{d} k}{k} \\simeq WP^{0}, \\end{equation} where the last term is in matrix notation and represents a numerical approximation to the integral, calculated using a modified version of CMBFAST (Seljak \\& Zaldarriaga 1996). The $C_{\\ell}$ spectrum then contains information about both the cosmological parameters and the initial power spectrum.  Various papers have presented reconstructions of the power spectrum from WMAP data (e.g. Bridle et al. 2003; Mukherjee \\& Wang 2003; Matsumiya et al. 2003; Shafieloo \\& Souradeep 2004; Kogo et al. 2004; Hannestad 2004). Most of these have confronted an \\textit{a priori} parametrised power spectrum with the data to obtain information about its shape and amplitude. Following the spirit of \\cite{gaw} and \\cite{tegzald}, our power spectrum is \\textit{not} described by a small set of parameters that incorporate features, but is finely discretized in $k$-space (see also Shafieloo \\& Souradeep 2004; Kogo et al. 2004). Our work represents the first attempt to reconstruct the power spectrum at high resolution in the full range $k\\sim0.01-0.1\\,h\\mathrm{Mpc^{-1}}$ together with an estimation of the error covariance matrix. Our findings are consistent with the results obtained with a different method by \\cite{kogo} that are limited to $k\\lesssim0.045\\,h\\mathrm{Mpc^{-1}}$, corresponding to $\\ell\\lesssim430$.  Generally speaking, it is not possible to invert the matrix $W$ to solve for the power spectrum, because each $C_{\\ell}$ embraces information about a limited range in $k$-space. However the inversion is feasible under some assumptions. Basically the lack of information can be remedied by the introduction of priors, such as for example the requirement of a certain degree of smoothness in the solution. To be specific, we consider a solution of the form \\begin{equation} \\overline P=MC^{WMAP}_{\\ell} \\end{equation} with an error covariance matrix given by \\begin{equation} \\Sigma=\\left[ W^{t}N^{-1}W+\\epsilon L^{t}L \\right]^{-1}. \\end{equation} Here: $M=\\left[ W^{t}N^{-1}W+\\epsilon L^{t}L \\right]^{-1}W^{t}N^{-1}$; $C^{WMAP}_{\\ell}$ are the mean values of WMAP data; $N$ is the WMAP covariance matrix (Verde et al. 2003) and $L$ is a discrete approximation to the first derivative operator.  Our approach follows the spirit of a similar solution proposed in \\cite{tegzald}; however we replace the identity operator by $L$.  The parameter $\\epsilon$ regulates the degree of solution smoothness, since the derivative operator $L$ acts in such a way to minimize the norm of the solution derivative. This solution can be effectively thought of as a linear least squares solution modified to accommodate smoothing.  In order to select a value for the parameter $\\epsilon$, we calculate the angular power spectrum $\\overline C_{\\ell}$ resulting from $\\overline P$. Then we form the expression \\begin{equation} \\chi^{2}=(\\overline C_{\\ell}-C^{WMAP}_{\\ell})^{t}N^{-1}(\\overline C_{\\ell}-C^{WMAP}_{\\ell}). \\end{equation} It is reasonable to fix the parameter $\\epsilon$ to render the $\\chi^{2}$ in the previous equation equal to the number of WMAP temperature data points, namely 899. In other words, a solution is chosen that gives an acceptably good fit to the data. Moreover, as will be explained below, this represents a conservative choice since the smoothing decreases the effective number of degrees of freedom.  To test the method, we reconstruct the power spectrum by using angular power spectra calculated from various power spectrum shapes with and without added errors. We obtain very good results, except when there are features of extension comparable to our $k$-space gridding $\\Delta k$, which we took to be approximately $10^{-4}$ in the range $0.01-0.1 \\,\\mathrm{Mpc}^{-1}$. The integration range is typically between $10^{-5}$ and $0.3\\,\\mathrm{Mpc}^{-1}$ and comprises about 1500 points. We also verify that the signals detected in the actual WMAP data (see below) do not depend on discretization issues by testing different resolutions. Our method yields a $k$-space resolution that is limited only by computing resources, at the price however of introducing correlations between neighboring points that we take fully into account into our statistical analysis.  In analogy with ordinary regression it is possible to define the effective number of degrees of freedom caused by smoothing by the expression $\\mathrm{Trace}\\left( WM \\right)$. It is easy to recognize that without smoothing this quantity would have been equal to the number of bins in $k$-space. We have found that our choice of the parameter $\\epsilon$ implies the effective usage of approximately 45 constraining informations out of the more than 1000 amplitudes that specify the power spectrum.  The procedure just described must be repeated for each given set of cosmological parameters. We proceed by first evolving the primordial power spectrum to redshift zero by multiplying it by the appropriate transfer function calculated by CMBFAST. Then we convolve the derived matter power spectrum with the SDSS window functions given by \\cite{sdss1}. Finally, we evaluate $\\chi^{2}$: \\begin{equation} \\chi^{2}=\\left(\\widetilde P-\\frac{P^{SDSS}}{b^{2}} \\right)^{t}\\left[\\widetilde\\Sigma+\\frac{N^{SDSS}}{b^{4}}\\right]^{-1}\\left( \\widetilde P-\\frac{P^{SDSS}}{b^{2}} \\right), \\label{cfr} \\end{equation} where $P^{SDSS}$ and $N^{SDSS}$ represent the SDSS mean data and (diagonal) covariance matrix (Tegmark et al. 2003a), $\\widetilde P$ and $\\widetilde \\Sigma$ are the SDSS-convolved WMAP-reconstructed matter power spectrum and covariance matrix (calculated by error propagation) and $b$ is the bias parameter. We follow the common practice of considering a constant bias which we set at the optimal value. ",
        "conclusions": "In summary we have shown that a slight discrepancy in the baryonic fraction that arises from SDSS and WMAP data could be resolved if the primordial power spectrum is not fixed \\textit{a priori} in the form of a power law, but is constrained to satisfy CMB data when confronted with large-scale structure information and letting key cosmological parameters vary. Our method is feasible due to a high resolution reconstruction of the power spectrum from the CMB angular power spectrum, that is carried out for the first time in the full range considered, and provides evidence in favour of intrinsic features in the primordial power spectrum. In particular a feature at $k\\sim0.05\\,h\\mathrm{Mpc^{-1}}$ seems to occur in both the WMAP and SDSS data sets. This kind of technique offers an effective way to break the degeneracy between the baryon abundance and the shape of the power spectrum (in the conventional approach parametrized by the spectral index, $n_{s}$) and relevantly brings also to a measure of the galaxy bias.  A power spectrum with the observed features brings to a good fit of SDSS data without the necessity of a large baryon fraction. More generally, if the initial power law hypothesis is relaxed, deviations seem to induce a somewhat better concordance picture from the combination of CMB, large-scale structure, BBN and Hubble constant measurements. This could have important consequences for fundamental physics."
    },
    "0402/astro-ph0402256_arXiv.txt": {
        "abstract": "{ We use a Monte Carlo code to generate synthetic near-IR reflection nebulae that resemble those (normally associated with a bipolar outflow cavity) seen towards massive young stellar objects (YSOs). The 2D axi-symmetric calculations use an analytic expression for a flattened infalling rotating envelope with a bipolar cavity representing an outflow. We are interested in which aspects of the circumstellar density distribution can be constrained by observations of these reflection nebulae.  We therefore keeep the line of sight optical depth constant in the model grid, as this is often constrained independently by observations. It is found that envelopes with density distributions corresponding to mass infall rates of $\\sim 10^{-4} $\\macco\\ (for an envelope radius of 4700 AU) seen at an inclination angle of $\\sim $\\degg{45} approximately reproduce the morphology and extension of the sub-arcsecond nebulae observed in massive YSOs. Based on the flux ratio between the approaching and receding lobe of the nebula, we can constrain the system inclination angle. The cavity opening angle is well constrained from the nebula opening angle. Our simulations indicate that to constrain the outflow cavity shape and the degree of flattening in the envelope, near-IR imaging with higher resolution and dynamic range than speckle imaging in 4m-class telescopes is needed. The radiative transfer code is also used to simulate the near-IR sub-arcsecond nebula seen in \\mon. We find indications of a shallower opacity law in this massive YSO than in the interstellar medium, or possibly a sharp drop in the envelope density distribution at distances of $\\sim$ 1000~AU from the illuminating source. ",
        "introduction": "\\label{IntroductionSection}  Bipolar outflows appear to be a ubiquitous phenomenon during the formation of stars in all mass ranges \\citep{BallyLada83,Henning00,RidgeMoore01,Beuther02}. Low mass young stellar objects (YSOs) show highly collimated bipolar jets from a few 10 ~AU \\citep{Burrows96} to several parsec \\citep{Reipurth97,Eisloeffel00} in length. These jets are thought to be magneto-hydrodynamically collimated in a wind formed at the inner star-disk system \\citep[e.g. X-wind, ][]{Shu94}. The jets are thought to drive the large scale molecular outflow \\citep{Masson94}.  The formation and collimation of outflows in massive YSOs is less well understood than in low mass YSOs. There appears to be a lack of highly collimated parsec-scale jets \\citep{Mundt94}. In the near-IR, searches for shock-excited \\molh\\ show traces of jets in massive star forming regions, but probably driven by low mass young stars located in the same cluster \\citep{Davis98,Wang03}. A recent search for optical shock-excited emission in the outer parts of the outflow, yielded no evidence of jet interaction (Alvarez \\& Hoare, in prep.). Very close to the driving source, there is also not clear evidence that jets are the rule in massive YSOs. In some cases, the free-free radio emission from the inner wind shows a jet morphology (e.g. HH80-81, \\citealt{Marti93}; Cep~A, \\citealt{Torrelles96}). Such jets would have to be magneto-hydrodynamically driven, even though the OB stars themselves are not magnetically active. Magneto-hydrodynamics in the infalling rotating cloud could set up bipolar flows \\citep{Tomisaka98}. In other cases, the ionised wind appears to be equatorial \\citep[e.g.][]{Hoare94,HoareMuxlow96,Hoare02}. Theoretical models show that radiation pressure in massive young stars can drive gas off the surface of a disk, producing a predominantly equatorial wind \\citep{Drew98, Drew00}. Any initial flow maybe hydrodynamically collimated into a bipolar flow by the flattened surrounding cloud \\citep[e.g. ][]{Delamarter00}. These alternative theories will predict different morphologies for the base of the outflow cavities carved out. These variations in morphology occur at scales of a few 100~AU, which at the typical distances to massive YSOs of $\\sim 1$~kpc, correspond with angular sizes of $\\sim$~0\\farcs1. Therefore, high resolution techniquies are fundamental to study the impact of the the outflow in the surrounding material.  In a related paper (Alvarez et al. in prep., hereafter Paper~I), we show high resolution near-IR speckle images which trace the circumstellar matter around massive YSOs at scales of a few 100~AU. The extended emission that is seen towards some of the sources can be interpreted as scattered light in an outflow cavity due to its monopolar morphology. This intepretation is supported in some cases by the blue colours of the nebula \\citep[e.g. \\mon, Paper I,][]{Preibisch02}. Furthermore, polarimetric speckle imaging of the reflection nebula in the massive star forming region \\sonef\\ \\citep{Schertl00} shows a centrosymmetric pattern which is typical of scattered light.  Intuitively, one can imagine that depending on the properties of the dust, the shape of the cavity, the density distribution and the orientation of the system with respect to the observer, the resulting reflection nebula will change. The morphology of the cavity is particularly important because it is shaped by the interplay between the infall and the outflow. For instance, it is expected that an equatorial or wide-angled wind will produce a cavity with a wide openig angle near the star. However, a jet is expected to open a rather narrow cavity.  Radiative transfer simulations have been widely used to generate synthetic nebulae that resemble the observations \\citep{Lazareff90,WhitneyHartmann92,Kenyon93,Fischer94,Fischer96,Whitney97,Lucas97,Lucas98,Wolf02}. The work by \\citet{Lazareff90} was based on a ray-tracing code and it was focused mainly on the effect produced by different disc models on the synthetic nebulae. The authors compared the general features of the model images with previous seeing-limited images of the low mass systems HL~Tau and L\\,1551\\,IRS5. \\citet{WhitneyHartmann92}, \\citet{Kenyon93} and \\citet{Whitney97} developed a Monte Carlo code to investigate how the nebula morphology and the near-IR colours of the synthetic images vary with different model parameters. In particular, \\citet{Whitney97} used their code to constrain the colours of the central source, the dust model and the envelope density distribution in a sample of $\\sim$\\,20 low mass YSOs. \\citet{Fischer94} and \\citet{Fischer96} developed a new Monte Carlo scattering code and they focused on exploring the effect of different dust models in the synthetic images. \\citet{Lucas97} and \\citet{Lucas98} compared synthetic nebulae produced with their Monte Carlo code with high resolution multi-colour observations of reflection nebulae associated with low mass YSOs. From this comparison, they could constrain some parameters defining circumstellar density distribution as well as the dust model for several sources. Recently, radiative transfer Monte Carlo codes have been developed to simulate scattering by non-spherical dust particles \\cite{Whitney02,Wolf02,Lucas03}.  These previous models have focused predominantly on low mass YSOs. Here, we apply the Monte Carlo code of \\citet{Lucas98} to high mass YSOs, where the infall rates are much higher. We also adopt an observational approach, by presenting a grid of models in which as each parameter is varied, the overall density is scaled too to keep the optical depth along the line of sight constant. This is because the line of sight optical depth is often well constrained from other data such as the optical depth of the 9.7~$\\mu$m silicate feature or the colour of the star. The models are decribed in Sect.~\\ref{ModelsSection}. The grid of models is presented in Sect.~\\ref{GridSection}. In Sect.~\\ref{IRS3ModelSection}, we use the models to constrain the density distribution in \\mon. Some concluding remarks are shown in Sect.~\\ref{ConclusionsSection}. ",
        "conclusions": ""
    },
    "0402/astro-ph0402126_arXiv.txt": {
        "abstract": "{We present the near-infrared and optical properties of the peculiar galaxy ESO~235-G58, which resembles a late-type ringed barred spiral seen close to face-on. However, the apparent bar of ESO~235-G58 is in reality an edge-on disk galaxy of relatively low luminosity.  We have analyzed the light and color distributions of ESO~235-G58 in the NIR and optical bands and compared them with the typical properties observed for other morphological galaxy types, including polar ring galaxies. Similar properties are observed for ESO~235-G58, polar ring galaxies, and spiral galaxies, which leads us to conclude that this peculiar system is a {\\it polar-ring-related} galaxy, characterized by a low inclined ring/disk structure, as pointed out by Buta \\& Crocker in an earlier study, rather than a barred galaxy. ",
        "introduction": "The southern galaxy ESO~235-G58 is classified in RC3 (de Vaucouleurs et al. 1991) as an SB(rs)d spiral based on its appearance on the ESO/SRC-J Southern Sky Survey. It is characterized by a quite elongated bar-like structure, for which the total extension is $\\sim 40''$ and the position angle (P.A.) of its major axis is $106^\\circ$. This apparent bar is encircled by a weak inner pseudo-ring and two faint outer arms, which reach a distance of about $1'$ from the center (Fig.\\ref{eso235B}). Using $H_0=70$ km s$^{-1}$ Mpc$^{-1}$, the galaxy has a distance of about $67$ Mpc. In this early classification the galaxy disk is seen nearly face-on, with a pseudo-ring at about 1' from the center, encircling an inner bar.  Closer inspection of ESO~235-G58 (Buta 1995), in the {\\it Catalogue of Southern Ringed Galaxies}, led to the suspicion that the apparent central bar is in reality an edge-on disk. With new optical data, Buta and Crocker (1993) showed that ESO~235-G58 is probably not a ringed barred spiral of late Hubble type, because the central elongated structure shows a nearly linear dust lane which splits the nuclear region in a way which is typical of edge-on disks (see Fig.~\\ref{eso235B}). This implies that the equatorial plane of the central structure is at a different angle to the line of sight than the outer ring.  The analysis of these optical data (Buta and Crocker 1993) suggested that ESO~235-G58 is composed of a low luminosity edge-on disk, with an equatorial dust-lane, and an outer ring/spiral at very large radii, seen at a lower inclination. Therefore, ESO~235-G58 may be a good candidate for Polar-Ring-related objects. Polar Ring galaxies (PRGs) are usually recognized because the two components (the host galaxy and ring) are nearly edge-on (Whitmore et al. 1990). It is more difficult to recognize cases where the inner component is edge-on, while the ring is more face-on: cases like that must exist and ESO235-G58 may be such kind of an object.  The main goal of the present work is to provide accurate optical and near-infrared (NIR) photometry of ESO~235-G58, and to compare the main properties of this galaxy with those typical of other morphological galaxy types, including PRGs.  NIR photometry is necessary to minimize the dust absorption which strongly affects the starlight distribution in the central galaxy. In addition, the study of optical and NIR integrated colors will provide information about the age and metallicity of the dominant stellar population in the different components of the system.  In this paper we present new NIR observations, obtained for ESO~235-G58 in J, H and Kn bands, and have applied the same procedures adopted to study PRGs by Iodice et al. (2002a, 2002b, 2002c). As shown by Iodice and collaborators, this kind of analysis allows us to obtain quantitative morphology of the main components in ESO~235-G58, in order to rigorously classify this object, and gain insight on its star formation history.  Observations and data reduction are presented in Sec.\\ref{obs}; the morphology, light and color distribution of the two components (host galaxy and ring) are discussed in Sec.\\ref{eso235_morph} and Sec.\\ref{eso235_phot}. In Sec.\\ref{eso235_dust} we describe the dust properties and how they compare with the typical properties in other galaxies and the Milky Way. In Sec.\\ref{eso235_col} the integrated colors derived for the central galaxy and ring are presented, and in Sec.\\ref{age_eso235} we give an age estimate of the two components. The two-dimensional model of the central galaxy light distribution is discussed in Sec.\\ref{eso235_g}, and in Sec.\\ref{eso235_r} we discuss the properties of the ring light distribution. A summary of the main results and conclusions is presented in Sec.\\ref{eso235_sum}. ",
        "conclusions": "\\label{eso235_sum}  We have discussed the NIR and optical properties of the peculiar galaxy ESO~235-G58: an accurate photometric study in the optical bands led Buta \\& Crocker (1993) to classify ESO~235-G58 as an interacting system related to the polar ring galaxies. For this peculiar object, we have analyzed the light and color distribution in the NIR and optical bands and we have compared them with the typical properties observed for other morphological galaxy types, including polar ring galaxies (described in Iodice et al. 2002a, 2002b, 2002c). The main results of this analysis are the following: \\begin{enumerate} \\item the P.A. and ellipticity profiles are quite different from those observed for nearly face-on barred galaxies, but they are, on the other hand, very similar to the typical P.A. and ellipticity profiles observed for edge-on disk galaxies; \\item the high-frequency residual images confirm that the central galaxy in ESO~235-G58 has an edge-on disk; \\item the analysis of the color and light distribution in the central component strongly suggests that ESO~235-G58 is not a nearly face-on barred galaxy and that it shows many similarities to edge-on spiral galaxies and to the host galaxy in PRGs; \\item the outer ring structure is almost undetectable in the NIR bands: it is characterized by very blue colors which are similar to those of dwarf irregular galaxies and polar rings; \\item the central galaxy in ESO~235-G58 last formed a significant number of stars between 1 to 3 Gyrs ago, very similar both to the host galaxy in PRGs and for spiral galaxies. The last episode of stellar burst in the ring structure may be as recent as $10^8 yr$, since its colors are comparable to those of the polar component in NGC~4650A and ESO~603-G21. \\end{enumerate}  Furthermore, this peculiar galaxy is characterized by a high amount of neutral hydrogen, mainly associated with the ring component, as is also usually observed in many PRGs (van Gorkom et al. 1987; Arnaboldi et al. 1997; van Driel et al. 2000; 2002): the total HI mass is about $3 \\times 10^9 M_{\\odot}$ (Buta \\& Crocker 1993; van Driel et al. 2000; 2002). We derive a total mass for the stellar component of about $4 \\times 10^7 M_{\\odot}$, assuming a mass-to-light ratio $M/L=0.04$ (see Sec.\\ref{age_eso235}).  The total mass for the stellar component of the central galaxy is about $2.1 \\times 10^9 M_{\\odot}$, assuming $M/L=1.3$ (see Sec.\\ref{age_eso235}).  The total {\\it baryonic mass} in the ring structure, i.e. stellar plus gaseous component, is then about 1.4 times larger than the total mass of the stellar component in the central galaxy. A value larger than unity for the {\\it total baryonic mass}-to-{\\it stellar mass} ratio is commonly observed for PRGs (see Iodice et al. 2002a, 2002c).  The low inclination angle between the central galaxy and ring ($40^\\circ$, by Buta \\& Crocker 1993), which is unusual in PRGs, may suggest that the interaction is recent and the ring may not yet have reached a stable configuration or it may be a transient structure. Connecting ESO235-G58 to PRGs is important to know this. The ring's low inclination with respect to the equatorial plane of the central galaxy makes ESO~235-G58 very similar to another {\\it polar-ring-related} object, {\\it NGC 660} (van Driel et al. 1995). According to the latest N-body simulations by Bournaud \\& Combes (2003), such objects are likely to be formed through an accretion mechanism, where gas is accreted by the host galaxy from a gas-rich donor. Bournaud \\& Combes' scenario can account for both nearly polar structures, and more shallowly-inclined rings, whose radial extension $\\Delta R /\\bar{R}$ can be up to $55\\%$.  These predictions are consistent with the observed values for these parameters in ESO~235-G58 (see Sec.\\ref{intro} and Sec.\\ref{eso235_r}).  The stability of such inclined structures depends on the ring mass and the dark matter distribution; for a ring at low inclination with respect to the equatorial plane whose mass is 30\\% of the visible total mass in the system, those numerical simulations (Bournaud \\& Combes 2003, see Section 6.3) predict that the ring is unstable, but will not disrupt within a timescale of 2.5 Gyrs in a nearly spherical dark halo. In the case of ESO~235-G58, the ring mass (stars plus gas) is about $60\\%$ of the visible total mass in the system, and information on the dark halo shape can be derived from the position of ESO~235-G58 in the $L_B - Log(W_{20})$ plane (where $L_B$ is the total luminosity of the whole system), as described by Iodice et al. (2003). The $L_B,\\,Log(W_{20})$ values for ESO~235-G58 fall on the average Tully-Fisher relation for bright disk galaxies (see Iodice et al. 2003, Fig.~2): this implies that the dark halo in ESO~235-G58 may be nearly spherical.  Thus the ring will probably persist for still longer than 2.5 Gyrs.  The present work leads us to conclude that ESO~235-G58 is a {\\it polar-ring-related} galaxy, characterized by a low inclined ring/disk structure. Such inclined structures are rarely observed, because they are less stable than those nearly polar and they can be easily confused with barred galaxies, as it seems to have happened in the case of ESO~235-G58. This peculiar galaxy provides us with an intriguing face-on view of the low-inclined disk component of a PRG-related object. Kinematic observations are urgently needed to further understand the structure of this enigmatic object. The analysis of high resolution observations would provide further insights into the effects of the interaction on both the inner component and the inclined disk."
    },
    "0402/astro-ph0402310_arXiv.txt": {
        "abstract": "{We explored the regions within a radius of 25\\arcsec~around 473 nearby, low-metallicity G- to M-type stars using $(VR)I$ optical filters and small-aperture telescopes. About 10\\%~of the sample was searched up to angular separations of 90\\arcsec. We applied photometric and astrometric techniques to detect true physical companions to the targets. The great majority of the sample stars was drawn from the Carney--Latham surveys; their metallicities range from roughly solar to [Fe/H]\\,=\\,$-3.5$\\,dex. Our $I$-band photometric survey detected objects that are between 0 and 5 mag fainter (completeness) than the target stars; the maximum dynamical range of our exploration is 9\\,mag. We also investigated the literature and inspected images from the Digitized Sky Surveys to complete our search. By combining photometric and proper motion measurements, we retrieved 29 previously known companions, and identified 13 new proper motion companions. Near-infrared 2MASS photometry is provided for the great majority of them. Low-resolution optical spectroscopy (386--1000\\,nm) was obtained for eight of the new companion stars. These spectroscopic data confirm them as cool, late-type, metal-depleted dwarfs, with spectral classes from esdK7 to sdM3. After comparison with low-metallicity evolutionary models, we estimate the masses of the proper motion companion stars to be in the range 0.5--0.1\\,$M_{\\odot}$. They are orbiting their primary stars at projected separations between $\\sim$32 and $\\sim$57000\\,AU. These orbital sizes are very similar to those of solar-metallicity stars of the same spectral types. Our results indicate that about 15\\%~of the metal-poor stars have stellar companions at large orbits, which is in agreement with the binary fraction observed among main sequence G- to M-type stars and T\\,Tauri stars. ",
        "introduction": "Halo subdwarfs are metal-deficient ([Fe/H]\\,$\\le$\\,$-1$\\,dex), high-velocity ($v_{\\rm tan}$\\,$\\ge$\\,200\\,km\\,s$^{-1}$) stars. They belong to the oldest known galactic population and are subluminous with respect to main-sequence stars of the same mass. Traditionally, they have been identified using proper motion surveys (e.g., Giclas \\cite{giclas71}) and, more recently, objective prism plates (e.g., Beers et al$.$ \\cite{beers85}). Such surveys are magnitude-limited and consequently biased towards intrinsically brighter stars. Only recently a subdwarf sequence of very-low mass halo subdwarfs has been identified with a cutoff at M($V$)\\,=\\,14.5 and $(V-I)$\\,=\\,2.8\\,mag (Monet et al$.$ \\cite{monet92}). The bolometric luminosities, effective temperatures and metallicities of these stars, and the faint cutoff, are an important constraint to stellar structural and evolutionary models. Gizis \\& Reid (\\cite{gizis97}) have stressed the importance of wide binary systems for checking the metallicity scale of low-mass dwarfs.  In 1991 we started a {\\sc ccd}-based imaging survey aimed at finding wide low-mass companions to halo subdwarfs. Wide binary systems are like small clusters because they offer us the possibility of studying two stars of different mass, but with the same age, distance and chemical composition. Early results of our survey were reported in Mart\\'\\i n \\& Rebolo (\\cite{martin92}) and Mart\\'\\i n et al$.$ (\\cite{martin95}).  In this paper, we present the full survey, including the discovery of companions to 13 metal-deficient dwarfs. Spectroscopic observations of the low-mass secondaries have been obtained for eight of them. We derive spectral types for these new objects ranging from esdK7 to sdM3.   \\begin{figure} \\centering \\includegraphics[width=8.5cm]{met.ps} \\caption{The metallicity distribution of our entire sample is plotted as a thin line. The metallicity distribution of the stars with wide companions is displayed with a thick line. The bin width represents a change of 0.25\\,dex in the metal logarithmic abundance.} \\label{metal} \\end{figure} ",
        "conclusions": "We explored the nearby regions around 473 low-metallicity G- to M-type stars searching for low-mass stellar companions that are orbiting their primary stars at wide separations (typically $\\ge$30\\,AU). The great majority of the target stars, with high proper motions and metallicities in the range [Fe/H]\\,=\\,[$-$3.5, 0.0], were selected from Carney et al$.$ (\\cite{carney94}) and Laird et al$.$ (\\cite{laird88}). All of them were imaged in the $I$-band with a 1-m class telescope. Optical $V$ and $R$ data were also collected for a large number of the targets. The dynamical range of the {\\sc ccd} detectors allowed us to detect companions up to 5 magnitudes fainter (completeness) than the target star. The physical link between a candidate and its primary star was assessed by means of photometric and proper motion measurements. We also searched the literature and various archives to complete our survey. We identified 13 new proper motion companions and retrieved 29 previously known companions. This suggests that about one third of the low-mass companions are missing in previous proper motion searches.  2MASS $JHK_s$ photometry is provided for a total of 39 companions out of 42. On the basis of optical and near-infrared colors, the 13 new companions are dwarfs. Two out of the 29 previously known companions are white dwarfs. We produced optical and near-infrared color-magnitude and color-color diagrams, and overplotted low-metallicity models from Baraffe et al$.$ (\\cite{baraffe97}). The agreement between observations and theory is reasonable, indicating that dwarf companions have masses in the range 0.1--0.5\\,$M_{\\odot}$.  Low-resolution optical spectra from 386 to 1000\\,nm were obtained for 8 of the new proper motion companions, for which we derived subdwarf spectral types esdK7.2--sdM3.0 (error bar of half a subclass). The spectra of significantly metal-poor companions are dominated by strong MgH and CaH molecular absorptions. At the resolution of our data, spectra appear featureless redward of 800\\,nm, except for the CaII triplet that remains detectable down to esdM3. The molecular features between 600 and 800\\,nm are quite sensitive to temperature changes in contrast to the blue spectra.  Proper motion pairs have projected separations between $\\sim$30 and $\\sim$57000\\,AU. Very wide companions are also identified among the most metal-depleted stars. These orbital sizes are similar to those of solar-metallicity binaries and multiple stars. After correcting for the effect of increasing distance to the most metal-deficient stars in our survey, we determined that 13--15\\%~of the low-metallicity G- to M-type stars harbor wide companions. This binary frequency is very similar to the binary fraction observed among stars of the solar vicinity and T\\,Tauri stars of star-forming regions, suggesting that metallicity is not a key parameter in the stellar formation of wide double and multiple systems."
    },
    "0402/astro-ph0402476_arXiv.txt": {
        "abstract": "We present optical spectra of the ionized `Outer Ejecta' of $\\eta$ Carinae that reveal differences in chemical composition at various positions.  In particular, young condensations just outside the dusty Homunculus Nebula show strong nitrogen lines and little or no oxygen --- but farther away, nitrogen lines weaken and oxygen lines become stronger.  The observed variations in the apparent N/O ratio may signify either that the various blobs were ejected with different abundances, or more likely, that the more distant condensations are interacting with normal-composition material.  The second hypothesis is supported by various other clues involving kinematics and X-ray emission, and would suggest that $\\eta$ Car is enveloped in a ``cocoon'' deposited by previous stellar-wind mass loss.  In particular, all emission features where we detect strong oxygen lines are coincident with or outside the soft X-ray shell.  In either case, the observed abundance variations suggest that $\\eta$ Car's ejection of nitrogen-rich material is a {\\it recent} phenomenon --- taking place in just the last few thousand years.  Thus, $\\eta$ Carinae may be at a critical stage of evolution when ashes of the CNO cycle have just appeared at its surface.  Finally, these spectra reveal some extremely fast nitrogen-rich material, with Doppler velocities up to 3200 km s$^{-1}$, and actual space velocities that may be much higher. This is the fastest material yet seen in $\\eta$ Car's nebula, but with unknown projection angles its age is uncertain. ",
        "introduction": "In the hot interiors of massive stars, the equilibrium CNO cycle converts most of the carbon and oxygen into nitrogen.  Through turbulent and rotational mixing (Maeder 1982), or perhaps by stripping of the outer hydrogen-rich envelope by a stellar-wind, these nitrogen-rich ashes from the core may eventually be observed in a massive star's atmosphere or its ejecta.  One spectacular example of N-rich ejecta in a massive star's circumstellar nebula, in addition to the rings around SN 1987a (Sonneborn et al.\\ 1997; Meaburn et al.\\ 1995), is the `Outer Ejecta' around the evolved massive star $\\eta$ Carinae.  Outside its dusty bipolar reflection nebula known as the `Homunculus', which was ejected in the mid-19th century, $\\eta$ Carinae is surrounded by a complex aggregate of ionized gas condensations known collectively as the `Outer ejecta', the `Outer Shell', or the `Outer Condensations' (e.g., Walborn 1976; Thackeray 1950; Meaburn et al.\\ 1996).  Since these ejecta are ionized, as opposed to mostly neutral and dusty like the Homunculus, they provide a practical way to estimate the chemical abundances of some material recently ejected by $\\eta$ Carinae.  Some of the more prominent nebular condensations in the outer ejecta have names that are identified in the {\\it HST}/WFPC2 image in Figure 1 (see Morse 1999; Morse et al.\\ 1998), following the nomenclature of Walborn (1976).  Proper motions suggest that some of these condensations may predate $\\eta$ Car's famous 19th century eruption by a few hundred years or more (Walborn et al. 1978, Walborn \\& Blanco 1988), while some ejecta closer to the Homunculus originated within a few decades of that eruption (Morse et al.\\ 2001).  Bright ionized condensations in the outer ejecta coincide with a soft X-ray emitting shell (Seward et al.\\ 2001; Chlebowski et al.\\ 1984), suggesting that they are excited predominantly by shocks, rather than photoionized.  Thus, young fast ejecta may be plowing into older material.  Davidson et al.\\ (1982, 1986) examined the UV and optical spectrum of the brightest condensation and the brightest X-ray source in the outer ejecta (the `S Condensation') and found it to be extremely nitrogen rich, being almost completely devoid of any oxygen or carbon lines.  Thus, ashes of the CNO cycle have been exposed at the star's surface, confirming that $\\eta$ Car is indeed an evolved massive star. Dufour (1989) and Dufour et al.\\ (1997) have also discussed the physical properties and abundances in the S Condensation and some similar nearby ejecta.  However, no suitable spectra for examining abundances of other condensations in the Outer Ejecta have been published; consequently, it is usually assumed that all of the Outer Ejecta are nitrogen rich.  Here we present long-slit optical spectra for several positions in $\\eta$ Car's Outer Ejecta.  We find that the chemical abundances are {\\it not} uniform, and we discuss interpretations and evolutionary implications of this conclusion.  We present our spectroscopic observations in \\S 2 and discuss the data in \\S 3, including a deductive chemical abundance analysis.  In \\S 4 we discuss implications of the non-uniform abundances in the young ejecta around $\\eta$ Car, as well as some interesting details of the observed kinematics. ",
        "conclusions": "\\subsection{Possible Interpretations}  Our observations suggest inhomogeneous abundances in the Outer Ejecta of $\\eta$ Carinae.  Specifically, there is a strong change in the apparent N/O ratio: while ejecta immediately outside the Homunculus are nitrogen rich and severely oxygen depleted, more distant material appears to have normal abundances.  If it is indeed true, how shall we interpret this apparent chemical abundance gradient?  Two different scenarios seem like obvious possibilities:  1.  Dense condensations ejected by $\\eta$ Carinae over the past few thousand years may have become progressively more enriched with nitrogen and depleted of oxygen as ashes from the CNO cycle appeared at the surface of the star, or as mass loss ate deeper into the processed material inside the star.  2.  All of the dense condensations that make up the Outer Ejecta may indeed be nitrogen rich (perhaps in varying degrees), but the older material farther from the star may be interacting with the normal composition material deposited by previous stellar-wind mass loss.  This would imply that $\\eta$ Car's Outer Ejecta are expanding into a cocoon or halo created by the loss of the star's original hydrogen-rich stellar envelope, as depicted in Figure 5.  Observed kinematics may help distinguish between these possibilities. For example, proper motion measurements may be able to determine if the ejecta are expanding freely or are being decelerated, and ejecta with different chemical abundances may show discrepant Doppler velocities.  Higher-quality data are needed for conclusive answers (high-dispersion spectroscopy of oxygen lines might be particularly useful), but some more immediate clues are discussed below. Proper-motion measurements in the S Ridge (Walborn \\& Blanco 1988; Morse et al.\\ 2001) show a very large scatter in the projected velocities (about 5 times larger than the Homunculus), suggesting interactions between slower condensations and faster ejecta or a wind.    \\subsection{Clues from Kinematics}  Of the two possibilities listed above, the second appears to be supported more closely by the observations. A prominent soft X-ray shell surrounds $\\eta$ Car, and the shock heating of this shell arises as fast ejecta overtake slower material.  Figure 3 shows contours of the soft X-ray emission observed by {\\it Chandra} ACIS-I (see Seward et al.\\ 2001), superposed on an optical [N~{\\sc ii}] image of the Outer Ejecta taken in June 1999 with {\\it HST}/WFPC2.  Condensations that show relatively bright oxygen lines in their spectra in Figure 2 (E5, W2, W Edge) coincide with or lie outside the soft X-ray shell, whereas oxygen-poor ejecta (S Ridge, NN Jet) are inside it.  The S Condensation, S Ridge, and NN Jet may comprise an `inner shell' (Meaburn et al.\\ 1996) that may be distinct from more distant ejecta (note, however, that the S Condensation and NN Jet appear to have been ejected in the Great Eruption, while the S Ridge originated somewhat earlier; Morse et al.\\ 2001).  Thus, if one wishes to estimate chemical abundances of $\\eta$ Car pertaining to its current evolutionary state, one must study the ionized ejecta of this `inner shell' which is not yet contaminated by swept-up material.  Examining long-slit spectra like the detail of [N~{\\sc ii}] and H$\\alpha$ lines shown in Figure 4 provides some additional clues. While the background H~{\\sc ii} region emission has been subtracted carefully, Figure 4 shows some faint H$\\alpha$ emission near the systemic velocity within $\\sim$30$\\arcsec$ of $\\eta$ Car.  This implies that the Outer Ejecta are embedded in a cocoon or halo of H-rich material.  Even with a careful subtraction of the spectrum on either side of a given condensation like E5, there is still some enhanced low-velocity H$\\alpha$ emission at the position of the condensation that contaminates the spectrum in Figure 2.  This may indicate some low-velocity H-rich gas associated with the E5 condensation itself; the oxygen lines in the E5 condensation also appear to have lower velocities than the nitrogen lines.  Obviously our analysis is hampered by low spectral resolution, so it is difficult to confirm any direct relation between kinematics and chemical abundances.  However, Meaburn et al. (1996) presented echelle spectra of H$\\alpha$ and the [N~{\\sc ii}] lines for several condensations and showed a clear trend: the [N~{\\sc ii}]/H$\\alpha$ ratio was stronger in fast ejecta just outside the Homunculus, and dropped by a factor of $\\sim$2 in the slower ejecta farther from the star.  Thus, even at high dispersion where line profiles are resolved, line intensities may be qualitatively consistent with an abundance gradient.    \\subsection{Extremely Fast Ejecta}  Figure 4$a$ also reveals some extremely fast nitrogen-rich ejecta not previously reported.  At one slit position northeast of the star, we see blueshifted emission up to about $-$3200 km s$^{-1}$ and redshifted emission as fast as $+$2000 km s$^{-1}$.  The extremely fast blueshifted material is faint, and was not seen in previous investigations (Meaburn et al.\\ 1987, 1993, 1996; Dufour 1989; Weis et al.\\ 2001).  Dufour (1989) detected a redshifted emission feature at $+$2150 km s$^{-1}$, but no blueshifted features faster than $-$1000 km s$^{-1}$ were reported.  The fastest blueshifted emission in Figure 4$a$ is particularly interesting.  First, this material is nitrogen-rich and the Doppler shifts up to $-$3200 km s$^{-1}$ are probably correct, since the two nitrogen lines and perhaps even H$\\alpha$ can be seen offset from one another by the expected amount (these features are identified as ``EFE'' for ``extremely fast ejecta'' in Figure 4).  Coincidentally, this fast blueshifted emission extends spatially away from the Homunculus as far as the position of the E Condensation 5 (roughly $-$25$\\arcsec$ along the slit in Figure 4), where it seems to end abruptly.  {\\it Is this fast material colliding with the E Condensation?}  If so, there are interesting implications for the geometry.  The E Condensation 5 has a trajectory tilted out of the plane of the sky by $\\alpha \\approx 13\\arcdeg$, since it has a Doppler shift of $-$140 km s$^{-1}$ (Meaburn et al.\\ 1996) and a tangential speed from proper motions of $\\sim$600 km s$^{-1}$ (Walborn et al.\\ 1978).  If the same value of $\\alpha$ applies to the EFE, the true space velocity away from $\\eta$ Car would be well above 10$^4$ km s$^{-1}$.  Such astonishingly high speeds would be unusual even for young supernova remnants (e.g., Fesen et al.\\ 1988).  Though $\\alpha$ is uncertain, the true space velocity is likely to be well above 3200 km s$^{-1}$, since the material is certainly not coming directly toward us -- the fastest material is seen $\\sim$25$\\arcsec$ away from the star. {\\it With speeds well above 3200 km s$^{-1}$, this is the fastest nebular material yet detected that is associated with $\\eta$ Car.}  Could this `EFE' signify a shock from the Great Eruption that is now causing the X-ray shell around $\\eta$ Car?  Obviously, better spectroscopic data are desirable to constrain the age of this interesting fast ejecta. Unfortunately, the proper motion of this emission cannot be measured with existing {\\it HST} data because it is Doppler shifted out of narrow imaging filters centered on H$\\alpha$ or [N~{\\sc ii}].  Regardless of the age, it is likely that this fast material was ejected by $\\eta$ Car itself (i.e. the evolved primary in the putative binary system that was responsible for the Great Eruption), because it is nitrogen rich, with [N~{\\sc ii}] $\\lambda$6548 $\\div$ H$\\alpha$ more than 6.  This indicates that under some circumstances, the primary star is indeed capable of ejecting material at speeds sufficient to account for the hard X-ray emission that varies with $\\eta$ Car's 5.5 year cycle (Pittard \\& Corcoran 2002).  Such high speeds are not seen in spectroscopy of the primary star's stellar wind, even at the poles where the outflow speed is highest (Smith et al.\\ 2003a).  If this fast material is sweeping-out a cavity and causing the soft X-ray emission shell in Figure 3, it may help to explain why there is no significant X-ray emission coming from the NN Jet and some other Outer Ejecta to the north and east of the Homunculus -- i.e. even though the NN Jet is moving very fast, it is plowing through material that is already moving outward at high speed, so the relative shock velocity is insufficient for X-ray production, and instead shows a relatively low-ionization optical spectrum.  This is analogous to the shock excitation in Herbig-Haro jets (Hartigan et al.\\ 1990; Morse et al.\\ 1994) and may explain why we see [N~{\\sc ii}] emission even though the observed speeds imply that the shocked material should be non-radiative.  We did not detect EFE south of the star in the RC Spec slit aperture offset to the southwest of $\\eta$ Car.  This absence is particularly interesting, since that position coincides with a distinct gap in the soft X-ray shell around $\\eta$ Car (see Figure 3).  Finally, the `EFE' in Figure 4 may help explain the peculiar morphology of the E Condensations (see Figs.\\ 1 and 3), which show teardrop-like shapes that sweep away from the central star.  If the E Condensations are dense knots that are being overtaken by a much faster and more tenuous outflowing wind or blast wave, their structure may arise from Rayleigh-Taylor instabilities.  The detailed structure in the E Condensations resembles some Rayleigh-Taylor instabilities in the supernova remnant Cas A, for example (Fesen et al.\\ 2001), or simulations of wind-cloud interactions (e.g., Klein et al.\\ 1994). Note that the E Condensations do coincide with a prominent feature in the soft X-ray shell (Figure 3).   \\subsection{Sudden Chemical Enrichment?}  Of the two possible scenarios to explain the observed N and O line intensities --- chemical abundance gradients in the dense knots vs.\\ modified abundances through swept-up material --- the second seems more likely for reasons described above.  In either case, however, the observed abundance gradient is significant, because it indicates that $\\eta$ Car's ejection of nitrogen-rich ashes of the CNO cycle is a {\\it recent} phenomenon, occurring in just the past few thousand years.  (The oldest of the outer condensations measured by Walborn et al.\\ (1978) have proper motions indicating ages of $\\sim$10$^3$ years.)  We have discussed several interesting aspects of the observed kinematics, but the main conclusion of this paper is as follows:  {\\it Our observations suggest that diffuse gas immediately outside the nitrogen-rich condensations seen in images of $\\eta$ Car --- the material they are running into --- has not been significantly processed through the CNO burning cycle.}  The surrounding cocoon of normal composition material that the outer ejecta are now plowing through (see Figure 5) probably corresponds to part of the star's original H-rich stellar envelope.  Lamers et al.\\ (2001) have inferred a similar type of abundance gradient or rapid enrichment deduced from N/O ratios in a sample of other LBV nebulae. Since the Outer Ejecta are so young ($\\la$10$^3$ years), this conjecture has interesting implications for stellar evolution theories for the most massive stars.  For example, stellar evolution models for stars above 50 to 60 M$_{\\odot}$ predict that internal turbulent mixing timescales are shorter than mass-loss timescales on the main sequence; a star with an initial mass of 120 M$_{\\odot}$ has a mixing timescale of only 1.6 Myr compared to a mass-loss timescale of 2.4 Myr, both of which are shorter than the time spent on the main sequence (Maeder 1982).  Rotation tends to enhance the mixing (Maeder \\& Meynet 2002; Maeder 2002), and there are indications that $\\eta$ Car has significant rotation (Smith et al.\\ 2003a).  Thus, we should expect the cores and envelopes of very massive stars to be well-mixed and to evolve quasi-homogeneously.  Instead, our observations suggest that $\\eta$ Car's outer layers were blown off the star before turbulent mixing was able to transport nitrogen-rich CNO products to the star's surface.  This might imply that in very massive stars with initial masses above 100 M$_{\\odot}$, the mass-loss timescale on the main sequence is shorter than the mixing timescale between the core and outer layers, shorter than the nuclear-burning timescale for the CNO cycle (e.g., Appenzeller 1970), or that violent sporadic mass-loss events like $\\eta$ Car's Great Eruption play a central role in a very massive star's evolution off the main sequence.  The cocoon around $\\eta$ Car may have been created during post-main sequence evolution immediately before its current tenure as an extreme luminous blue variable, consistent with the idea that $\\eta$ Car may be a rare example of a post-WNL type star (Walborn 1989; see also Langer et al.\\ 1994; Crowther et al.\\ 1995).   The apparent abundance variations we observe in $\\eta$ Car's Outer Ejecta provide grounds for interesting though inconclusive speculation regarding its current evolutionary state: Is there a link between $\\eta$ Car's fundamental instability and the critical point in the star's evolution when CNO products are first exposed at the surface?  Are catastrophic mass-loss events like the Great Eruption responsible for removing the outer H-rich envelope to expose the CNO ashes, or do they occur as a result of it (i.e. from a change in opacity)?  The Great Eruption ejected several solar masses of material from the star (Mitchell \\& Robinson 1978; Hackwell et al.\\ 1986; Cox et al.\\ 1995; Smith et al.\\ 1998, 2003b), and therefore removed a large fraction of the star's outer radius.  For example, models by Guzik et al.\\ (1999) predict extremely tenuous outer layers for a star like $\\eta$ Car, with 95\\% of the radius containing less than 1\\% of the total mass.  It is interesting to speculate that the Great Eruption itself, or perhaps a previous similar event, may have been the trigger that first released the CNO ashes from the star.  Walborn (1976) has already compared $\\eta$ Car's Outer Ejecta to the nitrogen-rich ``quasi stationary flocculi'' in the supernova remnant Cas A (Baade \\& Minkowski 1954; van den Bergh 1971; Chevalier \\& Kirshner 1978; Fesen et al.\\ 1987, 2001); these existed in the presupernova circumstellar environment and were probably ejected shortly before the progenitor star finally exploded.  The corresponding implications of the nitrogen-rich Outer Ejecta for $\\eta$ Car's evolutionary state and its near future are provocative."
    },
    "0402/astro-ph0402195_arXiv.txt": {
        "abstract": "We have performed the first measures of mass accretion rates in the core of the Orion Nebula Cluster. Four adjacent fields centered on the Trapezium stars have been imaged in the U- and B-bands using the Wide Field Planetary Camera~2 on board the {\\sl Hubble Space Telescope}. We obtained photometry for 91 stars in the U-band (F336W) and 71 stars in the B-band (F439W). The \\WFPCTwo\\ archive was also searched to obtain complementary V-band (F547M) and I-band (F791W) photometry. In this paper we focus our attention on a group of 40 stars with known spectral types and complete UBVI \\WFPCTwo\\ photometry. We locate each star on the HR diagram considering both the standard ISM reddening law with $R_V=3.1$ and the ``anomalous'' reddening law with $R_V=5.5$ more appropriate for the Orion Nebula. Then we derive the stellar masses and ages by comparing with the evolutionary tracks and isochrones calculated by D'Antona \\& Mazzitelli and Palla \\& Stahler. Approximately three quarters of the sources show excess luminosity in the U-band, that we attribute to mass accretion. The known correlation between the U-band excess and the total accretion luminosity, recalibrated for our photometric system, allows us to estimate the accretion rates, which are all found to be in the range $10^{-8}-10^{-12}$M$_\\odot$~yr$^{-1}$. For stars older than 1~Myr there is some evidence of a relation between mass accretion rates and stellar age. Overall, mass accretion rates appear lower than those measured by other authors in the Orion flanking fields or in Taurus-Auriga. Mass accretion rates remain low even in the vicinity of the $10^{-5}$~\\Msunyr\\, birth~line of Palla \\& Stahler, suggesting that in the core of the Trapezium cluster disk accretion has been recently depressed by an external mechanism. We suggest that the UV radiation generated by the Trapezium OB stars, responsible for the disk evaporation, may also cause the drop of the mass accretion rate. In this scenario, low-mass stars may terminate their pre-main sequence evolution with masses lower than those they would have reached if disk accretion could have proceeded undisturbed until the final disk consumption. In OB associations the low-mass end of the Initial Mass Function  may therefore be affected by the rapid evolution of the most massive cluster's stars, causing a  a surplus of ``accretion aborted\" very low-mass stars and brown dwarfs, and a deficit of intermediate mass stars. This trend is in agreement with recent observations of the IMF in the Trapezium cluster. ",
        "introduction": "One of the most relevant processes in the early stellar evolution is the interaction between young pre-main-sequence stars and their circumstellar disks. With the exception of the very initial phases of star formation, during which protostars may directly accrete low angular momentum material radially infalling from their parental cloud, most of the stellar mass build-up occurs through disk accretion (e.g.\\ Tereby, Shu, \\& Cassen 1984). Disk accretion regulates the final stellar mass removing angular momentum via viscous dissipation and feeding conspicuous mass loss through collimated outflows (e.g.\\ Hartmann 1998). There is evidence that the disk accretion process becomes less and less significant as the young star ages (Hartmann et~al.\\ 1998), and, for low mass stars ($M<1M_\\odot$), accretion must terminate within the disk lifetime (6--10~Myr, Haisch et~al.\\ 2001; Skrutskie et~al.\\ 1990), well before reaching the main sequence. However, a quantitative understanding of the evolution of the accretion process is still missing. In principle, mass accretion can be studied in a variety of ways. The infrared disk emission produced by internal viscous dissipation is potentially an important diagnostic tool, but in most disks direct irradiation of the surface from the central star is a stronger heat source (Kenyon \\& Hartmann 1987). A less ambiguous indicator is the recombination radiation produced when the disk material falls onto the stellar surface (Bertout 1989), through a boundary layer or, more probably, through accretion columns in the stellar magnetosphere (K\\\"onigl 1991; Camenzind 1990, Shu et~al.\\ 1994). The hot gas continuum adds up to the stellar continuum and alters the depths of photospheric lines, the so-called ``veiling.'' From the amount by which stellar absorption features are veiled, one can make quantitative estimates of the relative emission of star and hot continuum, obtaining the integrated accretion luminosity and, therefore, the mass accretion rate (Edwards et~al.\\ 1987; Hartigan et~al.\\ 1991; Bathala et~al.\\ 1996; Hartigan, Edwards \\& Ghandour 1995; Gullbring et~al.\\ 1998). The hot continuum emission becomes easily detectable shortward of the Balmer discontinuity as a characteristic ultraviolet excess (Kuhi 1974). Gullbring et~al.\\ (1998) have shown that the excess radiation in the Johnson U-band is related to the accretion luminosity derived from veiling measures. Calvet \\& Gullbring (1998) have reproduced the relation within the framework of the  magnetically driven accretion column model, showing that the correlation does not depend on the main stellar parameters (mass, effective temperature) at least for spectral types in the range from M3 to K5. The UV excess has been used by Gullbring et~al.\\ (1998, 2000) to determine the mass accretion of stars in the Taurus-Auriga association, by Hartmann et~al.\\ (1998) in Chamaeleon-I, and by Rebull et~al.\\ (2000) on the largely unexplored outer regions of the Orion Nebula.  This paper presents new \\HST\\ photometry of the Trapezium cluster, the nearest region of ongoing high-mass star formation. Located in the core of the Orion Nebula, this cluster contains at least 3000 members between $\\sim 40$~M$_\\odot$ and $\\sim 10$~M$_J$ (Lada \\& Lada, 2003). The availability of spectral type classification for $\\approx 1000$ stars (Hillenbrand 1997) provides a unique opportunity to build the largest database of accretion luminosities and mass accretion rate. In this work we present the results obtained for a sample of 40 isolated point sources with known spectral types and accurate \\HST\\ photometry. The observations are discussed in Section~2, whereas in Section~3 we illustrate the criteria leading to the sample selection. Our data analysis strategy is presented in Section~4. In Section~5 we present the results, in particular the location of our sources on the HR diagram and the derived mass accretion rates. Finally, in Section~6 we discuss our findings. Whereas our original goal was to build the first reliable map of the mass accretion rates vs. stellar mass and age, we find evidence of a link between the mass accretion rates and the extreme environmental conditions of the cluster core. This may have profound implications on the final stellar masses, affecting the Initial Mass Function (IMF) of the Cluster's core, especially at the low mass end. ",
        "conclusions": "We have performed the first quantitative analysis of the mass accretion rates in the core of the Trapezium cluster, focusing our attention on 40 unresolved stellar sources with known spectral types and accurate WFPC2 UBVI-band photometry. We have estimated the ultraviolet excess considering two different reddening laws, and derived the stellar parameters and mass accretion rates using the D'Antona \\& Mazzitelli (1998) and Palla \\& Stahler 1999) computations. Approximately 75\\% of the sources show appreciable excess luminosity in the U-band, that we attribute to accretion. We used the known correlation between the U-band excess and the total accretion luminosity (Gullbring et~al.\\ 1998), recalibrated to our photometric system, to estimate the mass accretion rates, all found to be in the range $10^{-8}-10^{-12}\\Msunyr$. We find some evidence for a decrease of $\\dot{M}$ with the stellar age for sources older than 0.1~Myr, but in general the mass accretion rates appear to be lower than those measured in Taurus or in the flanking fields of the Orion Nebula. We suggest that the photoionization generated by the Trapezium OB stars causes a drop of the disk mass accretion rate. Low mass stars therefore conclude their pre-main sequence evolution with lower masses, and the Initial Mass Function turns out to be affected by the rapid evolution of the most massive cluster's stars, with a surplus of ``accretion aborted'' stars or brown dwarfs, and a deficit of intermediate mass stars. This trend is in agreement with recent observations of the IMF in the Trapezium cluster."
    },
    "0402/astro-ph0402530_arXiv.txt": {
        "abstract": "{We report the discovery of a Narrow Line QSO located at about $1.3'$ from the Broad Line Radio Galaxy 3C445. The source, \\wga, although already revealed by ROSAT, has never been optically identified previously. An \\xmm observation of 3C445 has allowed, for the first time, an accurate  X-ray spectral study of \\wga, revealing an ultra-soft spectrum and fast flux variations typical of Narrow Line AGN. The 0.2-10 keV spectrum is well represented by a power law ($\\Gamma=2.5$) plus a black body component ($kT = 117$~eV) absorbed by Galactic \\nh. About $80\\%$ of the X-ray flux ($F_{0.2-10~\\rm keV} \\sim 3\\times10^{-13}$~\\erg) is emitted below $2$~keV. The $0.2-2$~keV flux is observed to decrease by about a factor $1.6$ in about $5000$~s.  The optical observations, triggered by the X-ray study,  confirm the Narrow Line AGN nature of this source. The continuum is blue with typical AGN emission lines, pointing to a redshift $z=0.46$. The full width half maximum of H$_\\beta$  is $2000$~km s$^{-1}$ and the flux ratio [O {\\sc iii}]/H$\\beta=0.21$. The optical luminosity  ($M_R=-23.2$) and the point-like appearance in the optical images identify \\wga as a Narrow Line QSO.  From the optical--UV--X--ray Spectral Energy Distribution we obtain a lower limit of the bolometric luminosity of \\wga ($L_{\\rm bol} \\ge 3\\times 10^{45}$ erg s$^{-1}$) implying, for accretion rates close to the Eddington limit, a black hole mass $M_{\\rm BH} \\ge 2.4 \\times10^{7}$~$M_{\\odot}$. ",
        "introduction": "The Narrow Line (NL) Type I AGN (Seyfert and QSO) are a class of objects which has focused the attention of the scientific community in the last years because of its unusual optical and X-ray properties. Their optical spectra have permitted lines which are slightly broader than the forbidden lines (i.e.full width at half maximum (FWHMH) H$_\\beta\\le 2000$ km $s^{-1}$), a flux ratio O[{\\sc III}]$\\lambda5007/H_{\\beta}<3$ and a prominent Fe{\\sc II} bump (Osterbrock $\\&$ Pogge 1985). In the X-ray range, they exhibit strong flux variability and very steep spectra (Boller et al. 1996, Forster \\& Halpern 1996, Molthagen et al. 1998, Dewangan et al. 2001). The most favored interpretation is that they represent an AGN class with very small black holes and very high accretion rates. Actually, emission lines of highly ionized iron, as expected to be produced by high accreting disks (Nayakshin \\& Kazanas 2001), have been discovered in several NL Seyfert 1 galaxies with ASCA and BeppoSAX (Pounds et al. 1995; Comastri et al. 1998; Turner et al. 1998; Leighly 1999, Gliozzi et al. 2001). Recently, Pounds et al. (2003a,b) have revealed several absorption lines in both \\xmm high and low resolution spectra, interpreted as signatures of a highly ionized wind. As suggested by King $\\&$ Pounds (2003) AGN with very high accretion rates should be able to produce an outflowing photosphere.   In this paper, we present  \\xmm and optical observations of the newly discovered NLQSO located only $1.3'$ far away from the Broad Line Radio Galaxy 3C445. This serendipitous source is included in the the WGA catalog (White, Giommi, Angelini 1994) with the name \\wga and discussed in a ROSAT 3C445 study of Sambruna et al. (1998). However no particular attention has been  paid to it. Its faintness in the ROSAT band (about a factor 3 below 3C445) assured a negligible contamination of the AGN spectrum.  Surprisingly enough, when \\xmm pointed the same field, the serendipitous source, this time clearly spatially resolved, appeared as bright as 3C445 in the soft energy band ($0.2-1$~keV). As we will show in this paper, optical observations of \\wga indicate that the source is actually a quasar at $z=0.46$ with spectral features typical of NL Type I AGN. The presence of this peculiar AGN in the vicinity of 3C445 should be taken into account when the X-ray spectra from previous and very poor spatial resolution satellites are interpreted.  \\begin{figure*}[!ht] \\centering \\includegraphics[scale=0.37]{0724fig1a.ps} \\includegraphics[scale=0.37]{0724fig1b.ps} \\includegraphics[scale=0.37]{0724fig2.ps} \\caption{EPIC-PN ({\\it  left panel}), OM/UVW2 ({\\it  middle  panel}) and Loiano Observatory  V band ({\\it right panel}) images of \\wga. The X-ray and Ultraviolet images are smoothed using a Gaussian filter with $\\sigma=2$~pixels (corresponding to $8''$ and $1''$ in the PN and OM, respectively).} \\label{om} \\end{figure*}   Throughout the paper, luminosities are calculated assuming isotropic emission, a Hubble constant of $H_0=75$~km~s$^{-1}$~Mpc$^{-1}$ and a deceleration parameter of $q_0 = 0.5$. ",
        "conclusions": "Our observations show that the source \\wga, located at about $1.3'$ from the bright radio Galaxy 3C445, is an AGN at redshift $z=0.46$. It is a quasar as attested by  its optical point-like aspect and  its high luminosity both in the optical ($M_{\\rm R}=-23.2$) and X-ray ($L_{0.2-10~\\rm keV}=2.3\\times 10^{44}$~erg s$^{-1}$) bands.   More interesting, both the optical and X-ray spectra indicate that \\wga is a NL QSO. The FWHM of H$_\\beta$ (2000 km s$^{-1}$), the [O {\\sc iii}]/H$_\\beta$ flux ratio (0.21) and the relatively strong FeII emission conform with the NL Type I classification. The X-ray spectrum of \\wga is very soft as observed in most of the NL Type I AGN (Boller et al. 1996, Leighly 1999). The steep power law ($\\Gamma=2.5$) necessary to reproduce the hard X-ray continuum was not sufficient to fit the soft emission below 1 keV. PN data need an extra component that, if we choose to model it with a black body emission, requires a temperature of  $kT=117$~eV.  Similar black body values characterize the NL QSOs studied by ROSAT and ASCA (Forster \\& Halpern 1996, Molthagen et al. 1997, Ulrich et al. 1999, Komossa et al. 2000, Dewangan et a. 2001).  Forster \\& Halpern (1996) discovered a clear correlation between the soft X-ray luminosity and the spectral slope in a large sample of NL Type 1 objects observed with ROSAT.  If we consider a broken power law to model the PN data, the fit formally acceptable ($\\chi^2=52$ for 49 d.o.f.) yields the following spectral parameters: $\\Gamma_ {\\rm soft} = 3.0^{+0.3}_{-0.1}$, $\\Gamma_ {\\rm hard} = 2.3^{+0.4}_{-0.2}$ with an energy break at $1.1$~keV. Our source parameters are in perfect agreement with that correlation. An object with the same intrinsic soft luminosity of \\wga $L_{0.1-2~\\rm keV} \\sim 10^{45}$ erg s$^{-1}$ (assuming $H_0$= 50 and $q_0=0$ to meet the cosmological assumptions of the authors) is expected to have a soft spectral slope of about 3, as actually observed.  Finally, the rapid change of the X-ray flux, rather common in ultra-soft AGN (Boller et al. 1996, Leighly 1999), is another hint in favour of its NL classification.   The ultra soft nature of \\wga becomes much more evident when the optical UV and X-ray data are combined to produce a Spectral Energy Distribution (SED). Unfortunately, neither radio nor infrared counterparts were found for this source.  In Fig.~5, the presence of a very prominent UV--soft--X--ray bump is unequivocal, taking also into account that the X-ray and the UV data are simultaneous. From Fig. 5, it is also evident that the thermal model used to parameterize the X-ray excess can not reproduce the entire UV bump, that is more intense and broad.  On the other hand, it is well known that a black body emission is a very rough assumption.  Plausible interpretations assume inverse-Compton scattering of cold disk photons by a warm/hot atmosphere (Janiuk et al. 2001, Vaughan et al. 2002) or high accretion slim disk characterized by a multi-colour black body emission (Mineshinge et al 2000).  A lower limit of the bolometric luminosity $L_{\\rm bol} \\ge L\\sim 3\\times10^{45}$ erg s$^{-1}$ can be directly estimated from Fig.~\\ref{sed}, using only the optical--UV--X data, implying a $0.2-10$~keV X--ray contribute to the total luminosity less than 7$\\%$. The bulk of the emission occurs between $10^{15}$ and $10^{17}$ Hz. By using simple arguments (uniform accretion at the Eddington limit), we deduce a lower limit to the mass of the central supermassive object of 2$\\times10^{7}$ M$_{\\odot}$.    \\begin{figure}[h] \\label{sed} \\centering \\includegraphics[angle=270, scale=0.35]{0724fig6.ps} \\caption{Optical-UV-X-ray spectral energy distribution of \\wga in the observer frame. A strong UV-soft X-ray bump is clearly visible.} \\end{figure}"
    },
    "0402/hep-ph0402005_arXiv.txt": {
        "abstract": "{In this talk I review briefly theoretical models and ideas on quantum gravity approaches entailing CPT violation. Then, I discuss various phenomenological tests of CPT violation using neutrinos, arguing in favour of their superior sensitivity as compared to that of tests using other particles, such as neutral mesons, or nuclear and atomic physics experiments. I stress the fact that there is no single figure of merit for CPT violation, and that the conclusions on phenomenological sensitivities drawn so far are highly quantum-gravity-model dependent.}  \\normalsize\\baselineskip=15pt ",
        "introduction": "There is a number of fundamental questions that one has to ask before embarking on a study of the phenomenology of CPT violation:    (I) Are there theories which allow CPT breaking?   (II) How (un)likely is it that somebody someday finds CPT violation, and why?  (III) What formalism does one has to adopt? How can we be sure of observing CPT Violation and not something else? our current phenomenology of particle physics is based on CPT invariance.   (IV) There does not seem to be a single ``figure of merit'' for CPT violation. Then how should we compare various ``figures of merit\" of CPT tests (e.g. direct mass measurement between matter and antimatter, $K^0$-${\\overline K}^0$ mass difference a la CPLEAR, Decoherence Effects, Einstein-Podolsky-Rosen (EPR) states in meson factories, neutrino mixing, electron g-2 and cyclotron frequency comparison, neutrino spin-flavour conversion {\\it etc}.)    In some of these questions we shall try to give  some answers in the context of this presentation. Because this is a conference on neutrinos, I will place emphasis on neutrino tests of CPT invariance. As I will argue below, in many instances neutrinos seem to provide at present the best bounds on possible CPT violation. However, I must stress that, precisely because CPT violation is a highly model dependent feature of some approaches to quantum gravity (QG), there may be models in which the sensitivity of other experiments on CPT violation, such as astrophysical experiments, is superior to that of current neutrino experiments.   My talk will focus on the following three major issues:  {\\bf (a)} \\underline{WHAT IS CPT SYMMETRY}: I will give a definition of what we mean by CPT invariance, and under what conditions this invariance holds.  {\\bf (b)} \\underline{WHY CPT VIOLATION ?}: Currently there are various {\\it Quantum Gravity Models} which may {\\it violate} Lorentz symmetry and/or quantum coherence (unitarity {\\em etc}), and through this CPT symmetry:  (i) space-time foam (local field theories, non-critical strings {\\em etc.}),  (ii) (non supersymmetric) string-inspired standard model extension with Lorentz Violation.  (iii)  Loop Quantum Gravity.  (iv) However, {\\em CPT violation} may also occur at a {\\em global scale}, {\\em cosmologically}, as a result of a cosmological constant in the Universe, whose presence may jeopardize the definition of a standard scattering matrix.   {\\bf (c)} \\underline{HOW CAN WE DETECT CPT VIOLATION?} : Here is a current list of most sensitive particle physics probes for CPT tests: (i) {\\em Neutral Mesons}: KAONS, B-MESONS, entangled states in $\\phi$ and $B$ factories.  (ii) {\\em anti-matter factories}: antihydrogen (precision spectroscopic tests on free and trapped molecules ),  (iii) Low energy atomic physics experiments, including ultra cold neutron experiments in the gravitational field of the Earth.  (iv) Astrophysical Tests (especially  Lorentz-Invariance violation  tests, via  modified  dispersion relations of matter probes {\\it etc.})  (iv) Neutrino Physics, on which we shall mainly concentrate in this talk.   I shall be brief in my description due to space restrictions. For more details I refer the interested reader to the relevant literature. I have tried to be as complete as possible in reviewing the phenomenology of CPT violation for neutrinos, but I realize that I might not have done a complete job; I should therefore apologize for possible omissions in references, but this is not intentional. I do hope, however, that I give a satisfactory representation of the current situation. ",
        "conclusions": "From this brief exposition it becomes clear, I hope, that CPT Violation may not be an academic issue, and indeed it may characterize a theory of quantum gravity. As I discussed, neutrino physics provide stringent constraints on CPT Violation, which in some cases are much stronger  than constraints from neutral meson experiments and factories. In this sense neutrinos may provide a very useful guide in our quest for a theory of Quantum Gravity. For instance, neutrino oscillation experiments provide stringent bounds on many quantum gravity models entailing Lorentz Invariance Violation. There are also plenty of low energy nuclear and atomic physics experiments which yield stringent bounds in models with Lorentz (LV) and CPT violation (notice that the frame dependence of LV effects is crucial for such high sensitivities). It is my firm opinion that neutrino factories, when built, will undoubtedly shed light on such important and fundamental issues and provide definitive answers to many questions related with LV models of quantum space time.   But, as I repeatedly stressed, Quantum Gravity may exhibit Lorentz Invariant (and hence frame independent) CPTV Decoherence. Theoretically the presence of an environment may be consistent with Lorentz Invariance~\\cite{mill}. This scenario is still compatible with all the existing $\\nu$ data, given that the parameters of such models are highly model dependent, and thus subject at present only to constraints by experiment. It is interesting to remark, though, that, in cases where quantum gravity induces neutrino oscillations between flavours or violates lepton number, the sensitivity of experiments looking for astrophysical neutrinos from extragalactic sources may exceed the order of $1/M_P^2$ in the respective figures of merit, and thus is far more superior than the sensitivities of meson factories and nuclear and atomic physics experiments, viewed as probes of quantum mechanics~\\footnote{However, as I remarked previously, the reader should be alert to the fact that there may be novel CPTV effects unrelated, in principle, to LV and locality violations, which are associated with modifications of EPR correlations. Such effects may be inapplicable to neutrinos, and thus testable only in meson factories~\\cite{bernabeu}.}.  Clearly much more work, both theoretical and experimental, is needed before definite conclusions are reached. Nevertheless, I personally believe that research on neutrinos could soon make important contributions to our fundamental quest for understanding the quantum structure of space time. Neutrino research certainly constitutes a very interesting area of fundamental physics, which will provide fruitful collaboration between astrophysics and particle physics, and which, apart from the exciting possibility of non-zero neutrino masses, may still hide even further surprises waiting to be discovered in the near future."
    },
    "0402/astro-ph0402238_arXiv.txt": {
        "abstract": "We measure the temperature of warm gas at planet-forming radii in the disk around the classical T Tauri star (CTTS) TW Hya by modelling the H$_2$ fluorescence observed in \\textit{HST}/STIS and \\textit{FUSE} spectra.  Strong \\Lya\\ emission irradiates a warm disk surface within 2 AU of the central star and pumps certain excited levels of H$_2$. We simulate a 1D plane-parallel atmosphere to estimate fluxes for the 140 observed H$_2$ emission lines and to reconstruct the \\Lya\\ emission profile incident upon the warm H$_2$.  The excitation of H$_2$ can be determined from relative line strengths by measuring self-absorption in lines with low-energy lower levels, or by reconstructing the \\Lya\\ profile incident upon the warm H$_2$ using the total flux from a single upper level and the opacity in the pumping transition. Based on those diagnostics, we estimate that the warm disk surface has a column density of $\\log N($H$_2)=18.5^{+1.2}_{-0.8}$, a temperature $T=2500^{+700}_{-500}$ K, and a filling factor of H$_2$, as seen by the source of \\Lya\\ emission, of $0.25\\pm0.08$ (all 2$\\sigma$ error bars). TW Hya produces approximately $10^{-3}$ $L_\\odot$ in the FUV, about 85\\% of which is in the \\Lya\\ emission line.  From the \\ion{H}{1} absorption observed in the \\Lya\\ emission, we infer that dust extinction in our line of sight to TW Hya is negligible. ",
        "introduction": "Classical T Tauri Stars (CTTSs) are roughly solar-mass pre-main sequence (PMS) stars that are accreting gas from their circumstellar disks. The extent and mass of dust in disks have been determined from IR spectral energy distributions (SEDs) and imaging \\citep[see review by][]{Zuc01}.   Cold gas at large radii is traced by observations of molecules such as CO and HCN \\citep[e.g.,][]{Aik02}.  Warmer neutral and ionized gases are used as diagnostics of accretion \\citep[e.g.,][]{Gom99,Joh02}. However, observing the gas in the disk close to the star, which is important for the formation and evolution of planets, has been difficult.  Gas at these radii can induce planet migration \\citep{Gol79}, dampen the eccentricity of terrestrial planets \\citep{Agn02}, and is neccessary for accretion onto giant planets \\citep{Lis93}.  In search of gas at planet-forming radii, \\citet{Cal91} investigated the formation of CO overtone lines at 2.3 $\\mu$m in disks around CTTSs.  These models predict that CTTSs as a group should show both CO emission and absorption, while Herbig Ae/Be stars should show CO in emission. Using high spectral resolution observations, \\citet{Car93} showed that the CO overtone emission from the probable Herbig Ae/Be star WL 16 arises in the inner disk, confirming the prediction of \\citet{Cal91}.  \\citet{Naj96} find the same result for the Herbig Ae/Be star 1548C27.  However, broader searches at high spectral resolution for CO overtone emission or absorption from CTTSs generally find no evidence for contributions to these lines from the disks around these stars \\citep{Cas96,Joh01b}. On the other hand, \\citet*{Naj03} find that CO fundamental emission near 4.6 - 4.9 $\\mu$m is common among CTTSs.  They suggest that overtone emission has not been detected at the observed CO temperatures (1100 -- 1300 K), mainly because the overtone bands have much lower transition rates.  However, CO is not expected to be the dominant gas constituent in the disks around CTTSs, and it is the major component of the gas, H$_2$,  that we seek to probe here.  Previous observations of H$_2$ have not provided a clear picture of the gas in the disk. \\citet{Thi01} detected cold H$_2$ emission in the pure rotational S(0) and S(1) lines from ISO observations of a large sample of CTTSs, Herbig Ae/Be stars, and debris-disk stars. However, in ground-based observations of the S(1) and S(2) lines, using much smaller apertures than ISO, \\citet{Ric01} and \\citet{She03} did not detect any H$_2$ emission from some of the same sources, indicating that the emission is probably extended. On the other hand, the emission in the H$_2$ 1-0 S(1) transition detected by \\citet*{Bar03} in Phoenix spectra of 4 T Tauri stars is narrow and centered at the radial velocity of the stars, presumably formed in low-velocity circumstellar gas.  While H$_2$ lines in the IR are typically weak and are observed against a strong dust continuum, fluorescent H$_2$ lines dominate the far-ultraviolet (FUV) spectrum of CTTSs at wavelengths longward of \\Lya. \\citet{Ard02} found that narrow H$_2$ lines can be blueshifted by up to 20 \\kms\\ in \\HST/GHRS spectra of CTTSs, observed with a $2\\arcsec\\times 2\\arcsec$ aperture. High-resolution echelle and long-slit spectroscopy with \\HST/STIS is providing the critical data needed to identify the source of the H$_2$ emission. Observations of T Tau reveal two components of H$_2$ emission.  \\citet{Sau03} and \\citet{Wal03} detect off-source UV H$_2$ emission, that is pumped near line-center of \\Lya, and most likely produced where stellar outflows shock the surrounding ambient molecular material.  The same set of fluorescent H$_2$ lines are detected in low-excitation HH objects such as HH43 and HH47 \\citep{Sch83,Cur95}, suggestive of a similar excitation mechanism. \\citet{Wal03} find that the on-source H$_2$ emission is photoexcited by a much broader \\Lya\\ emission line, likely produced by the accretion shock at the surface of T Tau.  In \\citet[][hereafter Paper I]{Her02}, we presented observations of H$_2$ fluorescence in the UV spectrum of TW Hya obtained with \\HST/STIS and \\FUSE.  Several lines of evidence suggested a disk origin for the H$_2$ emission: (i) the H$_2$ emission is not spatially extended beyond a point source in the cross-dispersion direction, (ii) the H$_2$ lines are emitted interior to TW Hya's wind because emission in one H$_2$ line is supressed by \\ion{C}{2} wind absorption, (iii) the H$_2$ line centroids have the same radial velocity as the photospheric lines of TW Hya, (iv) no H$_2$ absorption is detected against the \\Lya\\ emission, and (v) the TW Hya Association is isolated from large reservoirs of interstellar molecular material, making a circumstellar origin for the H$_2$ emission unlikely.  Because the H$_2$ emission is not extended beyond a point source, the spatial resolution of \\HST/STIS restricts the emitting region to be within 2 AU of the star.  Analysis of on-source IR H$_2$ emission lines to date has been limited because typically only one or two lines are observed.  The IR lines could be excited by shocks, UV fluorescence, X-rays, or thermal heating.  In contrast, over 140 H$_2$ lines from 19 different upper levels are detected in the FUV spectrum of TW Hya. The excitation of H$_2$ can be determined from relative line strengths in this rich spectrum either by measuring self-absorption in low excitation members of individual fluorescent progressions or by using an assumed Ly$\\alpha$ profile to reconstruct initial level populations in various pumping transitions.  In this paper, we calculate the temperature of warm gas at planet-forming distances by modelling warm H$_2$ in a disk around TW Hya. We use the H$_2$ emission to reconstruct the Ly$\\alpha$ profile incident on the fluoresced H$_2$. The observed Ly$\\alpha$ profile differs from the incident profile because of interstellar absorption along our line of sight to TW Hya.  We discuss the strength of Ly$\\alpha$ emission and its affect on the circumstellar disk. ",
        "conclusions": "  1.  The FUV continuum rises shortward of 1700 \\AA, which is indicative of an H$_2$ dissociation continuum, although it could also be produced by H$_2$ fluorescence due to FUV continuum pumping or an additional accretion or activity component.  The FUV continuum is not well explained by simple models of a pure hydrogen slab, which are commonly invoked to analyze the excess NUV continuum.  2.  The extinction towards TW Hya is negligible, based on the hydrogen column density in our line of sight and assuming an interstellar gas-to-dust ratio.  3.  Self-absorption of Lyman-band transitions involving low excitation energy levels weakens the flux in these lines. We model this effect by simulating \\Lya\\ emission entering a plane-parallel atmosphere and pumping the H$_2$.  4.  Using the observed H$_2$ fluxes and our fluorescence models, we reconstruct the \\Lya\\ line profile incident upon the warm molecular layer and compare it to the observed \\Lya\\ profile, for a range of assumed temperatures and column densities. Undetected progressions rule out a large region of parameter space in this model.  The reconstructed \\Lya\\ profile is similar to the observed profile in the wings, but shows a much narrower absorption feature than is observed. This narrow absorption component in the reconstructed profile could be a self-reversal or a component of the wind between the source of \\Lya\\ emission and the warm molecular region.   5.  Our models indicate that a molecular layer with a kinetic temperature of $2500^{+700}_{-500}$ K and a column density of $\\log N$(H$_2)=18.5^{+1.2}_{-0.8}$ ($2\\sigma$ error bars) absorbs \\Lya\\ radiation in the surface layer and inner edge of the disk within 2 AU of the central star. The \\Lya\\ pumping leads to a small H$_2$ dissociation rate and  does not cause significant disk dissipation. The filling factor of the warm H$_2$ around TW Hya is $\\eta=0.25\\pm0.08$, although significant uncertainties in the geometry of the fluorescent H$_2$ weaken our confidence in the large filling factor.  6.  The warm H$_2$ most likely resides in a warm surface layer of the disk. This surface layer may be analgous to a classic PDR, although in this case the FUV radiation field is dominated by emission lines rather than a continuum.  In particular, \\Lya\\ comprises about 85\\% of the FUV radiation field below 2000 \\AA, and it controls the excitation and ionization of the disk surface.  7.  Some of the observed H$_2$ lines are pumped from high rotational levels that cannot be excited thermally or by fluorescence. Formation of rotationally excited H$_2$ by reactions with H$_3^+$ may explain these emission lines."
    },
    "0402/astro-ph0402524_arXiv.txt": {
        "abstract": "\\par We investigate the origin of the small, chemically rich molecular clumps observed along the main axis of chemically rich outflows such as CB3 and L1157.  We develop a chemical model where we explore the chemical evolution of these clumps, assuming they are partially pre-existing to the outflow, or alternatively newly formed by the impact of the outflow on the surrounding medium. The effects of the impact of the outflow are reproduced by density and temperature changes in the clump. We find that the observed abundances of CH$_3$OH, SO and SO$_2$ are best reproduced by assuming a scenario where the dense molecular gas observed is probably pre-existing in the interstellar medium before the formation of their exciting (proto)stars and that the clumpiness and the rich chemistry of the clumps are a consequence of a pre-existing density enhancement and of its interaction with the outflow. ",
        "introduction": "At an early stage in their evolution stars eject material in the form of outflows. In fact, the first signs of an outflow are coupled with infall motion and therefore with the first stages of star formation (e.g. Bachiller 1996; Richer et al. 2000). Once the protostar is formed, outflows are the main means of removing the material left over from the collapse of the cloud. \\par Small molecular clumps ($\\sim$ 0.1 pc) are detected in association with several outflows; it is generally believed that the clumps are generated by episodic mass loss of the forming object.  To date, there is no detailed understanding of the role of the clumpiness in outflows. In fact, two main kinds of clumps have been observed: (i) high-velocity clumps or the so-called molecular bullets which are well defined entities travelling at velocities larger than 100 km s$^{-1}$. The prototype is L1448 (Bachiller, Mart\\'{\\i}n-Pintado \\& Fuente 1991), where bullets appear in pairs with the members of a pair being symmetric in both position and velocity with respect to the star. These bullets are most likely associated with mini-bow shocks formed by the outflow propagation (e.g. Dutrey, Guilloteau \\& Bachiller 1997). The molecular lines emitted by this kind of clump are relatively weak, so their chemical composition remains unknown. (ii) A second kind of clump, observed along a few outflows associated with low- and intermediate-mass stars; these outflows stand out because of their association with chemically rich clumps at definitely lower velocity, such as L1157, BHR71, and CB3 (Bourke et al. 1997; Codella \\& Bachiller 1999; Bachiller et al. 2001).  The origin of these chemically rich clumps is not yet clear.  In this paper, we investigate the origin and nature of the second kind of clumps by the use of a chemical model that simulates the clump formation and its subsequent interaction with the outflow. We consider here two main scenarios: (1) pre-existing clumps, affected by the outflow, and (2) newly-formed clumps, created by the outflow.  {\\it Note that we will use the word pre-existing to indicate a density structure formed before the advent of the outflow.} Our definition does not imply that the observed abundances (e.g Codella \\& Bachiller 1999; Bachiller et al. 2001) are pre-existing to the outflow. \\par In what follows, we introduce our two scenarios for the formation of the clumps.  \\begin{itemize}  \\item[(1)] The clumps are pre-existing to the outflow; for example they may be either remnant material of the collapsing parent cloud or completely independent of the star formation process, but present homogeneously in the dark molecular cloud (e.g. Falle \\& Hartquist 2002; Morata et al. 2003; Garrod et al. 2003). In this scenario, we observe these clumps in association with molecular outflows because, as the outflows travel through the molecular cloud, they interact with the clumps and shock them. If this is the case, then the clumps may not be an indication of the episodic nature of outflows.   \\item[(2)] The cloud material is homogeneous and low in density but as an episodic outflow forms, its interaction with the surrounding homogeneous material will lead to a compression (increase in density) and increase in temperature (with subsequent evaporation of the grains mantles) in localized regions {\\it only}, hence the observed clumpiness (e.g Arce \\& Goodman 2001, 2002, and reference therein). In this scenario, the clumps are a direct manifestation of the episodic nature of outflows.  \\end{itemize}   Note that these scenarios do not necessarily exclude each other: in fact, the morphology of the regions where outflows and jets propagate have a very complicated geometry and structure by their very nature (e.g Hester et al. 1998). \\par The main aim of this study is to investigate the nature of the chemically rich clumps by modelling their chemical evolution and by comparing the models with observations. In particular, we attempt to identify observable species that can be used as discriminants between the two scenarios. In this paper, we will focus our attention on the clumps located in the intermediate-mass star forming region CB3 (Codella \\& Bachiller 1999). In addition, a low mass case, represented by the L1157 outflow, will be briefly discussed. Our model is described in Section 2, and our results are presented, discussed and compared with observations in Section 3. A brief conclusion is given in Section 4. ",
        "conclusions": "We have presented here a detailed time-dependent chemical model of the chemically rich clumps observed along outflows.  This preliminary study was aimed at finding some observable tracers that could help us understanding the origin of the clumps with respect to the outflow. We also attempted a qualitative comparison of our models with some observations, in particular with the clumps observed along the CB3 outflow (Codella \\& Bachiller 1999). Despite large uncertainties in the mode of formation and several chemical assumptions we made, we believe that we are able to constrain some of physical and chemical parameters of the CB3 clumps.  Our conclusions are:  \\begin{enumerate} \\item[1.] The initial sulphur abundance of the gas forming the clump can not be solar. We find a depletion factor of $\\sim$ 100, confirming the findings of other studies (Oppenheimer \\& Dalgarno, 1974; Ruffle et al. 1999). \\item[2.] A substantial freeze out must occur during the formation of the clump, regardless of its mode of formation. \\item[3.] Our models indicate that the most likely explanation for the outflow clumps is that they are pre-existing, meaning $only$ that their density structure is, at least partly, formed prior of the advent of the outflow. This does $not$ exclude the general explanation that the outflow clumps are mainly made of swept-up ambient gas and that therefore the clumps are an indication of the episodic nature of the outflows.  \\item[4.] It is probable that, with the advent of the outflow, not only the temperature of the clumps increases and reaches the one observed, but the clumps also undergo a period of non-dissociative shock (and therefore high temperatures, $\\sim$ 1000 K).  \\item[5.]The rich chemistry of the clumps observed along CB3, and L1157, seems to be a consequence of a pre-existing density enhancement (either uniform or already in clumps) and of its interaction with the outflow. The latter, most likely, shocks and accelerate the gas, and possibly, if episodic, induces its clumpiness. This is indicated by the high abundance of methanol and some of the sulphur-bearing species. In fact these molecules are formed by a combination of freeze out and surface reactions, and shocked chemistry: both most efficient when the outflow compresses already dense material at its passage. \\item[6.] We find that it is not possible for the outflow clumps to have a uniform high density - a density gradient is needed in order to account for the observed emission of most species. \\item[7.] CO, CH$_3$OH, CS and SO$_2$ are most likely emitted from the lower density components, while SO and OCS come from an intermediate density component. This chemical stratification supports the findings of Bachiller et al. (2001). \\item[8.] At late times ($t > 10,000$ yr) H$_2$CO is $always$ over-abundant by at least a couple of orders of magnitudes.  This discrepancy is similar to that found for the clumps ahead of Herbig-Haro objects (Viti et al. 2003): these objects however do not share a common chemistry and we find no chemical reason why the abundance of H$_2$CO predicted by our models should be much larger than apparently observed. A possible explanation may be that the observed clumps are smaller than the size implied by the observations. In fact, when computing the theoretical column densities for smaller sizes (see Section 3.3.4), the match between theory and observations improves.  \\end{enumerate}   In conclusion, we suggest that interferometric observations of outflow clumps closer to us than CB3 are performed. This may reveal the real structure of these clumps and help constrain the models."
    },
    "0402/astro-ph0402654_arXiv.txt": {
        "abstract": "We present the first laboratory experiments using a notch-filter mask, a coronagraphic image mask that can produce infinite dynamic range in an ideal Lyot coronagraph according to scalar diffraction theory. We fabricated the first notch-filter mask prototype with .25 $\\micron$ precision using an e-beam lithography machine.  Our initial optical tests show that the prototype masks generate contrast levels of 10$^{-5}$ at 3$\\lambda/D$ and 10$^{-6}$ at $\\sim 8 \\lambda/D$, with a throughput of 27\\%.  We speculate on the ``as-is'' performance of such a mask in the Hubble Space Telescope. ",
        "introduction": "\\label{s1} Directly imaging extrasolar terrestrial planets in reflected light requires facing the extremely high predicted contrast ratios between planets and their host stars, e.g., $\\sim 10^{-10}$ for an Earth analog orbiting a solar type star at quadrature. A planet-finding coronagraph must realize this contrast within a few diffraction widths ($\\lambda/D$, where $\\lambda$ is the wavelength of light, and $D$ is the long axis of the primary mirror) of the stellar image.  Though several coronagraph designs can achieve this contrast according to scalar diffraction theory \\citep{ks03}, substantial work on mask design and laboratory investigation probably remains before this contrast can be achieved in practice.  Some coronagraph designs use image-plane masks to absorb on-axis light and diffract it away \\citep{malbet96,makidon01,kuchner02}.  Other designs use shaped or apodized pupils which benefit from combining aperture shape and the pupil intensity distribution to remove the wings of a circular aperture's PSF \\citep{kasdin1,kasdin2,debes1,debes2,jian1}.  Image masks offer the advantages that they explicitly remove starlight from the beam, and that they can provide high contrast at small angles from the optical axis, given sufficient control over low-spatial frequency modes.  Scattered light, wavefront errors, and mask construction errors can all degrade the contrast of a coronagraph.  For example, for any coronagraphic image mask, mid-spatial frequency intensity errors near the center of the mask must be $\\lesssim 10^{-9}$ \\citep{kuchner02}.  Some of these errors can be controlled using active optics, but these corrections will necessarily apply only over a limited range of wavelengths.  Notch-filter masks offer a promising choice for planet-finding coronagraphs \\citep{kuchner02b}. These image masks absorb most of the light from an on-axis point source, and diffract all of the remainder onto a matched Lyot stop. While Lyot coronagraphs with Gaussian image masks must have a throughput of $\\lesssim 1/2$ to reach 10$^{-10}$ contrast, linear notch-filter masks have unlimited dynamic range according to scalar diffraction theory, and a throughput of $\\sim(1-\\epsilon)$, where $\\epsilon$ is typically $\\sim 0.3-0.5$.  Other coronagraph designs besides notch filter masks can create perfect subtraction of on-axis light.  However, those designs based on masks with odd symmetry \\citep{rouan2,rouan1} or interferometrically synthesized masks with odd symmetry \\citep{baudoz1, baudoz2} create nulls that degrade as $\\theta^2$, where $\\theta$ is the angle from the optical axis.  This rapid degradation means that the finite size of a real star causes the coronagraph to leak light at levels unsuitable for terrestrial planet detection. Other designs, like the dual phase coronagraphic mask with an apodized entrance pupil \\citep{soummer03a,soummer03b}, produce the needed null depth, but must use masks with special chromatic behavior and require precise, achromatic aperture apodization. Notch-filter masks are intrinsically achromatic and like the dual phase coronagraph, they create nulls of order $\\theta^4$ or slower \\citep{kuchner02}.  Notch-filter masks can be designed such that they are binary--regions of the mask are either opaque or transparent.  This is a great advantage as intensity errors are not an issue so long as the mask is sufficiently opaque, leading to manufacturing constraints that are orders of magnitude smaller. However, the shape of the mask must be precisely reproduced, to the level of $\\lambda f_\\#$/3600 for a contrast of $10^{-10}$ within the search area.  For an instrument with $f/100$ and working $\\lambda\\sim.66\\micron$, this corresponds to a tolerance on the order of 20 nm.  Nanofabrication techniques are required to reach this precision.  As part of a joint university-industry study partly funded by Ball Aeorospace and in collaboration with the Penn State Nanofabrication facility (Nanofab), we have fabricated a prototype notch-filter mask and tested it in an experimental setup.  We discuss briefly the mask fabrication process in Section \\ref{s2}, describe methods for modeling performance in Section \\ref{s2a}, review the experiments and results in Section \\ref{s3}, and discuss ways of improving performance in Section \\ref{s4}. ",
        "conclusions": "\\label{s4} Our experiments did not attain the mask performance levels expected from scalar diffraction theory.  In this section we will quantify the effect of some errors that degraded the contrast, and speculate on the potential uses of this mask for space-based planet searches.  MT, the finite size of the point source, mask alignment errors, and mask fabrication errors all combine to explain the degraded performance of the notch filter mask.  These effects can be estimated and collected into an error budget to guide further testing of the mask and drive improvements in our setup.  The MT for dark parts of the mask should be less than the contrast requirements. Degradation from light transmission can be estimated by assuming a $\\lambda/D$ by $\\lambda/D$ hole in the mask with fractional transmission $f$.  The central intensity of the leakage would be $\\sim .25f$ as found in \\citet{kuchner02}.  The transmission flux measured in Table \\ref{tab1} is $3\\times10^{-3}$ times the unblocked point source, giving a transmission peak intensity of 7.5 $\\times10^{-4}$.  This is larger than what is observed at the center, but one can estimate what would be expected further away--the PSF is $\\sim$10$^{-2}$ the peak at the first Airy ring, which for the transmission gives an intensity of $\\sim10^{-5}$, which is more consistent with what is seen further from the center.  The mask may not be uniformly transmissive and slightly thicker toward the center, which could account for the suppression of the peak core. The opaque parts of the mask are covered by a 105 nm thick layer of chromium; if this layer is doubled or tripled it will push the MT to $\\sim10^{-8}$  If the error in alignment with respect to the mask is larger than the physical size of the point source, then the leakage is $\\sim (\\Delta\\theta/\\theta_{1/2})^{4}$ where $\\Delta\\theta$ is the error in alignment and $\\theta_{1/2}$ is the half power position of the mask \\citep{kuchner02}.  The size of our single mode fiber core, 5 $\\micron$, ensures that the leakage due to it is $\\ll$ the leakage due to misalignment of the mask. We have measured the half power of the mask to be $\\sim$8 pixels or 192 $\\micron$ in the focal plane.  Our precision stage had an estimated accuracy of $\\sim$16 $\\micron$, based on half the value of the smallest movement possible in the focal plane.  The leakage would be $\\sim4.8\\times10^{-5}$ that of the unblocked point source.  The surface roughness of our lenses will dictate the levels of scattered light we should observe and allow us to estimate the contribution of scattered light to the degradation in contrast.  We measured the surface roughness of one of our lenses with a profilometer at Nanofab and obtained an estimate of the RMS roughness (See \\citet{elson1}).  Scattered light levels are $\\propto$ $\\delta_{rms}^2$, assuming a collection of plane gratings that diffract (scatter) light into angles of interest. This formalism is for an opaque surface that reflects light.  However, the results for a series of uncorrelated surfaces (i.e. an optical setup of many lenses) give similar results provided that the roughness is separated on scales $>$ $2\\lambda$ \\citep{elson2}. We find that the RMS roughness of the lens is .4 nm, which can be compared to the measured roughness of HST, $\\sim$5.5 nm.  Therefore we estimate that the scattered light levels should be $\\sim(\\delta/\\delta_{HST})^2$ less than that of HST, corresponding to a contrast level of $3\\times10^{-7} (x/14.5)^{-2.19}$, where $x$ is in multiples of $\\lambda/D$ \\citep{bb1,malbet95}.  This corresponds to a scattered light level of $\\sim$9$\\times 10^{-6}$ at 3 $\\lambda/D$ and $1\\times 10^{-6}$ at $8\\lambda/D$. More accurate measurements and analysis are needed to better quantify the limitations of scattered light in the lab, as the above comparison is not necessarily accurate with such small scattering angles \\citep{bb1}.  MT and scattered light dominate the source of errors at $\\sim$3$\\lambda/D$, which is consistent with what is seen in Figure \\ref{fig:comp}.  The resulting PSF with the notch-filter mask resembles the MT PSF, close to the PSF core, where the residual Airy Pattern of the MT dominates.  Further from the core the Airy pattern is less distinct, most likely due to speckles from light scattered from the microroughness of our lenses.  A Lyot stop of diameter $\\sim$2.4 mm should have sufficed for a contrast of 10$^{-6}$.  Experimentally we found that an undersized Lyot stop with 75\\% the diameter of the theoretical design appeared more efficacious.  This was based on an initial belief that the degradation was caused by excess scattered light or slight misalignments of the Lyot stop and the optical beam.  In those cases, undersizing the Lyot Stop would compensate for low levels of leakage.  However, since it is apparent that the main cause of the degradation in contrast is due to the MT, undersizing the stop simply reduces throughput.  The design ``as-is'' already could have significant science benefits in space.  Observations at the scattered light limit of HST coupled with PSF subtraction (shown to give an improvement of contrast of around a factor of 50-100) could yield contrast levels of $\\sim10^{-7}$ \\citep{schneider,grady}.  For HST, the Lyot stop would need to be designed such that the central obscuration and support pads would be adequately blocked at a cost in throughput.  The Lyot stop would be the overlap of three HST pupils, just as in the ideal case. If we assume that with sufficient integration time we can reliably detect planets at this contrast level we can speculate how useful HST would be for a planet search. An instrument on HST optimized for coronagraphy could become a test bed for future TPF coronagraph technology.    This setup would allow a limited extrasolar planet direct imaging survey around nearby stars and white dwarfs.  As an example we consider our reported contrast at 3$\\lambda/D$ in the J~band on HST with PSF subtraction. Given the best results one could expect $\\Delta$J=17.5 and observe 1~Gyr old 3~M$_{Jup}$ planets 10~AU from their host stars out to 30~pc and a 10-100~Myr old 2 $M_{Jup}$ at 6.3 AU around $\\beta$ Pictoris \\citep{burrows}."
    },
    "0402/astro-ph0402148_arXiv.txt": {
        "abstract": "{We present further observations of molecular gas in head-on collisions of spiral galaxies, this time of the CO($J=1\\rightarrow 0$) and CO($J=2\\rightarrow 1$) lines in the UGC~813 -- UGC~816 system. UGC 813/6 are only the second known example of head-on spiral-spiral collisions, the first example being the UGC 12914/5 pair. Strong CO emission is present in the bridge between UGC 813 and 816, unassociated with stellar emission, just as in UGC 12914/5. The CO emission from the UGC 813/6 bridge, not counting the emission from the galaxies themselves, is at least that of the entire Milky Way. Collisions of gas-rich spirals are really collisions between the interstellar media (ISMs) of the galaxies.  We show that collisions between molecular clouds bring H$_2$ into the bridge region.  Although the dense clouds are ionized by the collisions, they cool and recombine very quickly and become molecular again even before the galactic disks separate. Because the clouds acquire an intermediate velocity post-collision, they are left in the bridge between the separating galaxies.  The star formation efficiency appears low in the molecular clouds in the bridges. We speculate that the pre-stellar cores in the molecular clouds may expand during the cloud collisions, thus retarding future star formation. Because the ISM-ISM collisions discussed here require a very small impact parameter, they are rare among field spirals.  In clusters, however, these collisions should be an important means of ejecting enriched gas from the inner parts of spirals. \\\\ ",
        "introduction": "A new class of colliding galaxies has recently come to be recognized -- galaxy encounters in which the neutral interstellar media (ISMs) of the systems hit each other in collisions with very small impact parameters, irrespective of the relative inclinations of the galactic planes. These systems are characterized by a bridge with strong synchrotron emission and abundant atomic hydrogen \\citep[][ hereafter C93]{Condon93}. Two galaxy pairs have clearly undergone such a collision -- the UGC 12914/5 \\citep[C93,][]{Jarrett99} and UGC 813/6 \\citep[][ also VV 769]{Condon02} pairs.  Contrary to initial expectations (C93), extremely strong CO emission was detected throughout the UGC 12914/5 bridge \\citep{Braine03,Gao03}, showing that a large quantity of molecular gas is present as well.  In fact, the CO emission from the UGC 12914/5 bridge was found to be nearly 5 times that of the entire Milky Way.  Due to the appearance of the synchrotron brightness contours and the magnetic field, the UGC 12914/5 system is also known as the \"Taffy galaxies\".  The UGC 813/6 pair shares these properties and thus is a member of the class of taffy galaxies.  In this work, we show that the UGC 813/6 bridge contains a lot of molecular gas, again like UGC 12914/5.  The centers of UGC 813 and 816 were previously observed in CO(1--0) by \\citet{Zhu99}.  It is commonly believed that molecular clouds are too small and dense to hit each other \\citep[{\\it i.e.} they have a low filling factor, {\\it e.g.} ][]{Jog92} or to be affected by collisions with more diffuse atomic or ionized hydrogen clouds.  In addition to presenting the new results for UGC 813/6, the goal of this work is to show that indeed collisions of molecular clouds are plausible and even inevitable in collisions of gas-rich galaxies. As the surface filling factor of the molecular gas in the spiral disks increases, a head-on collision of the galaxies becomes radically more efficient at drawing dense gas out of the parent disks and into the region between the two galaxies which separate after collision.  The fate of the gas is unknown and probably quite dependent on whether the galaxies merge post-collision or not.  Such collisions are, however, likely the only means of forcing much of the gas out of the inner regions of spirals because tidal forces primarily affect the less bound external regions.  The stars are not affected by collisions of the ISMs of the galaxies and of course stars do not collide with each other.  However, the optical appearance can be perturbed by the tidal forces which affect both stars and gas.  The slower the relative velocities of the galaxies in the collision, the stronger the effect of the tidal forces will be.  The UGC 813/6 bridge system is actually part of a triplet, the third member being CGCG 551-011 about 50 kpc away and at the systemic velocity of the UGC 813/6 pair, about 5200 km s$^{-1}$.  Following \\citet{Condon02} we use the optical convention ($v = c z$) for all velocities.  The atomic hydrogen (H{\\sc i}) column densities are extremely high in the bridge region, reaching about $3 \\times 10^{21}$ cm$^{-2}$.  Two H{\\sc i} peaks, H{\\sc i}$_{\\rm E}$ and H{\\sc i}$_{\\rm W}$, are also present in the outer parts of the common H{\\sc i} envelope and are probably tidal features.  The western H{\\sc i} peak coincides with a blue object.  The eastern H{\\sc i} peak has no obvious optical counterpart but is close to the optical tidal tail.  \\citet{Condon02} suggests that UGC 813 and UGC 816 are undergoing their first collision (some 40 to 50 Myrs ago) and are now separating at about 400 -- 500 km s$^{-1}$. Contrary to the UGC 12914/5 system, the UGC 813/6 disks are rotating in the same direction.  CGCG 551-011 is also a spiral galaxy but has been stripped of most of its H{\\sc i} and has a strong compact radio continuum source.  Without numerical simulations, we cannot determine whether the tidal features could have been generated by a collision between UGC~816 and CGCG~551-011 some 200 Myr ago, stripping CGCG~551-011 of most of its H{\\sc i} and feeding the central source.  After describing the observations of the UGC~813/6 system, we present the CO spectra and compare these with the H{\\sc i} data from \\citet{Condon02}. Using the information available for both the UGC 12914/5 and UGC 813/6 pairs, we then show that molecular clouds can collide and drive large amounts of dense gas out of the spiral disks, despite earlier assumptions to the contrary.  The differences with respect to previous work are explained and a general scenario proposed for how the gas, particularly the dense gas, gets into the bridge region for this type of collision. Taking H$_0 = 70$ km s$^{-1}$ Mpc$^{-1}$, we assume a distance to the UGC 813/6 group of 75 Mpc.  \\section {Observations}  The UGC~813/6 bridge, galaxy centers, and the outlying H{\\sc i} maxima were observed in the CO(1$\\rightarrow$0) and CO(2$\\rightarrow$1) to search for molecular gas. The observations were carried out with the 30 meter millimeter-wave telescope on Pico Veleta (Spain) run by the Institut de RadioAstronomie Millim\\'etrique (IRAM) in August 2003.  The CO(1$\\rightarrow$0) and CO(2$\\rightarrow$1) transitions at 115 and 230 GHz respectively were observed simultaneously and in both polarizations.  A bandwidth of over 1300 km s$^{-1}$ was available in both transitions using the two $512 \\times 1$ MHz filterbanks at 3mm and the two $256 \\times 4$ MHz filterbanks at 1mm.  The band was centered around $cz = 5200$ km s$^{-1}$, where $z$ is the redshift, corresponding to redshifted frequencies of 113.3059 and 226.6074 GHz.  \\begin{figure*}[t] \\begin{center} \\includegraphics[angle=270,width=18cm]{0732fig1.ps} \\caption{CO(1$\\rightarrow$0), CO(2$\\rightarrow$1), and H{\\sc i} spectra at the positions observed in the UGC 813/6 galaxies and bridge, using respectively black solid (white solid in the bridge), red dashed (yellow dashed in the bridge), and green dotted lines.  The circle indicates the size of the CO(2--1) beam. UGC 813 is the edge-on spiral to the lower right and UGC 816 is to the left, less inclined and at a higher velocity.  The asterisk marks a bright foreground star.  For the CO spectra, the intensity scale is main beam antenna temperature in mK (-30 to 100 mK for the galaxy centers and -12 to 40 mK for the bridge positions).  The H{\\sc i} spectra are plotted such that equal CO(1$\\rightarrow$0) and H{\\sc i} intensity correspond to equal H-atom column densities in H$_2$ and in H{\\sc i}.  It is thus apparent that in the bridge, the H{\\sc i} and H$_2$ column densities are similar, assuming the $\\ratio$ factor is correct.  The underlying image is an R band image taken with the Instituto de Astrof\\'{\\i}sica de Andaluc\\'{\\i}a (IAA) 1.5--m telescope on Pico Veleta.  } \\end{center} \\end{figure*}  System temperatures were typically 200 -- 300 K at 3mm and twice as high for the CO(2--1) transition (T$_A^*$ scale).  The forward (main beam) efficiencies at Pico Veleta are currently estimated at 0.95 (0.74) at 115 GHz and 0.91 (0.54) at 230 GHz.  At the redshifted frequencies the half-power beamwidths are 22$''$ and 11$''$. All observations were done in wobbler-switching mode, usually with a throw of 100$''$ but sometimes more or less depending on the position observed, in order to be sure not to have emission in the reference beam. Pointing was checked on the nearby quasar 0133+476 every 60 -- 80 minutes with pointing errors usually less than 3$''$.  Data reduction was straightforward.  After eliminating the few obviously bad spectra or bad channels, the spectra for each position and each transition were summed.  Only zero-order baselines ({\\it i.e.} continuum levels) were subtracted to obtain the final spectra.  Where no CO line was obvious, the line windows were based on the H{\\sc i} spectra.  \\section {Results}  The centers of both galaxies and 4 of the 5 bridge regions were clearly detected in both transitions.  In Fig. 1 we show the CO(1--0) and CO(2--1) spectra for the galaxy centers and bridge along with the H{\\sc i} spectra at those positions at 16$''$ resolution, intermediate between the CO(1--0) and CO(2--1) angular resolutions.  The northernmost bridge position shows positive flux at the appropriate velocities.  No line is obvious in CO(2--1) although it covers an optically brighter region (a tidal arm) than the bridge center (Fig. 2).  All of the pointing centers are shown in Fig. 2.  In Fig. 3 we show the CO(1--0), CO(2--1), and H{\\sc i} spectra of the eastern and western H{\\sc i} peaks.  No CO emission was detected despite the strong H{\\sc i}.  Comparison of the CO(1--0) observations at 22$''$ (this paper) and 55$''$ \\citep{Zhu99} suggests that in UGC 813 the molecular gas is more centrally concentrated than in UGC~816 because in UGC 813 the line is much brighter when observed at high resolution.  \\begin{figure*}[t] \\begin{center} \\includegraphics[angle=270,width=18cm]{0732fig2.ps} \\caption{R band image taken with the IAA 1.5meter telescope of the entire UGC 813/6 system with H{\\sc i} column density contours from \\citet{Condon02}. The centers of UGC 813 and 816 are marked by black triangles. The bridge positions observed are the five white triangles between the galaxies.  These are the positions whose spectra are shown in Fig. 1. The triangles to the extreme West and East, slightly north of UGC 816, mark the positions of the H{\\sc i} peaks whose spectra are shown in Fig. 3. Asterisks mark the positions of some galactic stars which could be confused with other features.  The spiral galaxy seen nearly edge-on in the lower right is CGCG~551-011.} \\end{center} \\end{figure*}  \\begin{table*} \\begin{center} \\begin{tabular}{llllllll} Source & RA & Dec & offset & I$_{\\rm CO(1-0)}$ & I$_{\\rm CO(2-1)}$ & N(H$_2$) & N(H{\\sc i}) \\\\ & (J2000) & (J2000) & arcsec & K km s$^{-1}$ & K km s$^{-1}$ & 10$^{20}$ cm$^{-2}$& 10$^{20}$ cm$^{-2}$\\\\ \\hline UGC 813 & 01:16:16.45 & 46:44:25 & (0,0) & 12.2$\\pm 0.9$ & 10.3$\\pm 1.3$ & 24 & 38 \\\\ UGC 816 & 01:16:20.6 & 46:44:52 & (0,0) & 11.1$\\pm 0.6$ & 13.9 $\\pm 2.1$ & 22 & 37 \\\\ Bridge  & 01:16:18.4 & 46:44:39 & (0,0) & 7.0$\\pm 0.2$ & 3.7$\\pm 0.4$ & 14 & 35 \\\\ Bridge  &&& (7,-10) & 6.6$\\pm 0.4$ & 7.4$\\pm 0.5$ & 13 & 27 \\\\ Bridge  &&& (-7,10) & 2.5$\\pm 0.3$ & 0.7$\\pm 0.4$ & 5 & 32\\\\ Bridge  &&& (10,7) & 6.7$\\pm 0.4$ & 4.5$\\pm 0.7$ & 13 & 30\\\\ Bridge  &&& (-10,-7) & 10.5$\\pm 0.5$ & 8.3$\\pm 0.6$ & 21 & 35 \\\\ H{\\sc i}$_{\\rm E}$ & 01:16:24.31 & 46:45:24 & (0,0) & 0$\\pm 0.14$ & 0$\\pm 0.4$ & $\\la 1$ & 22 \\\\ H{\\sc i}$_{\\rm E}$&&& (8,-8) & 0$\\pm 0.10$ & 0$\\pm 0.2$ & $\\la 0.6$ & 8\\\\ H{\\sc i}$_{\\rm W}$ & 01:16:06.1 & 46:45:31 & (0,0) & 0$\\pm 0.08$ & 0$\\pm 0.2$ & $\\la 0.5$ & 10 \\\\ \\end{tabular} \\caption[]{Complete list of positions observed in CO, corresponding to the triangles in Figure 2.  The offsets are in arseconds with respect to the (0,0) position of the source. Columns 3 and 4 give the CO(1--0) and CO(2--1) integrated intensities and cols. 5 and 6 the H$_2$ and H{\\sc i} column densities. The CO line window for the \"H{\\sc i}$_{\\rm E}$\" positions is 5370 to 5470 km s$^{-1}$ and 5170 to 5270 km s$^{-1}$ for the \"H{\\sc i}$_{\\rm W}$\" position, both based on the H{\\sc i} spectra shown in Fig. 3.  Uncertainties are 1$\\sigma$ based on the rms noise and the line width.  Upper limits to N(H$_2$) are $3\\sigma$.} \\end{center} \\end{table*}  In Table 1 we give the CO intensities for each transition and the corresponding H$_2$ and H{\\sc i} column densities for each of the positions observed.  While the CO emission from the galaxies is stronger than in the bridge, the CO emission from the bridge is at least that of the Milky Way. For a CO -- H$_2$ conversion factor of $\\ratio = 2 \\times 10^{20}$ molecules cm$^{-2}$ (K km s$^{-1}$)$^{-1}$ \\citep[e.g.][]{Dickman86}, the bridge is about half molecular and half atomic gas -- some 2 $\\times 10^9$M$_\\odot$ of each.  This is based on the central bridge position alone.  For the other positions, it is difficult to separate the emission from the disks from that of the bridge, such that the above represents a lower limit to the true bridge CO emission.  \\begin{figure}[t] \\begin{center} \\includegraphics[angle=270,width=8.8cm]{0732fig3.ps} \\caption{CO(1--0), CO(2--1) and H{\\sc i} spectra of the \"TDG candidate\" positions.  The brightness scale is for the CO(1--0) line and should be doubled for the CO(2--1) line ({\\it i.e.} the  CO(2--1) spectra have been divided by two). The H{\\sc i} scale has been divided by 4 such that if the column density of molecular gas were 25\\% of the atomic gas column density, then the spectra would be of the same height.  It is immediately obvious that the CO intensities indicate low H$_2$ column densities compared to H{\\sc i}. The lower left panel shows the sum of the two eastern positions shown in the top two panels.  The positions are shown in Fig. 2 where they are indicated by triangles. } \\end{center} \\end{figure}  As can be readily seen from Fig. 1, the molecular and atomic gas line profiles are very similar.  Furthermore, assuming the $\\ratio$ factor used is appropriate, the hydrogen column densities are also similar in the bridge region (the H{\\sc i} spectra in Fig. 1 are scaled so that the same line area as in CO represents the same H-atom column density). The molecular gas dominates in the galactic centers and in the densest parts of the bridge.  This is obviously not true of tidal tails, which preferentially bring material out of the H{\\sc i} - dominated outer regions of spirals.  The CO line intensities given in Table 1 for the bridge positions are unheard of for tidal material.  The two H{\\sc i} peaks, at about 1-2 arcmin from the bridge and marked as H{\\sc i}$_{\\rm E}$ and H{\\sc i}$_{\\rm W}$, appear to be at the ends of tidal arms or tails and indeed the H{\\sc i} emission is strong but no CO has been detected there so far (Fig. 3). The H$_2$/H{\\sc i} mass ratio in the tidal material is less than 10\\%, compatible with the less evolved Tidal Dwarf Galaxies (TDG) for which CO detections have been obtained \\citep{Braine_tdg2}.  Figure 4 presents an H$\\alpha$ image of the UGC~813/6 system with H{\\sc i} contours superposed. No star formation is observed near the northeastern H{\\sc i} peak (H{\\sc i}$_{\\rm E}$), which is clearly not, or at least not yet, a TDG.  To further distinguish between bridge and tidal features, we plot the synchrotron emission \\citep{Condon02} over the R band image in Fig. 5. Just as in the Taffy galaxies, the CO emission follows the synchrotron more than any tidal features.  Strong extra-disk synchrotron emission is one of the characterizing features of ISM-ISM collisions and not of tidal arms or tails.  In the Taffy galaxies (UGC 12914/5), the CO/H{\\sc i} ratio in the bridge is higher.  The CO and H{\\sc i} line profiles in the UGC 12914/5 bridge are quite different; this is, however, probably a projection effect because we see the H{\\sc i} from UGC 12914 as if it were in the bridge while it is actually far behind \\citep[see ][]{Braine03}. UGC 12914 and 12915 are generally bigger and more gas-rich than the UGC 813/6 system.  As we will see, this is important for the amount of molecular gas that can be stripped through cloud -- cloud collisions.  \\section {Why is H$_2$ found in the bridges?}  In the following we will address the question of the origin of the H$_2$ found in the bridge between UGC~813 and UGC~816.  As this galaxy pair is very similar to the Taffy galaxies UGC~12914/5 and we also found abundant molecular gas there, we will attempt to provide a more general explanation which applies to any Taffy--type interaction.  We will assume that the $\\ratio$ factor adopted here is valid for the bridge molecular gas. \\citet{Gao03} adopt a higher $\\ratio$ factor in their work on UGC 12914/5; \\citet{Braine03} discuss why the factor they adopt is probably too high. The first question is perhaps whether the molecular gas could form from the atomic gas in the bridge.  \\subsection{Could the H$_2$ form out of the H{\\sc i} in the bridge?}  The H{\\sc i} column density is typically $2 - 4 \\times 10^{21}$ cm$^{-2}$ in the bridges.  If the depth of the bridge is similar to its extent perpendicular to the collision direction, then the neutral gas is spread over some 10 -- 20 kpc. The average volume density is then $\\la 0.1$.  \\citet{Braine_tdg2} estimate that $10^7/n$ years ({\\it i.e.} $\\ga 10^8$ yrs in the bridges) is necessary to convert 20\\% of the H{\\sc i} into H$_2$.  Once some of the H{\\sc i} has become H$_2$ the process becomes less efficient.  While local density enhancements could allow some H$_2$ formation from the atomic gas, it is unlikely that the large quantities of H$_2$ observed could have formed from H{\\sc i}.  H$_2$ is in fact the {\\it dominant} phase in the Taffy bridge and is roughly half of the UGC~813/6 bridge ISM. If the grains are warm then the ability of H-atoms to form H$_2$ on the grain surfaces is greatly reduced \\citep{Hollenbach71a}.  Furthermore, the efficiency of H{\\sc i} to H$_2$ conversion is obviously reduced if grains are destroyed in the strong shocks generated in cloud collisions (whether in atomic or molecular clouds).  There is a further reason for why clouds which are not initially dense will not be able to contract.  In both systems, and particularly the Taffy galaxies, the spirals are massive (several $10^{11}$ M$_\\odot$ each from C93 and \\citet{Condon02}).  As a result, the tidal forces they exert on clouds are stronger than in the Milky Way.  Following \\citet{Bania86}, for clouds to resist against tidal shear, their density must be $$ n \\ga 200 \\, ({{3 {\\rm \\, kpc}} \\over{R}})^2 {\\rm nucleons \\, \\, cm}^{-3}$$ where R is the distance from the cloud to the galactic centers. Being able to resist tidal forces is not sufficient for cloud survival but is certainly required.  Thus, clouds which are not sufficiently dense after collision to form H$_2$ rapidly may not be dense enough to survive to form H$_2$ subsequently in the bridge.  \\subsection{GMC collisions as part of galaxy collisions}  The observations show that when spiral galaxies physically collide, with the inner parts of their disks passing through each other, more than $10^9$ M$_\\odot$ of molecular gas is found in the resulting bridge region between the two galaxies.  While tidal forces can drive gas out of spiral galaxies, they act on stars in the same way, such that it is difficult to imagine how a bridge with so much gas but no stars could have been formed by tidal forces alone. It has commonly been accepted \\citep[e.g.][]{Jog92} that Giant Molecular Clouds (GMCs) do not collide whereas the more diffuse clouds of atomic gas do. The goal of this section is to demonstrate that in gas-rich but not exceptional spiral galaxies GMC-GMC collisions occur and are likely the source of the large quantities (1 -- 10 $\\times 10^9$M$_\\odot$) of molecular gas present in the bridges.  \\begin{figure*}[t] \\begin{center} \\includegraphics[angle=270,width=18cm]{0732fig4.ps} \\caption{H$\\alpha$ image after broadband continuum subtraction taken with the San Pedro M\\'artir, Baja California (Mexico), 2.12--m telescope with the same H{\\sc i} column density contours as in Fig. 2. The dark vertical stripe through the bridge is due to the bright star above UGC~816. A fairly bright H{\\sc ii} region is seen in the southern part of the bridge.  No H$\\alpha$ emission is observed near the NE H{\\sc i} peak.} \\end{center} \\end{figure*}  \\begin{figure*}[t] \\begin{center} \\includegraphics[angle=270,width=18cm]{0732fig5.ps} \\caption{R band image with 1.4 GHz synchotron emission contours superposed.  Triangles indicate the positions observed in CO and the asterisk a star.  While the optical emission in the NW bridge position (near the northern tidal arm) is stronger than in the others, the CO emission is much weaker there. } \\end{center} \\end{figure*}  In the calculations we will suppose that molecular clouds are spherical and obey the so-called \\citet{Larson81} relation with an average column density of N(H$_2$) $= 10^{22}$ cm$^{-2}$.  This is typical and corresponds exactly to what \\citet{Jog92} assume. We define the following quantities: \\\\ $A_{cl}$, the projected cloud area, equal to $\\pi R_{cl}^2$ for identical spherical clouds where $R_{cl}$ is the cloud radius; \\\\ $n_{cl}$, the number density of such clouds; \\\\ $M_{cl}$, the mass of a cloud of radius $R_{cl}$; \\\\ $z_{\\rm H_2}$, the half thickness of the layer of molecular clouds;\\\\ $\\Sigma_{\\rm H_2}$, the face-on H$_2$ surface density of the spiral galaxies, expressed in M$_\\odot$ pc$^{-2}$, typically averaged over a telescope beam;\\\\ $\\Sigma_{cl}$, the H$_2$ surface density of a cloud which is simply 160 M$_\\odot$ pc$^{-2}$ for N(H$_2$)  $= 10^{22}$ cm$^{-2}$; \\\\ and the velocities $V_{\\perp}$, taken to be the encounter speed perpendicular to the galactic planes and $V_{\\parallel}$, the relative cloud velocities in the planes (which are assumed parallel as in the Taffy galaxies). The link to extragalactic observables is provided by the integrated CO(1$\\rightarrow$0) line intensity I$_{\\rm CO}$, expressed in K km s$^{-1}$, via the conversion factor taken to be $\\ratio = 2 \\times 10^{20}$ H$_2$ cm$^{-2}$ per K km s$^{-1}$.  The mean free path $\\lambda_{\\rm mfp}$ of a spherical particle is $\\lambda_{\\rm mfp} = 1/(4 A_{cl} n_{cl})$. For a gas of particles in random motion, the mean free path is actually $\\sqrt{2}$ shorter.  The factor 4 comes from the fact that clouds collide if they pass within 2$R_{cl}$ of each other.  The number density of clouds is $$n_{cl} = {{\\Sigma_{\\rm H_2}} \\over {2 z_{\\rm H_2}}} \\times {{1} \\over {M_{cl}}}, \\, \\, \\, {\\rm where} \\, \\, \\, M_{cl} = \\Sigma_{cl} \\times A_{cl}$$ such that the mean free path is simply $$\\lambda_{\\rm mfp} = {{2 z_{\\rm H_2} \\Sigma_{cl}} \\over {4\\Sigma_{\\rm H_2}}} $$ such that when $\\Sigma_{\\rm H_2} = \\Sigma_{cl}$, the mean free path of a cloud is 1/4 of the thickness of the molecular cloud layer $2 z_{\\rm H_2}$.  This is in some ways a trivial result because $\\Sigma_{\\rm H_2} = \\Sigma_{cl}$ means that the galactic planes are completely covered by a layer of molecular clouds so it seems obvious that they should collide when the disks pass through each other. CO observations show that many spirals have disk CO intensities of order I$_{\\rm CO} \\ga 20$ K km s$^{-1}$, corresponding to $\\Sigma_{\\rm H_2} \\ga 65$ M$_\\odot$ pc$^{-2}$, such that the mean free path is less than the thickness of the molecular cloud layer.  The mean free path formula above gives collisions even when only parts of the spherical clouds overlap. However, when there is a transverse velocity, such as in the Taffy galaxies where the counter-rotation yields high encounter speeds parallel to the galactic planes, then the path of a cloud in the molecular layer of the other system is increased from 2 $z_{\\rm H_2}$ to 2 $z_{\\rm H_2} \\sqrt{V_{\\parallel}^2 + V_{\\perp}^2} / V_{\\perp}$. When the parallel and perpendicular velocities are roughly equal, this lengthens the path and increases the probability of cloud collisions by $\\sqrt{2}$.  While the magnetic field energy in a cloud is greatly less than the kinetic energy of the collision, magnetic fields are nonetheless quite strong in dense gas, with the field strength increasing as $\\sqrt{\\rm density}$, reaching about 100 $\\mu$ G \\citep{Crutcher99}. As such, they could be expected to increase the effective cloud cross-section somewhat.  If the galactic planes are not parallel and/or if the collision is not face-on then the cloud paths within the galactic planes are increased.  The above discussion only considers collisions between molecular clouds, showing that rather than being rare, GMC - GMC collisions on large scales are inevitable in collisions of gas-rich spirals. After (inelastic) collision, the resulting cloud and possibly cloud fragments will have a velocity intermediate between the velocities of the galactic disks.  They are thus left behind in the bridge region between the separating galaxies.  Whether or not the clouds fall back onto one of the galaxies depends on their velocity with respect to the galaxies and thus the collision velocity.  \\subsection{Why previous results were different}  We choose to compare here with the work of \\citet{Jog92} who explored this question and have been cited for why GMCs should {\\it not} collide. Furthermore, the cloud properties they assume are in accord with our hypotheses -- the only significant number, the average H$_2$ column density (or projected H$_2$ surface mass density) is the same as our adopted value. \\citet{Jog92} define the mean free path by introducing a volume filling factor for molecular clouds, which is the cloud volume multiplied by the number of clouds per unit volume in the molecular layer of a spiral galaxy.  Hence \\begin{eqnarray} \\nonumber f & = & 4 \\pi \\, R_{cl}^3 \\, n_{cl} /3 \\\\ \\nonumber & = & (4 \\pi R_{cl}^3 /3) \\times \\Sigma_{\\rm H_2}/(2 z_{\\rm H_2} \\Sigma_{cl} A_{cl}) \\\\ \\nonumber & = & (4 R_{cl} / 3) \\times \\Sigma_{\\rm H_2} / (2 z_{\\rm H_2} \\Sigma_{cl}) \\end{eqnarray} after making the substitutions using the formulae above for identical spherical clouds (the hypothesis adopted by \\citet{Jog92}).  Their mean free path is then $$ \\lambda_{\\rm mfp} = 2 R_{cl} / 3 f = 2 R_{cl} / 3 \\times 1 / (4 \\pi / 3) R_{cl}^3 \\, n_{cl} = 1/2 A_{cl} n_{cl} $$ which is twice the value we use.  This may be to ensure that only collisions with comparable molecular column densities are included, such that grazing encounters are entirely excluded.  This is not the origin of the discrepancy.  \\citet{Jog92} specify the volume filling factor $f$ to be $f_{\\rm GMC} = 0.01$ which we now express in terms closer to observables. Using the above definitions, \\begin{eqnarray} \\nonumber \\Sigma_{\\rm H_2} & = & f \\times 2 z_{\\rm H_2} \\Sigma_{cl} A_{cl} \\times 3 / (4 \\pi R_{cl}^3) \\\\ \\nonumber & \\approx & {{f} \\over {0.01}} \\times {{2 z_{\\rm H_2}} \\over {130{\\rm pc}}} \\times {{\\Sigma_{cl}} \\over {R_{cl}}} {\\rm M}_\\odot {\\rm pc}^{-2} \\\\ \\nonumber & \\approx & 6 {\\rm M}_\\odot {\\rm pc}^{-2} \\end{eqnarray} for  $2 z_{\\rm H_2} = 130$pc, $\\Sigma_{cl} = 160 {\\rm M}_\\odot {\\rm pc}^{-2}$, and $R_{cl} = 25$pc, as in \\citet{Jog92}.  $6 {\\rm M}_\\odot {\\rm pc}^{-2}$ corresponds to I$_{\\rm CO} = 2$ K km s$^{-1}$, which is weak for the inner parts of a spiral galaxy.  {\\it This is the origin of the discrepancy: the \\citet{Jog92} calculations assume a very low molecular gas surface density.} The value they assume is roughly the molecular gas surface density of the Milky Way averaged over the entire optical disk.  In the molecular ring of our Galaxy the H$_2$ surface density is much higher.  This is why grazing encounters should not be efficient at generating GMC-GMC collisions but collisions of the inner parts of spirals will be. ",
        "conclusions": "Consistent with the Taffy galaxies, UGC 12914/5, the second such system, UGC 813/6, also contains large amounts of molecular gas in the bridge region between the galaxies.  UGC 813 and 816 are considerably smaller and less gas-rich than UGC 12914 and 12915, and the CO emission from the UGC~813/6 bridge is about 4 times less than from the UGC 12914/5 bridge.  Nonetheless, the CO emission from the UGC~813/6 bridge is similar to that from the entire Milky Way and the H$_2$/H{\\sc i} mass ratio is roughly unity.  It is shown that for reasonably gas-rich, but not extraordinary, spiral galaxies a head-on collision will result in many collisions of molecular clouds. Some earlier calculations had reached the opposite conclusion but they had supposed very gas-poor disks.  Although the neutral (both atomic and molecular) clouds are ionized in collisions, the dense gas cools very rapidly allowing it to recombine and become molecular while still in the colliding disks.  The new cloud is then left in the bridge because the collision has left it with an intermediate velocity.  Some H$_2$ may form through collisions of atomic gas clouds just as some of the originally molecular (or atomic) gas will remain ionized after collision.  In addition to showing that GMC -- GMC collisions occur in cases where the nuclei of the galaxies have passed within a few kpc of each other, several factors lead us to conclude that GMC -- GMC collisions play a major role in bringing molecular gas into the bridge regions.  From a theoretical point of view, the study by \\citet{Hollenbach89} indicates that a column density of about $4 \\times 10^{21}$ cm$^{-2}$ is necessary to form H$_2$.  This is a large column density for purely atomic clouds.  Furthermore, if H{\\sc i} - H{\\sc i} cloud collisions produced the H$_2$, the molecular gas would be as extended as the H{\\sc i} in the bridge -- this is not the case.  Being much denser than atomic gas clouds, GMCs will be able to cool much more quickly after collision and ionization, forming the molecular gas we now observe in the bridge.  The star formation in the bridge material is below that expected based on the molecular gas mass.  We speculate that the cloud collisions may have eliminated many of the pre-stellar cores, such that in the few $10^7$ years since the collision the cores have not had time to reform and generate the strong stellar luminosity expected.  ISM-ISM collisions may provide an efficient mechanism for driving enriched gas out of the inner parts of spirals in galaxy clusters.  Given that clusters were even denser in the past, most cluster galaxies have probably undergone such a collision at least once.  Only a very small fraction of field spirals, even at high redshift, would have suffered a head-on collision.  The future of the bridge material is as yet unclear.  With a velocity of $\\la 300 \\kms$ with respect to one or another of the spirals, some of the gas is probably destined to fall back onto the parent spirals. Some however may remain in between the receding galaxies."
    },
    "0402/astro-ph0402181_arXiv.txt": {
        "abstract": "We present an updated phase-coherent timing solution for the young, energetic pulsar \\psr\\ for twenty years of data. Using a partially phase-coherent timing analysis, we show that the second frequency derivative is changing in time implying a third frequency derivative of $\\nudotdotdot= (-9 \\pm 1) \\times 10^{-32} $s$^{-4}$. This value is consistent with the simple power law model of pulsar rotation. ",
        "introduction": "Radio pulsars are powered by rotational kinetic energy and experience losses due to the emission of electromagnetic energy while spinning down. The spin-down of radio pulsars is described by (Manchester \\& Taylor, 1977) \\begin{equation} \\dot{\\nu} \\propto -{\\nu}^n , \\end{equation} where $\\nu$ is the spin frequency, $\\dot{\\nu}$ is the frequency derivative and $n$, the braking index, equals 3 for a spinning magnetic dipole in a vacuum. Taking a derivative of (1), we find that $n = \\nu \\ddot{\\nu}/{\\dot{\\nu}^2}$. Taking further derivatives of (1), and defining the second braking index as $m_0=n(2n-1)$\\cite{Kaspi}, we find an expression $m = {{\\nu^2} \\nudotdotdot}/{\\dot{\\nu}^3}$. Measuring a value of $m$ that agrees with the predicted value $m_0$ is an excellent check on the validity of the power law spin-down model of radio pulsars given in equation (1). Kaspi et al. (1994) reported a value of $n=2.837 \\pm 0.001$ and $m=14.5 \\pm 3.6$ for \\psr. ",
        "conclusions": "Performing a phase-coherent timing analysis on twenty years of radio timing data for \\psr\\ obtained from the Molongolo Observatory Synthesis Telescope and the Parkes 64m telescope, we found the second frequency derivative to be $\\ddot{\\nu} = 1.962(1) \\times 10^{-21}$~s$^{-3}$, implying a braking index of $n = 2.814(1)$. Significant low-frequency timing noise in the data prohibited a phase-coherent analysis to determine an accurate value of the third frequency derivative. This prompted a partially phase-coherent timing analysis by breaking up the data into subsets of $\\sim 2$ yr (such that phase residuals were `white' in each interval following a fit of $\\nu$, $\\dot\\nu$ and $\\ddot\\nu$) and plotting the values of $\\ddot{\\nu}$ against time, the results of which are shown in Figure 1. This method is less sensitive to timing noise which dominates on long time scales. We found a value $\\nudotdotdot = (-9 \\pm 1) \\times 10^{-32} $s$^{-4}$ by a weighted least-squares fit to the data. This value for $\\nudotdotdot$ is in agreement with that predicted from the spin-down law of $\\nudotdotdot = -9.37 \\times 10^{-32} $s$^{-4}$. This measured value for $\\nudotdotdot$ implies a second braking index of $m= 12.7 \\pm 1.4$, also in agreement with the predicted value of $m=13.3$, though less than the value $m=15$ expected for an ideal spinning dipole. Our measurement of the $\\nudotdotdot$ implies that the spin-down law correctly predicts the behaviour of \\psr, reinforcing the use of the power law given in equation (1) to describe the rotation of pulsars.  \\begin{figure} \\begin{center} \\centerline{\\psfig{file=livingstonem_1.ps,width=8cm}} \\end{center} \\caption{Phase-coherent subsets plotted versus epoch. The slope of the line is the third derivative and has a value of $(-9 \\pm 1) \\times 10 ^{-32} $s$^{-4}$. } \\end{figure}"
    },
    "0402/astro-ph0402462_arXiv.txt": {
        "abstract": "We explore the pulsar slot gap (SG) electrodynamics up to very high altitudes, where for most relatively rapidly rotating pulsars both the standard small-angle approximation and the assumption that the magnetic field lines are ideal stream lines break down. We address the importance of the electrodynamic conditions at the SG boundaries and the occurrence of a steady-state drift of charged particles across the SG field lines at very high altitudes. These boundary conditions and the deviation of particle trajectories from stream lines determine the asymptotic behavior of the scalar potential at all radii from the polar cap (PC) to near the light cylinder. As a result, we demonstrate that the steady-state accelerating electric field, $E_{\\parallel}$, must approach a small and constant value at high altitude above the PC.  This $E_{\\parallel }$ is capable of maintaining electrons moving with high Lorentz factors ($\\sim {\\rm a~few}~ \\times 10^7$) and emitting curvature $\\gamma $-ray photons up to nearly the light cylinder.  By numerical simulations, we show that primary electrons accelerating from the PC surface to high altitude in the SG along the outer edge of the open field region will form caustic emission patterns on the trailing dipole field lines.  Acceleration and emission in such an extended SG may form the physical basis of a model that can successfully reproduce some pulsar high-energy light curves. ",
        "introduction": "There is no doubt that pulsars are accelerating particles up to relativistic energies in their magnetospheres, and that these particles are primarily responsible for the pulsar radio- to high-energy non-thermal emission. It is also believed that the energetics of this acceleration, as well as the main physical processes involved in production of high-energy photons, are more or less understood. However, the ambiguity in interpretation of pulsar timing observations in terms of emission site mapping in a pulsar magnetosphere makes it difficult to answer the basic question of where the pulsar high-energy emission originates.  In their recent attempt to explain the observed high-energy light curves of pulsars, Dyks \\& Rudak (2003) concentrated on a purely geometrical model by postulating that the emission is produced in a relatively narrow region along the last open magnetic field lines of a pulsar magnetosphere. The interesting result of their study is the occurrence of caustic emission zones (Morini 1983), i.e. the phase shifts of radiation emitted at radii between $\\sim 0.1 - 0.7$ times the light cylinder radius, parallel to field lines on the trailing edge of the polar cap (PC), are cancelled by phase shifts due to relativistic effects of aberration and time-of-flight.  Radiation emitted over a large range of altitudes thus arrives in phase, forming two narrow peaks in the light curves, very similar to those of known $\\gamma$-ray pulsars (e.g. Thompson 2001).  In our previous paper (Muslimov \\& Harding 2003 [MH03]) we began discussing the regime of acceleration of particles and production of high-energy emission within the pulsar slot gap (SG), a narrow region on the boundary of the open field lines, where the electric field drops to zero.  The SG is a pair-free region of slower acceleration, in which the parallel electric field is unscreened. Pair cascades develop along the inner edge of the SG at several stellar radii above the NS surface. Even though the SG regime in pulsars was originally introduced in the electrodynamic model of Arons \\& Scharlemann (1979), it was not considered a viable high-energy emission region (see e.g. Arons 1996). The revised version of the SG regime proposed by MH03 incorporates the effect of relativistic frame dragging (Muslimov \\& Tsygan 1992 [MT92]) and, more importantly, the effect of SG boundaries on the strength of the accelerating electric field within the SG.  MH03 demonstrated that the primary electrons tend to accelerate up to higher altitudes before pair production begins, and pair cascades continue along the inner boundary of the SG until the magnetic field becomes too low.  The resulting radiation from the pair cascades forms a wide, hollow cone of high-energy radiation due to the flaring of field lines.  Adhering to the small-angle approximation, MH03 restricted their study of the SG regime to altitudes less than four-five stellar radii.  However, since the parallel electric field in the SG is not screened on field lines close to the open-field boundary, acceleration may continue to much higher altitudes.  Particle acceleration and radiation in such an extended SG may therefore provide a physical basis for the two-pole caustic model of Dyks \\& Rudak (2003).  Formation of a SG requires the production of enough pair multiplicity to screen the parallel electric field above the pair formation front.  We have found from our previous studies (Harding \\& Muslimov 2001 [HM01], 2002 [HM02]) that the youngest and most energetic pulsars can produce pairs from curvature radiation (CR) of primary electrons, which are numerous enough to screen the electric field.  Older, less energetic pulsars, those below the CR pair death line, can produce only pairs from inverse Compton radiation of primary electrons scattering thermal X-rays from the NS surface.  The inverse Compton pairs are not numerous enough to completely screen the parallel electric field.  A necessary condition for formation of a SG is thus the ability to produce pairs from CR, and the expression for the CR death line (given by Eqn [52] of HM02) defines the boundary in the $P$-$\\dot P$ diagram of pulsars capable of having SGs. Such pulsars include the Crab, Vela, Geminga and most of the $\\gamma$-ray pulsars detected by EGRET, but not the majority of millisecond pulsars.  The extension of the regime of SG acceleration to much higher altitudes is the main subject of the present paper. In the Sections below we outline our approach to constructing an appropriate steady-state physical solution that can be used up to very high altitudes in the SG. We also discuss the immediate consequences of our proposed extended SG solution: acceleration of particles (electrons, positrons) and high-energy emission up to nearly light-cylinder radius, and the possibility of occurrence of high-altitude caustic emission on trailing field lines.  The paper is organized as follows. In \\S 2 we discuss the electrodynamics within the SG regions of pulsars. We address the physical constraints on the scalar potential (\\S 2.1) and equipotentiality of SG boundaries and derivation of effective Poisson's equation (\\S 2.2) in the outermost section of SG. In \\S 3 we present the electrodynamic solution within the SG at very high altitudes. In \\S 3.1 we illustrate the possibility of extended acceleration within SG in the regime where the acceleration is balanced by the curvature-radiation reaction. Our numerical calculations are discussed in \\S 3.2. Finally, in \\S 4 we discuss our main results and draw our principal conclusions. ",
        "conclusions": "In this paper we extend our previous study (MH03) of pulsar SGs to investigate the regime of acceleration of primary electrons and high-energy emission to very high altitudes.  Incorporating the effect of cross-field motion of charges near the light cylinder on the distribution of accelerating potential within the entire SG, we derived the explicit expressions for the accelerating electric field in the space-charge-limited flow approximation.  We have modeled the particle acceleration up to high altitudes and generation of high-energy emission within the SG, illustrating the resulting high-energy radiation pattern and light curves for the case of the Crab pulsar.  The primary (space-charge-limited) flow within the SG becomes radiation-reaction limited at high altitudes, such that the energy gain from acceleration is nearly compensated by the CR losses.  The resulting emission pattern exhibits both the hollow cones centered on each magnetic pole (the corresponding phenomenological model was first discussed in connection with radio emission by Komesaroff 1970, Radhakrishnan 1969, and Radhakrishnan \\& Cooke 1969), due to pair cascades on the inner edge of the SG, as well as the caustics along the trailing field lines at high altitude noted in previous studies (Morini 1983, Romani \\& Yadigaroglu 1995, Cheng et al. 2000, Dyks \\& Rudak 2003).  The extended SG acceleration allows the caustic emission to be viewed from both magnetic poles simultaneously, naturally producing the widely-spaced double-peaked profiles seen in many high-energy pulsars.  Note that in a vacuum NS magnetosphere with the strong magnetic field prohibiting transfield motion, the GJ space charge may serve as a source of electric field accelerating test particles (e.g. electrons/positrons) along the magnetic field.  In the SG geometry with space-charge-limited flow (nearly GJ near the surface), the occurrence of a steady-state cross-field drifting at very high altitudes tends to reduce the GJ space charge (see condition [\\ref{Eperp}] and related discussion) within the SG, maintaining the SG boundaries as equipotential surfaces. The reduction of the GJ space charge is accompanied by a steady-state distribution of the effective flux, $\\Delta F$, within the SG. Our model requires that the source of a steady-state accelerating electric field within the SG establishes at relatively low altitudes (e.g. over a few stellar radii above the PC) and remains constant along the magnetic stream lines up to very high altitudes.  Whether or not a steady-state particle flow can be achieved on a given field line will depend on the relative importance of the terms $\\propto \\cos \\chi$ and $\\propto \\theta_0 \\sin \\chi \\cos \\phi _{pc}$ near the NS surface, which depends on the pulsar inclination angle $\\chi$, the magnetic azimuth $\\cos \\phi _{pc}$ and the PC opening angle $\\theta_0$. If the term $\\propto \\cos \\chi$ dominates, then $E_{\\parallel}$ accelerates the same sign of charge at all altitudes and steady-state flow can be maintained.  However, if the term  $\\propto \\theta_0 \\sin \\chi \\cos \\phi _{pc}$ begins to dominate near the surface for large $\\chi$ on unfavorably curved field lines ($\\cos \\phi _{pc} < 0$), then $E_{\\parallel}$ changes sign and may prevent a steady-state flow of charge (or any charge flow) along that field line (see Section 3.2.1).   Even if $E_{\\parallel}$ does not change sign above the PC, the possibility we have explored in this paper of a steady-state regime of acceleration of primaries from the PC surface up to the light cylinder implicitly assumes a free supply of charge from the NS.  We do not exclude a scenario with non-stationary or insufficient supply of primaries from the PC surface, in which case there could be a regime of intermittent charge flow. Note that we have briefly discussed a similar possibility in the case of acceleration of primaries over the main part of the PC (see HM01).  Outer gap models (e.g. Cheng et al. 1986, Romani 1996), that assume vacuum gaps form along field lines above the null charge surface, require little or no current flow into the gaps from outside.  In the model presented in this paper, steady current flow can exist from the PC to the light cylinder along even some unfavorably curved field lines where outer gaps are assumed to exist.  In this case the space-charge limited current flow of nearly GJ surface value, combined with cross-field particle motion, would prevent the formation of an outer gap.  In the case of very high inclination angle, where either non-steady current flow or insufficient charge supply from the surface exists, outer gaps may still form (see Section 3.2.1). In this case, however, the radiation which forms caustic emission above the null surface can be viewed from only one pole (Romani \\& Yadigaroglu 1995, Cheng et al. 2000).  The geometry of the extended SG emission is thus clearly distinct from that of the outer gap emission and will have observable consequences, as has been discussed by Dyks \\& Rudak (2003) and Dyks et al. (2003) and will be detailed in future publications."
    },
    "0402/astro-ph0402132_arXiv.txt": {
        "abstract": "We study numerically particle acceleration by the electric field induced near the horizon of a rotating supermassive ($M\\sim 10^9-10^{10}M_{\\odot}$) black hole embedded in the magnetic field $B$. We find that acceleration of protons to energy $E\\simeq 10^{20}$~eV is possible only at extreme values of $M$ and $B$.  We also find that the acceleration is very inefficient and is accompanied by a broad band MeV-TeV radiation whose total power exceeds at least by a factor of 1000 the total power emitted in ultra-high energy cosmic rays (UHECR). This implies that if $O(10)$ nearby quasar remnants were sources of proton events with energy $E>10^{20}$~eV, then each quasar remnant would overshine e.g. the Crab nebula by more than two orders of magnitude in the TeV energy band. Recent TeV observations exclude this possibility. A model in which $O(100)$ sources are situated at $100-1000$~Mpc is not ruled out and can be experimentally tested by present TeV $\\gamma$-ray telescopes. Such a model can explain the observed UHECR flux at moderate energies $E\\sim (4-5)\\times 10^{19}$~eV. ",
        "introduction": "A conventional hypothesis of ultra-high energy cosmic ray (UHECR) acceleration in extragalactic astrophysical objects has two important consequences. First, it predicts the Greisen-Zatsepin-Kuzmin (GZK) cutoff \\cite{GZK} in the spectrum of UHECR at energy of order $5\\times 10^{19}$~eV.  Whether such a cutoff indeed exists in nature is currently an open question \\cite{AGASA,hires}. Second, it implies that the observed highest-energy cosmic rays with $E>10^{20}$~eV should come from within the GZK distance $\\sim 50$~Mpc. Moreover, under plausible assumptions about extragalactic magnetic fields supported by recent simulations \\cite{IGMF}, the propagation of UHE protons over the GZK distance is rectilinear and the observed events should point back to their sources. While sub-GZK UHECR were found to correlate with BL Lacertae objects \\cite{bllacs,EgretBllacs}, no significant correlations of cosmic rays with energies $E \\gsim 10^{20}$ eV with nearby sources were found \\cite{AGASA2}.  In view of the latter problem, a question arises whether there exist UHECR accelerators which can produce super-GZK protons and are quiet in the electromagnetic (EM) channel. If such quiet accelerators existed, they could explain the apparent absence of sources within $\\sim 50$~Mpc in the direction of the highest-energy events. This idea was advocated, e.g., in Ref.~\\cite{dead} where sources of UHE protons were associated with supermassive black holes in quiet galactic nuclei (the so-called ``dead quasars'').  However, it was pointed out \\cite{Levinson:nx} that most of the energy available for particle acceleration in such an environment is spent for EM radiation by the accelerated particles. As a consequence, the flux of TeV $\\gamma$-rays produced by  such an accelerator may be at a detectable level.  Recent observations by HEGRA/AIROBICC \\cite{hegralimit}, MILAGRO \\cite{milagro} and TIBET \\cite{tibet} arrays improved substantially upper limits on the flux of $\\gamma$-rays above $10$ TeV from the point sources in the Northern hemisphere. This may exclude completely a possibility to explain the observed super-GZK cosmic rays by the acceleration near supermassive black holes. The purpose of this paper is to analyze this question quantitatively. To this end we study particle acceleration near the black hole horizon numerically. Following Refs. \\cite{dead,Levinson:nx} we restrict ourselves to the case of protons. The case of heavy nuclei acceleration, propagation and detection is phenomenologically very different and requires separate consideration \\cite{longpaper}. In particular, heavy nuclei can be easily desintegrated already at the stage of acceleration.  We would like to stress that it is not our pupose to construct a realistic model of a compact UHE proton accelerator, but rather to find out whether quiet compact accelerators may exist, even if most favorable conditions for acceleration are provided. To this end we minimize the energy losses of accelerated particles by considering acceleration in ordered electromagnetic field and neglect all possible losses related to scattering of the accelerated particles off matter and radiation present in the acceleration site. However, we take into account, in a self-consistent way, the synchrotron/curvature radiation losses which are intrinsic to the acceleration process.  Clearly, this approximation corresponds to most favorable conditions for particle acceleration. In realistic models the resulting particle energy will be smaller, while emitted EM power larger. Therefore, our results should be considered as a lower bound on the ratio of electromagnetic to UHECR power of a cosmic ray accelerator based on a rotating supermassive black hole.  We find that the flux produced by a nearby UHE proton accelerator of super-GZK cosmic rays in the energy band $E_{\\gamma}>10$~TeV should be at least $100-1000$ times larger than that of the Crab nebula. The existence of such sources is indeed excluded by recent observations \\cite{hegralimit,milagro}.  At the same time, the constraints on the sources of sub-GZK cosmic rays are weaker or absent, see Sect.~VII for details.  The paper is organized as follows. In Sect.~II we describe our minimum loss model in more detail. In Sect.~III we present the analytical estimate and the numerical calculation of maximum particle energy. In Sect.~IV the self-consistency constraints on the parameters of this model are considered which arise from the requirement that there is no on-site $e^+e^-$ pair production caused by emitted radiation.  In Sect.~V the calculation of the EM luminosity of the accelerator is presented.  In Sect.~VI observational constraints are derived. Sect.~ VII contains the discussion of the results and concluding remarks. ",
        "conclusions": "The model of particle acceleration near the horizon of a supermassive black hole which we have discussed in this paper is based on a number of assumptions: maximum rotation moment of the black hole, low matter and radiation density in the acceleration volume, absence of back reaction of the accelerated particles and their radiation on the EM field, uniform magnetic field at large distance from the black hole. These assumptions have one common feature: they facilitate acceleration to higher energies and minimize losses (and, therefore, the produced radiation). We have found that even in these idealized conditions the acceleration of protons to energy $E= 10^{20}$~eV requires extreme values of parameters, $M\\simeq 10^{10} M_{\\odot}$ and $B\\simeq 3\\times 10^4$~G. Moreover, the acceleration is very inefficient: the total power emitted in TeV gamma rays is $100-1000$ times larger than in UHECR. In view of recent TeV observations, this rules out some UHECR models based on this mechanism of acceleration, e.g. the model of several nearby dormant galactic nuclei (``dead quasars'') which was aimed to explain observed UHECR flux with energy $E>10^{20}$~eV.  In a more realistic case the above conditions may not be satisfied completely, and the acceleration of protons to energy $E\\simeq 10^{20}$~eV in the continuous regime may not be possible. The synchrotron losses due to the presence of random component $B_{\\rm rand}$ of the magnetic field can be neglected if \\begin{equation} B_{\\rm rand}\\ll \\frac{B}{\\cal R}\\,\\frac{m}{m_p}\\simeq 10^{-2}\\,\\frac{m}{m_p}\\, B\\; , \\end{equation} where ${\\cal R}$ is given by (\\ref{eq:E_EM/E_UHECR}). This means that the presence of a tiny (1\\% level) random magnetic field will lead to the decrease of maximal energies of accelerated protons and increase of electromagnetic luminosity of the accelerator. Note that the synchrotron radiation is emitted in this case at energies \\begin{equation} \\label{synch1} \\epsilon_{\\rm synch}\\le \\frac{m}{e^2}\\,\\frac{B}{B_{\\rm rand}} \\simeq 0.1\\,\\frac{m}{m_p}\\,\\frac{B}{B_{\\rm rand}} \\mbox{ TeV}\\; . \\end{equation} The power is still given by the Eq. (\\ref{eq:TeV_luminosity}).  Even if the strength of the random component of magnetic field is as small as $10^{-5}\\;B$, for electrons, which are inevitably present in the accelerator, the synchrotron losses will dominate over the curvature losses. The electromagnetic power emitted by electrons will be in the 100~MeV - 10~TeV energy band (see Eq. (\\ref{synch1})). Assuming that the density of electrons is of the same order as the density of protons, one obtains the same estimate (\\ref{eq:TeV_luminosity}) for the 100 MeV luminosity of the accelerator. This means that such an accelerator is not only a powerful TeV source, but is also, in fact, an extremely powerful EGRET source.  Even if the idealized conditions are realized in Nature, corresponding objects must be extremely rare. Thus, only a very small fraction of (active or quiet) galactic nuclei could be stationary sources of UHE protons with energies above $10^{20}$~eV.  If the parameters of the model are not precisely tuned to their optimal values, one expects the maximum energies of accelerated protons to be somewhat below $10^{20}$~eV. It is therefore interesting to note that most of the correlations of UHECR with the BL Lacertae objects come from the energy range $(4-5)\\times 10^{19}$~eV. The central engine of BL Lacs is thought to consist of a supermassive black hole; it is possible that the acceleration mechanism considered above is operating in these objects \\cite{neronov}.  This mechanism may also operate in the centers of other galaxies which may possess (super)massive black holes, including our own Galaxy where it may be responsible for the production of cosmic rays of energies up to $\\sim 10^{18}$~eV \\cite{Levinson:2000zz,longpaper}.  The constraints from TeV observations are different in this case.  First, cosmic rays of lower energies propagate over cosmological distances, so the UHECR flux is collected from a much larger volume, and the number of sources may be larger. Correspondingly, the TeV luminosity of each source is smaller. Second, TeV radiation attenuates substantially over several hundred megaparsecs. Third, at $E<10^{20}$~eV the ratio $L_{\\rm EM}/L_{UHECR}$ is smaller. Consider e.g. the case of $O(100)$ sources located at $z\\sim 0.1$ with typical maximal energy at accelerator $E\\sim 5\\times 10^{19}$ eV. According to Fig.~\\ref{fig:energy}, in this case the typical energy of produced $\\gamma$-rays is $\\sim 4$ TeV. The flux of $\\gamma$-rays in this energy range is attenuated by a factor $10 - 100$, while according to Eq.~(\\ref{eq:E_EM/E_UHECR}), $L_{\\rm EM}/L_{UHECR} \\sim 50$.  Therefore, one may expect $F_{\\rm EM} \\sim (0.01 - 0.1)\\, F_{\\rm Crab}$ for the TeV flux from each of these sources. This is within the range of accessibility of modern telescopes. For example, the TeV flux from the nearby ($z=0.047$) BL Lac 1ES 1959+650, which correlates with arrival directions of UHECR \\cite{EgretBllacs,Gorbunov:2004bs}, is at the level of $0.06F_{\\rm Crab}$ during quiet phase and rises up to $2.9 F_{\\rm Crab}$ during flares. Several other BL Lacs which are confirmed TeV sources have fluxes $\\approx 0.03F_{\\rm Crab}$, see e.g. Ref.~\\cite{Aharonian:2004kd}.  This paper mainly concerns the stationary regime of acceleration when the acceleration volume is not polluted by the creation of $e^+e^-$ pairs. To ensure this condition we required that magnetic field does not exceed the critical value (\\ref{eq:discharge_cond}). If the magnetic field is larger, the acceleration by the mechanism considered here can only happen during flares which are interrupted by the creation of $e^+e^-$ plasma and neutralization of the electric field as discussed in the end of Sect.~\\ref{sect:discharge}. Although we do not have a quantitative model of a flare, some features of this regime and its consequences for UHECR production can be understood qualitatively. Since there is no constraint on the magnetic field in this regime, the maximum energies of accelerated protons may exceed $10^{20}$~eV. However, the efficiency of the acceleration during flares must be much lower than in the stationary case. First, as follows from Fig.~\\ref{fig:ratio}, the ratio $L_{\\rm EM}/L_{\\rm UHECR}$ is larger at large $B$. Second, the dominant part of the EM radiation is produced by the created electrons and positrons whose number density exceeds by far the number density of protons. Thus, one expects that the ratio $L_{\\rm EM}/L_{\\rm UHECR}$ for this sources is much larger than in Eq.~(\\ref{eq:TeV_luminosity}).  An UHECR accelerators operating in the flaring regime would produce approximately constant UHECR flux at the Earth. The reason is the time delay of protons due to random deflections in the extragalactic magnetic fields. This delay is of the order of $\\sim 10^5 [\\alpha/1^{\\circ}]^2$~yr for a source at 100~Mpc, where $\\alpha$ is the deflection angle. Since the time scale of flares (light crossing time) is of the order of day(s), the variations of UHECR flux would average away. On the contrary, the TeV radiation from such a source would be highly variable, with powerful ``TeV bursts'' and {\\em average} energy flux in TeV band exceeding that in UHECR by a factor of $10^4$ or higher. Note that there exist tight constraints on transient TeV sources: the energy flux of a TeV burst of duration $10^5$~s has to be less than $10^{-10}\\mbox{erg/cm$^2$s}\\sim 10\\, F_{\\rm Crab}$ \\cite{milagro,tibet}. As in the case of stationary accelerator, this constraint excludes the possibility to explain observed UHECR flux by a few nearby proton accelerators operating in the flaring regime.  The hypothesis of several hundred remote sources is not constrained by TeV observations.  To summarize, the model of compact UHE proton accelerators which operate near the horizons of supermassive black holes in galactic nuclei can explain only the sub-GZK events. A large number (several hundred) of sources situated at cosmological distances is required.  Production of UHECR in such sources may be associated with the blazar-type activity, TeV $\\gamma$-radiation being an important signature of the model, testable by the existing $\\gamma$-ray telescopes."
    },
    "0402/astro-ph0402418_arXiv.txt": {
        "abstract": "We present X-ray spectra taken with {\\em XMM-Newton} of \\rxj, the third brightest in the class of nearby, thermally emitting neutron stars.  In contrast to what is the case for the brightest object, RX J1856.5$-$3754, we find that the spectrum of \\rxj\\ cannot be described well by a pure black body, but shows a broad absorption feature at 27\\,\\AA\\ (0.45\\,keV).  With this, it joins the handful of isolated neutron stars for which spectral features arising from the surface have been detected.  We discuss possible mechanisms that might lead to the features, as well as the overall optical to X-ray spectral energy distribution, and compare the spectrum with what is observed for the other nearby, thermally emitting neutron stars.  We conclude that we may be observing absorption due to the proton cyclotron line, as was suggested for the other sources, but weakened due to the strong-field quantum electrodynamics effect of vacuum resonance mode conversion. ",
        "introduction": "\\label{sec:introduction}  The nearby, thermally emitting, radio-quiet neutron stars offer the best possibilities for measuring spectra directly from a neutron-star surface, uncontaminated by emission due to accretion and/or magnetospheric processes.  Since the serendipitous discovery of the first of these in \\citeyear{wwn96} by Walter et al., six (possibly seven) more have been uncovered in the ROSAT All-Sky Survey (see reviews by \\citealt{ttz+00,hab04}).  For four sources, optical counterparts have been identified (\\citealt{wm97,mh98,kvk98,kkvk02,kkvk03}).  The high X-ray to optical flux ratios leave no model but an isolated neutron star.  At present, it is not clear what is the source of the thermal emission.  The possibilities considered range from slow accretion from the interstellar medium to release of residual heat, to decay of strong magnetic fields.  Most important for the present purposes, however, is that all six sources appear to have X-ray spectra that, as far as one can tell from current observations, are entirely thermal, thus offering the hope that it will be possible to do a `standard' model-atmosphere analysis, and infer precise values of the temperature, surface gravity, gravitational redshift and magnetic field strength.  Given their interest, the nearby, thermally emitting neutron stars have been among the prime targets for spectroscopy with the {\\em Chandra X-ray Observatory\\/} and {\\em XMM-Newton}.  First {\\em XMM} results on the second-brightest in the class, \\object{RX J0720.4$-$3125} (\\citealt{pmm+01}), were somewhat disappointing, as no lines were found.  The first {\\em Chandra} spectra, of the brightest thermally emitting neutron star, \\object{RX~J1856.5$-$3754}, also revealed no lines (\\citealt{bzn+01}).  Indeed, even a much longer (500~ks) integration failed to show evidence for any features, showing instead a spectrum remarkably well described by a black body affected only by interstellar extinction (\\citealt{dmd+02,br02,bhn+03}). Similarly, {\\em Chandra} data on the \\object{Vela pulsar} (\\citealt{pzs+01}) and \\object{PSR B0656+14} (\\citealt{ms02}), both of which have strong thermal components in their X-ray spectra, failed to show any features.  The only exceptions came recently.  First, \\citet{spzt02} discovered two absorption features in {\\em Chandra} observations of 1E 1207.4$-$5209, the pulsating central source in the supernova remnant PKS 1209$-$51/52.  Variability in these features as a function of pulse phase was seen in {\\em XMM} observations by \\citet{mdlc+02}, and confirmed by \\citet{bclm03}, who also reported a third and possibly even a fourth feature, all harmonically spaced.  The nature of the absorption features is not yet clear, with suggestions ranging from cyclotron lines to transitions of Helium or mid-Z atoms in a strong magnetic field (\\citealt{spzt02,hm02}).  Second, \\citet{hsh+03} discovered a broad absorption feature in {\\em XMM} spectra of RX J1308.6+2127, a nearby, thermally emitting neutron star.  As before, it is not clear what causes the feature; \\citeauthor{hsh+03} speculate it might be due to proton cyclotron absorption.  Third, while we were revising of our manuscript, a preprint by \\citet{hztb03} showed that another nearby neutron star, RX J0720.4$-$3125 had a broad, but weaker absorption feature as well, contrary to earlier claims (\\citealt{pmm+01,pzs+01,kvkm+03}).  Here, we present the discovery of a broad absorption feature in a third nearby, thermally emitting neutron star, \\rxj.  We describe our observations and reduction in \\Sref{data}.  In \\Sref{spectrum}, we analyze and characterize the spectrum, and in \\Sref{timing}, we derive limits to any periodic variations.  We discuss implications for our understanding of \\rxj, as well as the nearby, thermally emitting neutron stars in general, in \\Sref{speculation}.  \\begin{deluxetable}{lllllll} \\tablecaption{Summary of X-ray Observations\\label{tab:log}} \\tabletypesize{\\small} \\tablewidth{0pt} \\tablehead{ \\colhead{ObsID}& \\colhead{Instrument}& \\colhead{Mode}& \\colhead{Start Time}& \\multicolumn{2}{c}{Exposure}& \\colhead{Count Rate\\tablenotemark{a}}\\\\ & & & &\\colhead{Raw} &\\colhead{Filtered} & \\\\ & & &\\colhead{(UT)} &\\colhead{(ks)} &\\colhead{(ks)}& \\colhead{(${\\rm s}^{-1}$)}} \\startdata 2791 & ACIS-I & Standard & 2002~Jan~07.17 & 20.4 & 20.4 & 0.153\\tablenotemark{b}\\\\[0.1in] 0073140301 & EPIC-PN & Timing &       2002~Jan~10.04 & 26.0 &14.3 &3.118(16)\\\\ & RGS1/RGS2 & Standard & & 33.6 & 20.9/19.4 & 0.139(3)/0.130(3)\\\\[0.1in] 0073140201 & EPIC-PN & Timing &       2002~Jan~15.99 & 26.4 &24.2 &3.083(12)\\\\ & RGS1/RGS2 & Standard & & 30.4 & 27.8/27.0 & 0.138(3)/0.127(3)\\\\[0.1in] 0073140501 & EPIC-PN & Timing &       2002~Jan~19.97 & 29.8 &18.4 &3.069(14)\\\\ & RGS1/RGS2 & Standard & & 33.9 & 22.2/21.5 & 0.141(3)/0.133(3)\\\\[0.1in] 0157360401 & EPIC-PN & Large-window & 2003~Jan~17.91 & 33.2 &25.7 &2.385(10)\\\\ & RGS1/RGS2 & Standard & & 41.7 & 31.3/31.2 & 0.144(2)/0.133(2)\\\\[0.1in] 0157360601 & EPIC-PN\\tablenotemark{c} & Large-window & 2003~Feb~26.80 & 16.8 & 8.8 &1.477(13)\\\\ & RGS1/RGS2 & Standard & & 32.2 & 10.1/\\phn9.5& 0.148(5)/0.147(6)\\\\ \\enddata \\tablenotetext{a}{All count-rates are background subtracted and only from the filtered time intervals.  Count-rates for EPIC-pn are for single events with energies $>0.3$~keV.  Numbers in parentheses indicate uncertainties in the last digit.} \\tablenotetext{b}{Piled-up.} \\tablenotetext{c}{Taken with the thick filter.  All other EPIC observations are taken with the thin filter.} \\end{deluxetable} ",
        "conclusions": ""
    },
    "0402/astro-ph0402242_arXiv.txt": {
        "abstract": "CAIRNS (Cluster And Infall Region Nearby Survey) is a spectroscopic survey of the infall regions surrounding nine nearby rich clusters of galaxies.  In Paper I, we used redshifts within $\\sim 10\\Mpc$ of the centers of the clusters to determine the mass profiles of the clusters based on the phase space distribution of the galaxies.  Here, we use 2MASS photometry and an additional \\ncziinew redshifts to investigate the environmental dependence of near-infrared mass-to-light ratios. In the virial regions, the halo occupation function is non-linear; the number of bright galaxies per halo increases more slowly than the mass of the halo.  On larger scales, the light contained in galaxies is less clustered than the mass in rich clusters.  Specifically, the mass-to-light ratio inside the virial radius is a factor of $1.8\\pm0.3$ larger than that outside the virial radius.  This difference could result from changing fractions of baryonic to total matter or from variations in the efficiency of galaxy formation or disruption with environment.  The average mass-to-light ratio $M/L_K = 53\\pm 5 h$ implies $\\Omega _m = 0.18\\pm 0.03$ (statistical) using the luminosity density based on 2dFGRS data.  These results are difficult to reconcile with independent methods which suggest higher $\\Omega _m$.  Reconciling these values by invoking bias requires that the typical value of $M/L_K$ changes significantly at densities of $\\lesssim$3$\\rho_c$. ",
        "introduction": "The relative distribution of matter and light in the universe is one of the outstanding problems in astrophysics.  Clusters of galaxies, the largest gravitationally relaxed objects in the universe, are important probes of the distribution of mass and light. \\citet[][]{zwicky1933} first computed the mass-to-light ratio of the Coma cluster using the virial theorem and found that dark matter dominates the cluster mass.  Recent determinations using the virial theorem yield mass-to-light ratios of $M/L_{B_j}\\sim 250 \\mlsun$ \\citep[][and references therein]{g2000}.  Equating the mass-to-light ratio in clusters to the global value provides an estimate of the mass density of the universe \\citep{oort58}; this estimate is subject to significant systematic error introduced by differences in galaxy populations between cluster cores and lower density regions \\citep{cye97,g2000}.  Indeed, some numerical simulations suggest that cluster mass-to-light ratios exceed the universal value \\citep[][but see also Ostriker et al.~2003]{diaferio1999,kk1999,bahcall2000,benson00}.  Determining the global matter density from cluster mass-to-light ratios therefore requires knowledge of the dependence of mass-to-light ratios on environment.  \\citet{bld95} show that mass-to-light ratios increase with scale from galaxies to groups to clusters.  Ellipticals have larger overall values of $M/L_B$ than spirals, presumably a result of younger, bluer stellar populations in spirals. At the scale of cluster virial radii, mass-to-light ratios appear to reach a maximum value.  Some estimates of the mass-to-light ratio on very large scales ($>$10$\\Mpc$) are available \\citep[see references in][]{bld95}, but the systematic uncertainties are large.  There are few estimates of mass-to-light ratios on scales between cluster virial radii and scales of 10$\\Mpc$ \\citep[][]{elt97,small98,kaiserxx,rines2000,rines01a,bg03,katgert03,kneib03}. On these scales, many galaxies near clusters are bound to the cluster but not yet in equilibrium \\citep{gunngott}.  These cluster infall regions have received relatively little scrutiny because they are mildly nonlinear, making their properties very difficult to predict analytically.  However, these scales are exactly the ones in which galaxy properties change dramatically \\citep[][and references therein]{ellingson01,lewis02,gomez03,treu03,balogh03}.  Variations in the mass-to-light ratio with environment could have important physical implications; they could be produced either by a varying dark matter fraction or by variations in the efficiency of star formation with environment.  In blue light, however, higher star formation rates in field galaxies compared to cluster galaxies could produce lower mass-to-light ratios outside cluster cores resulting only from the different contributions of young and old stars to the total luminosity \\citep{bahcall2000}.  Because clusters are not in equilibrium outside the virial radius, neither X-ray observations nor Jeans analysis provide secure mass determinations at these large radii. There are now two methods of approaching this problem: weak gravitational lensing \\citep{kaiserxx} and kinematics of the infall region \\citep[][hereafter DG97 and D99]{dg97,diaferio1999}. \\citet{kaiserxx} analyzed the weak lensing signal from a supercluster at $z \\approx 0.4$; the mass-to-light ratio ($M/L_B$=280 $\\pm$40 $h$ for early-type galaxy light) is constant on scales up to $6~\\Mpc$. \\citet{wilson01} finds similar results for weak lensing in blank fields; \\citet{gray02} obtain similar results for a different supercluster.  Recently, \\citet{kneib03} used weak lensing to estimate the mass profile of CL0024+1654 to a radius of 3.25$\\Mpc$.  \\citet{kneib03} conclude that the mass-to-light ratio is roughly constant on these scales.  Galaxies in cluster infall regions produce sharp features in redshift surveys \\citep{kg82,shectman82,dgh86,kais87,ostriker88,rg89}. Early investigations of this infall pattern focused on its use as a direct indicator of the global matter density $\\Omega_m$. Unfortunately, random motions caused by galaxy-galaxy interactions and substructure within the infall region smear out this cosmological signal \\citep[DG97, ][]{vh98}.  Instead of sharp peaks in redshift space, infall regions around real clusters typically display a well-defined envelope in redshift space which is significantly denser than the surrounding environment \\citep[][hereafter Paper I, and references therein]{cairnsi}.  DG97 analyzed the dynamics of infall regions with numerical simulations and found that in the outskirts of clusters, random motions due to substructure and non-radial motions make a substantial contribution to the amplitude of the caustics which delineate the infall regions \\citep[see also][and references therein]{vh98}. DG97 showed that the amplitude of the caustics is a measure of the escape velocity from the cluster; identification of the caustics therefore allows a determination of the mass profile of the cluster on scales $\\lesssim 10\\Mpc$.  DG97 and D99 show that nonparametric measurements of caustics yield cluster mass profiles accurate to $\\sim$50\\% on scales of up to 10 $h^{-1}$ Mpc.  This method assumes only that galaxies trace the velocity field. Indeed, simulations suggest that little or no velocity bias exists on linear and mildly non-linear scales \\citep{kauffmann1999a,kauffmann1999b}.  \\citet[][hereafter GDK]{gdk99}, applied the kinematic method of D99 to the infall region of the Coma cluster.  GDK reproduced the X-ray derived mass profile and extended direct determinations of the mass profile to a radius of $10~\\Mpc$.  The caustic method has also been applied to the Shapley Supercluster \\citep{rqcm}, A576 \\citep[][hereafter R00]{rines2000}, AWM7 \\citep{kg2000}, the Fornax cluster \\citep{drink}, A1644 \\citep{tustin}, A2199 \\citep{rines02}, and six other nearby clusters (Paper I).  \\citet{bg03} applied the caustic technique to an ensemble cluster created by stacking redshifts around 43 clusters from the 2dF Galaxy Redshift Survey.  R00 found an enclosed mass-to-light ratio of $M/L_R \\sim 300 h$ within 4$~\\Mpc$ of A576.  \\citet{rines01a} used 2MASS photometry and the mass profile from GDK to compute the mass-to-light profile of Coma in the K-band.  They found a roughly flat profile with a possible decrease in $M/L_K$ with radius by no more than a factor of 3.  \\citet{bg03} find a decreasing ratio of mass density to total galaxy number density.  For early-type galaxies only, the number density profile is consistent with a constant mass-to-light (actually mass-to-number) ratio.  Here, we calculate the infrared mass-to-light profile within the turnaround radius for the CAIRNS clusters (Paper I), a sample of nine nearby rich, X-ray luminous clusters.  We use photometry from the Two Micron All Sky Survey \\citep[2MASS,][]{twomass} and add several new redshifts to obtain complete or nearly complete surveys of galaxies up to 1-2 magnitudes fainter than $M^*_{K_s}$ \\citep[as determined by][hereafter K01 and C01]{twomasslfn,twomdflfn}.  Infrared light is a better tracer of stellar mass than optical light \\citep[][]{gpb96,zibetti02}; it is relatively insensitive to dust extinction and recent star formation.  Despite these advantages, there are very few measurements of infrared mass-to-light ratios in clusters \\citep{tustin,rines01a,lin03}.  Mass-to-light ratios within virial regions (where the masses are more accurate than in the infall regions) provide interesting constraints on the distribution of dark matter and stellar mass \\citep[see also][hereafter L03]{lin03}.  The virial masses in our sample span an order of magnitude in mass. More massive clusters have larger mass-to-light ratios.  Cluster virial regions also provide potentially important constraints on the halo occupation distribution \\citep[e.g.,][and references therein]{peacock00,berlind02,berlind03}, the number of galaxies in a halo of a given mass \\citep[see ][for a recent review]{2002PhR...372....1C}.  The motions of galaxies and hot gas yield estimates of the dynamical mass independent of the number of galaxies (provided enough galaxies are present to yield a virial mass).  Our mass profiles in Paper I are among the first to extend significantly beyond $r_{200}$.  Thus, they should provide accurate estimates of $r_{200}$.  Also, the recent release of 2MASS allows us to count galaxies based on their near-infrared light, which is close to selecting galaxies by stellar mass.  Thus, both the masses and galaxy numbers are better defined than the few previous direct estimates of the halo occupation function \\citep{peacock00,marinoni02,lin04}.  We describe the cluster sample, the near-infrared photometry, and the spectroscopic observations in $\\S$ 2. We discuss the galaxy properties (luminosity functions and broadband colors) within and outside the virial radius and compare both populations to field galaxies in $\\S 4$.  We calculate the number density and luminosity density profiles in $\\S 5$ and compare them to simple theoretical models.  We compute radial profiles of the mass-to-light ratio in $\\S 5$.  In $\\S 6$ we constrain the halo occupation distribution for the CAIRNS clusters and explore the dependence of mass-to-light ratios on halo mass.  We discuss possible systematic uncertainties and the implications of our results in $\\S 7$ and conclude in $\\S 8$.  We assume a cosmology of $H_0 = 100~h~\\kms, \\Omega_m = 0.3, \\Omega _\\Lambda = 0.7$ except as noted in $\\S 7$. ",
        "conclusions": "We discuss some of the first estimates of radial variations in mass-to-light ratios on scales of 1-10$\\Mpc$ using near-infrared photometry from 2MASS and mass profiles from the kinematics of infalling galaxies.  Because cluster infall regions contain the transition from cluster galaxies to field galaxies \\citep[][and references therein]{ellingson01,lewis02,gomez03,treu03,balogh03}, mass-to-light ratios in infall regions should closely resemble the global value.  To summarize our results: \\begin{itemize} \\item{Infall regions contain more bright galaxies (to a fixed absolute magnitude limit) than cluster virial regions.} \\item{The near-infrared luminosity functions for bright galaxies ($M_{K_s}\\lesssim -22 + 5 \\mbox{log} h$) in the CAIRNS cluster virial regions and infall regions do not differ significantly from the field galaxy luminosity function.  Clusters contain an excess of extremely bright galaxies above the predictions of a Schechter function.} \\item{Optical-near-infrared colors in A576 show no radial dependence. This lack of a color gradient shows that the stellar populations do not change dramatically with radius.  It is likely that the mild gradients found in the photometric study of \\citet{goto04} may be enhanced in optically selected samples as compared to near-infrared selected samples such as CAIRNS.} \\item{Galaxies in cluster virial regions and infall regions exhibit a near-infrared color-magnitude relation with a shallower slope than at optical wavelengths.  These galaxies also exhibit little scatter in $J-K_s$ colors, indicating that the stellar populations are fairly homogeneous and that internal dust extinction and/or emission is important for only a few galaxies. } \\item{Both the surface number density profiles and surface luminosity density profiles of CAIRNS members indicate that galaxies and stellar light are more extended than mass. } \\item{Near-infrared mass-to-light ratios generally decrease with radius by a factor of $1.8\\pm0.3$ in the infall regions of the CAIRNS clusters.  This result agrees with previous results based on individual clusters and optical photometry. The presence of decreasing mass-to-light profiles even at $K_s$ band suggests that the decrease is not due to changes in stellar populations.} \\item{Near-infrared mass-to-light ratios calculated at $r_{200}$ using caustic mass estimates agree quite well with mass-to-light ratios calculated at $r_{500}$ from X-ray mass estimates.  This agreement suggests that the decreasing mass-to-light profiles are not monotonic; the mass-to-light ratio is roughly constant inside $r_{200}$.} \\item{We derive some of the first constraints on the halo occupation function using cluster masses and near-infrared selected galaxy samples. The number of bright galaxies $N_{200}$ projected within $R_{200}$ increases as $N_{200} \\propto M_{200}^{0.70\\pm0.09}$, significantly shallower than $N_{200}\\propto M_{200}$.  Earlier studies of the halo occupation distribution suggest that a halo occupation distribution shallower than $N_{200}\\propto M_{200}$ is necessary to reproduce the observed clustering properties of galaxies \\citep[e.g.,][]{berlind02,berlind03}.} \\item{No such non-linear relation is evident between $N_{200}$ and $L_{200}$, the $K_s$ band luminosity inside $R_{200}$.  This result shows that the non-linearity of the halo occupation function is not driven by variations in the luminosity function.} \\item{More massive virialized halos have larger mass-to-light ratios. This result follows logically from the two prior points.  Our $M/L-M$ relation agrees with previous determinations \\citep[][L03]{bahcall02}. These results signify that the efficiency of galaxy formation decreases (and/or that the efficiency of galaxy disruption increases) with increasing halo mass and/or virial temperature.} \\item{We investigate possible systematic effects and conclude that dark matter is more concentrated than stellar mass {\\it contained in galaxies}.  This result could arise either from different clustering properties of dark matter and baryonic matter or from variations in the efficiency of converting baryonic matter into galaxies. The cluster environment seems to be either less efficient at converting baryons into galaxies or more efficient at disrupting galaxies than less dense environments.  Such a difference is predicted by simulations of $\\Lambda$CDM cosmologies where processes such as tidal stripping and dynamical friction disrupt galaxies in clusters \\citep{kk1999,1999ApJ...523...32C} and supported by observations of significant numbers of intergalactic stars in clusters. Alternatively, the heating of the intracluster medium may cut off the supply of cold material needed to form stars \\citep[e.g.,][ and references therein]{1999ApJ...522..590B,bnm2000}, thus lowering the star formation efficiency in cluster galaxies.}  \\item{Assuming the mass-to-light ratios at large radii are similar to the global value, we estimate $\\Omega_m = 0.10\\pm0.03$ (1-$\\sigma$ statistical uncertainty) using the SDSS luminosity density with appropriate color corrections or $\\Omega_m = 0.13\\pm0.03$ (1-$\\sigma$ statistical uncertainty) from the 2dFGRS.  We suggest that the 2dFGRS and the CfA/SSRS2 surveys sample local underdensities.  Uncertainties in the luminosity density, especially at infrared wavelengths, contribute a significant amount of the systematic uncertainty in estimating $\\Omega_m$.  These estimates of $\\Omega_m$ are small compared with other recent estimates from the microwave background, the galaxy power spectrum, and supernovae.  However, they agree well with other estimates based on cluster mass-to-light ratios \\citep{cnoc96,cye97,bahcall2000,g2000,bahcall02}, cluster abundances \\citep[][but see Schuecker et al.~2003]{hiflugcs,bahcall03a} and weak lensing \\citep{kaiserxx,wilson01,hoekstra01,gray02}.  We discuss possible systematic effects that could cause our result to be anomalously low. Reconciling these estimates of $\\Omega_m$ by invoking bias requires that the typical value of $M/L_K$ at the smallest densities we probe $\\approx 3 \\rho_c$ is a factor of 2-3 smaller than the global value. For instance, if galaxy formation occurs nearly exclusively above a density threshold $\\delta\\sim 10$, the mass-to-light ratios in cluster outskirts may underestimate the global value. }  \\end{itemize}  One promising future direction is to study clusters at moderate redshifts where weak lensing provides an independent mass estimate \\citep{kneib03}.  Comparing lensing mass profiles to caustic mass profiles will constrain unknown systematics in both techniques.  For instance, a sheet of mass of uniform density produces no lensing signal (the mass-sheet degeneracy), but this mass should be evident in the galaxy kinematics.  Conversely, foreground and background structures may produce a weak lensing signal but would not affect the kinematics of the infall region.  Infall regions are interesting environments for studying the evolution of galaxy populations.  We show here that infall regions are important in constraining models of galaxy bias or antibias and $\\Omega_m$.  If other methods yield a precise measurement of $\\Omega_m$, the changes in mass-to-light ratios with environment provide important clues to the formation and evolution of galaxies."
    },
    "0402/astro-ph0402597_arXiv.txt": {
        "abstract": "In these lectures, the properties of dense hadronic and quark matter and its relation to compact stars will be discussed.  In a bottom--up approach we start with nuclear and hypernuclear physics at low density and extrapolate hadronic matter to large densities. The matching to the quark matter phase is performed in a top--down approach starting at asymptotically large densities. Implications for the mass--radius relation of compact stars and the existence of a new family of solutions will be outlined. ",
        "introduction": "Neutron stars are created by supernovae type II explosions and are the final endpoint of evolution of massive stars. The compact remnants of the core collapse supernovae have masses in the range of 1--2 solar masses and radii of the order of 10 km. The interior of neutron stars consists of matter under extreme densities, several times the density of normal nuclear matter, $n_0 = 3 \\cdot 10^{14}$ g/cm$^3$. The study of neutron stars has considerably advanced during the last years. With new telescopes, ground-based and in satellites, one measures spectra of supernova remnants like the crab nebula not only in the optical but also in the x-ray (Chandra, XMM--Newton), in radio as well as in the infrared band. The Hubble Space Telescope and the Chandra Telescope has even published a movie of the crab nebula on--line! These movies demonstrate how the rotating neutron star, the crab pulsar, pushes out energetic wisps into the crab nebula in the equatorial plane as well as jets of matter along the polar axis.  More than 1000 pulsars are known today. The masses of pulsars were measured most precisely from a few binary neutron star systems \\cite{Thorsett99}, especially from the Hulse-Taylor pulsar with $M=(1.4411\\pm0.00035)M_\\odot$. The shortest known rotation period so far measured is 1.557 ms for the pulsar PSR 1937+21.  Our knowledge about compact stars got a new twist by the discovery of isolated, non-pulsating neutron stars. The first one seen and the closest one known is RX J1856 \\cite{Walter2001} being radio-quiet and with no pulsations. The thermal spectrum indicates a temperature of 49 eV in the optical band. The x-ray spectrum, however, shows a nice Planck curve with a temperature of 60 eV as measured by the Chandra satellite \\cite{Drake02}. Most surprisingly, no spectral lines from elements in the atmosphere of the neutron star have been found in the spectra! These puzzles of the spectra of the isolated neutron star are still not fully resolved (see \\cite{Burwitz02} for a possible explanations). The revised and improved parallax measurement of RX J1856 gives a distance of $D=117\\pm 12$ parsec \\cite{Walter02}. An ideal black-body emitter would have a very small radius of $R_\\infty = 4-8$ km at that distance which is much smaller than the canonical value for a neutron star of 10 km \\cite{Drake02}. Corrections from an atmosphere resulted in an apparent radius of $R_\\infty = 15\\pm 3$ km which would be compatible with most of the modern neutron star models on the market \\cite{Walter02,Pons2002}.  The structure of neutron stars encompasses several distinctly different zones when going from the surface to the centre (for a review see e.g.\\ \\cite{Haensel03}). First, there is a thin atmosphere up to about $10^4$ g/cm$^3$, mainly iron but could be also hydrogen or helium by accretion. Then the outer crusts or the envelope begins which consists of free electrons and nuclei forming a Coulomb lattice. The sequence of nuclei starts with iron and then continues to neutron-rich nuclei stopping at the neutron drip-line at about $10^{11}$ g/cm$^3$ \\cite{BPS}. The inner crusts consists of free neutrons in addition to nuclei and electrons where the neutrons are in a superfluid state. The nuclei can now form the pasta in the crust, various inhomogeneous phase structures as bubbles (meat-balls), rods (spaghetti), and plates (lasagna) immersed in the neutron and electron fluid. The sequence can then reverse, so that the neutron fluid forms the geometrical structures immersed in a background of extremely neutron-rich and superheavy nuclei \\cite{Negele73}.  At about half times normal nuclear matter density, the matter distribution will be uniform, consisting of mainly neutrons with a small admixture of protons and electrons, where the protons can now form a superconductor. Still, the end of the crust is located far away from the centre, as the crust is only a few hundred meters thick for massive neutron stars!  The core of the neutron star consists of matter under extreme densities. New forms of matter have been proposed to exist under these condition in the very core of compact stars: Bose condensation of pions or kaons, phase transitions to hyperon matter (hyperon star) and quark matter (hybrid star). If strange quark matter is absolutely stable, then the corresponding compact star is dubbed a strange star and the quark phase extends all the way from the centre to the neutron-drip line, as all free neutrons are swallowed by the true ground state of matter, while nuclei are saved by virtue of the Coulomb barrier.  Now let us start with a simple consideration: let us assume that the high-density equation of state can be described by free (known) particles. The stable baryons known in vacuum are the nucleons (n,p) and hyperons (the $\\Lambda$ and $\\Sigma^{-,0,+}$ with one strange quark, the $\\Xi^{-,0}$ with two strange quarks, and the $\\Omega^-$ with three strange quarks). The masses of the hyperons increases with the number of strange quarks. Besides those, there are spinless mesons which are stable against strong interactions (charged pions, and kaons with one anti-strange or one strange quark). They can form a Bose condensate. A calculation of neutron star matter for a free gas of particles shows, that first the negatively charged $\\Sigma^-$ appears at about $4n_0$ and then the $\\Lambda$ at about $8n_0$ \\cite{Ambart60}. The $\\Sigma^-$ appears before the $\\Lambda$ despite the fact that it is slightly heavier, because it is negatively charged. The presence of the $\\Sigma^-$ can then take over the r\\^ole of the electron in making the overall matter charge neutral, thereby lowering the Fermi energy of electrons and the total energy of the system. One can compute that there will be no other particles in the composition of a free gas of hadrons up to $20n_0$ which is well beyond the maximum density in the interior of a compact star (besides that our approach will be not applicable anymore, as hadrons are composite particles). The corresponding equation of state will be even slightly softer than that of a a free neutron gas due to the appearance of additional particles.  Hence, the maximum mass is the one given by Oppenheimer and Volkoff for a free gas of neutrons which is about $0.7M_\\odot$ \\cite{OV39}. That maximum mass is more than a factor two smaller than required by the measurement of the mass of the Hulse-Taylor pulsar of $1.44 M_\\odot$!  Neutron stars can {\\em not} be described by a approximately free gas of particles, contrary to white dwarfs. Interactions between hadrons are crucial for explaining such a large neutron star mass. ",
        "conclusions": ""
    },
    "0402/astro-ph0402074_arXiv.txt": {
        "abstract": " ",
        "introduction": "IP~Pegasi is a deeply eclipsing dwarf nova ($P_{orb}= 3.8$~hr) which shows 1-2 weeks-long, $\\simeq 2$~mag outbursts every 60-120 days. During outbursts, tidally induced spiral shocks form in its accretion disc as the disc expands and its outer parts feel more effectively the gravitational attraction of the companion star (Steeghs, Harlaftis \\& Horne 1997). Here we present the first results of a time-resolved ultraviolet (UV) spectral mapping experiment of IP~Pegasi 9-13 days after onset of the May 1993 outburst. ",
        "conclusions": ""
    },
    "0402/astro-ph0402505_arXiv.txt": {
        "abstract": "SMA observations of the massive star-forming region IRAS\\,18089-1732 in the 1\\,mm and 850\\,$\\mu$m band reveal outflow and disk signatures in different molecular lines. The SiO(5--4) data show a collimated outflow in the northern direction. In contrast, the HCOOCH$_3$(20--19) line, which traces high-density gas, is confined to the very center of the region and shows a velocity gradient across the core. The HCOOCH$_3$ velocity gradient is not exactly perpendicular to the outflow axis but between an assumed disk plane and the outflow axis. We interpret these HCOOCH$_3$ features as originating from a rotating disk that is influenced by the outflow and infall. Based on the (sub-)mm continuum emission, the mass of the central core is estimated to be around 38\\,M$_{\\odot}$. The dynamical mass derived from the HCOOCH$_3$ data is 22\\,M$_{\\odot}$, of about the same order as the core mass. Thus, the mass of the protostar/disk/envelope system is dominated by its disk and envelope. The two frequency continuum data of the core indicate a low dust opacity index $\\beta \\sim 1.2$ in the outer part, decreasing to $\\beta \\sim 0.5$ on shorter spatial scales. ",
        "introduction": "Unambiguous proof for disks in massive star formation is still missing. Millimeter continuum observations suggest flattened structures without providing velocity information (e.g., \\citealt{shepherd2001}), and molecular line studies suggest rotational motions but are often confused outflows and ambient gas (e.g., \\citealt{zhang1998} and Beuther et al., this volume). Maser studies show disk signatures in some cases but are mostly not unambiguous as well (e.g., \\citealt{churchwell2002}). The best evidence yet for genuine disk emission comes from CH$_3$CN observations in IRAS\\,20126+4104 \\citep{cesaroni1999}. In this case, the velocity gradient defining the presence of the disk is aligned perpendicular to the bipolar outflow, consistent with the common disk/jet paradigm.  To further investigate possible disk emission and its association with molecular jets, we used the Submillimeter Array (SMA) to observe the jet tracer SiO(5--4) and the hot-core tracer HCOOCH$_3$(20--19) in a massive star-forming region.  The source IRAS\\,18089-1732 is a young High-Mass Protostellar Object (HMPO) which has been studied in detail over recent years. The source is part of a sample of 69 HMPOs selected mainly via infrared color-color criteria and the absence of strong cm emission \\citep{sridha}. IRAS\\,18089-1732 is approximately at a distance of 3.6\\,kpc\\footnote{The kinematic distance ambiguity is solved by associating the region via the near- and mid-infrared surveys 2MASS and MSX on larger scales with sources of known distance (Bontemps, priv. comm.).} and its bolometric luminosity is about $10^{4.5}$\\,L$_{\\odot}$ \\citep{sridha}. Millimeter continuum observations reveal a massive core $>2000$\\,M$_{\\odot}$ with H$_2$O and CH$_3$OH maser emission, and a weak 1\\,mJy source is detected at 3.6\\,cm \\citep{beuther2002a,beuther2002c}. As part of a single-dish CO outflow study, wing emission indicative of molecular outflows was detected but the CO map was too confused to define a bipolar outflow \\citep{beuther2002b}. During these observations, \\citet{beuther2002b} also observed SiO(2--1) at 3\\,mm, and bipolar structure was detected in the north-south direction. Furthermore, \\citet{sridha} reported the detection of the hot-core-tracing molecules CH$_3$CN and CH$_3$OH. This letter focuses on the jet/disk observations and the (sub-)mm continuum data. A description of the line forest observed simultaneously is presented in an accompanying paper (Beuther et al., this volume). ",
        "conclusions": "The combined SiO(5--4) and HCOOCH$_3$(20--19) observations toward IRAS\\,18089-1732 support a massive star formation scenario where high-mass stars form in a similar fashion as their low-mass counterparts, i.e., via disk accretion accompanied by collimated jets/outflows. The HCOOCH$_3$ observations still barely resolve the disk/envelope structure and the interpretation is not entirely unambiguous, but the data indicate rotation which might stem at least partly from an accretion disk. Higher resolution observations in different disk tracers are needed to investigate the disk/envelope conditions in more detail. Furthermore, the continuum data indicate a lower $\\beta$ than the standard value of 2, and we observe a decreasing $\\beta$ with decreasing spatial scales. As nearly all the gas observed in dust emission takes part in the assumed dynamical rotation observed in HCOOCH$_3$ this low $\\beta$ might be due to grain growth or high opacity within a disk-like structure (e.g., \\citealt{beckwith2000})."
    },
    "0402/astro-ph0402219_arXiv.txt": {
        "abstract": "{High resolution broad and narrow band images and long slit spectroscopy of the peculiar galaxy IC~1182 are  presented. The analysis of the broad band images  reveals a distorted morphology with a large, heavily obscured disk-like structure and several knots in the central region. Galactic material, some of it in the form of two slender tails, is detected well beyond the main body of the galaxy. The second, fainter tail and several knots are reported here for the first time. The galaxy has color indices of an early type object except U$-$B, which is significantly bluer than what is typical for this kind of galaxy.  The narrow band images centered on different emission lines show that the galaxy is a very powerful emitter. Most of  the knots detected in the central region and in the prominent tail emerging eastward from the galaxy are very luminous in H$\\alpha$, and have typical sizes about 1 kpc (FWHM). The emission in the main lines extends all over the galaxy, with plumes and arc-like structures seen in H$\\alpha$ at large distances from the center. The observed, uncorrected H$\\alpha$ flux corresponds to a total luminosity of 3.51$\\times$10$^{41}$ erg s$^{-1}$, about 3 times that of the starburst galaxy Arp~220. We have found that the internal extinction deduced from the observed Balmer decrement is high all along the slit, with E$_{B-V}\\approx$ 1, so the corrected SFR could amount to 90~M$_{\\odot}$ per year. On this basis IC~1182 is found to be a very powerful starburst galaxy. Surprisingly, the source is not in the IRAS Point Source Catalogue.  The emission knots detected in the central region of the galaxy have line ratios that place them close to the border of the region occupied by active nuclei in the diagnostic diagrams. Using the best determined diagnostic ratio, [OIII]/H$\\beta$ {\\sl vs}  [NII]/H$\\alpha$, they can still be classified as extreme HII-like regions. We notice that the same kind of line ratios are also measured at different places in the galaxy, adding to the idea that the nuclear line ratios can be explained in terms of stellar photoionization. The metallicity we have measured for the ionized gas in the two brightest central knots is low, 0.1~Z$_{\\odot}$ and 0.06~Z$_{\\odot}$ respectively, and their measured helium abundance is also lower than solar.  In the main body of the galaxy, besides the reported knots, the distribution of the ionized gas resembles that of an inclined disk about 12 kpc in size. The spectroscopic data show however a complex rotation pattern. We interpret them  as corresponding to two identifiable disk galaxies with observed rotation amplitudes of 200 \\kms and 100 \\kms respectively. The stellar absorption lines detected in the bigger system do not show any clear rotation pattern.  The data presented here indicate that IC~1182 is a high luminosity starburst system. Its global properties and peculiarities can be understood as corresponding to two systems that can still be recognized, in the process of merging, with two tidal tails emerging from the central region of the galaxy. In the main tail there are several candidates forming tidal dwarf galaxies. The measured low metallicity of the ionized gas, together with the low amplitude of one of the systems, suggests that the process involves a late-type, gas-rich spiral galaxy that is supplying most of the gas to the system.  \\keywords {Galaxies: interactions -- Galaxies: star burst -- Galaxies: active -- Galaxies: dwarfs}  } ",
        "introduction": "IC~1182 (Mkr~298) is a bright, early-type (S0p) galaxy belonging to the Hercules cluster. It is placed at the end of the stream of galaxies running from the center of the cluster (defined by its brightest galaxies, NGC~6041A, B) toward the East. Its peculiar morphology is apparent on DSS images, where a long, linear, knotty structure extending 88\\arcsec~ from the main body of the galaxy to the E is clearly visible. In the picture by Arp (1972) diffuse material associated with the galaxy is also visible toward the NW. These are clear signs of interaction. The impression that the galaxy is involved in some interaction or merging process is reinforced by the discovery reported by Rafanelli et al. (1999) of  two bright knots separated by $\\approx$~3\\farcs3, at PA $\\approx$ 130 in the central region of the galaxy.  IC~1182 has some other peculiarities. Bothun et al. (1981) found that it has total colors that correspond to an early type object, except U$-$B, that is too blue for an E or S0 galaxy. They also reported (see also  Salpeter  \\& Dickey 1985) that IC~1182 has an unusually large HI mass fraction for an early type galaxy. Their data reveal the contribution to the observed HI profile of a close companion, most likely interacting with IC~1182, but even after taking into account that contamination, IC~1182 is still very HI-rich, with log(M$_{HI}$/L$_B)\\approx-$0.5. On the other hand, IC~1182 is not associated with a powerful radio-continuum source. Dickey and Salpeter (1984) reported the  detection at 1415~MHz of a rather weak source, probably extended, at the position of IC~1182.  Another sign of peculiarity of IC~1182 is the sizeable polarization it shows (p = 1.09\\%, at $\\theta$= 104.5\\degr, Martin et al. 1983). It is also an  X-ray source. Huang and  Sarazin (1996) reported detection of an individual source coincident with IC~1182 when making high resolution ROSAT observations of the Hercules cluster. Rafanelli et al. (1999) pointed out that the soft X-ray luminosity of the nucleus of IC~1182 is close to the mean value for Seyfert 2 galaxies. They also argued that the existing data on the X-ray variability of the source would imply a source size too small to contain the supernovae necessary to produce the observed luminosity. This would be strong evidence for the presence of an active nucleus at the center of IC~1182.  The nature of the nucleus of IC~1182 has been the subject of some discussion. It was first classified as Seyfert 2 by Khachikian and Weedman (1974), but the line ratios reported by Koski (1978) are closer to those of a LINER (see Viegas-Aldrovandi  \\&  Gruenwald 1990). Heckman et al. (1981) reported that the observed line profiles are due to the superposition of several components that are spatially and kinematically distinct, and V\\'eron et al. (1997) classify it as a composite spectrum object.  The reported properties and the high internal extinction and high H$\\alpha$ luminosity we find in the present work (see below) should correspond to a high IR luminosity galaxy. Surprisingly, the only IR source detected by IRAS at the position of IC~1182 is one close to the bright knot at the end of the long tail emerging toward the East of the main body of the galaxy, with  f$_{60}$ = 240~mJy in the IRAS PSC (Beichman et al. 1986), corresponding to L$_{FIR}\\sim$ 10$^7$ L$_{\\odot}$ at the distance of the galaxy.  We present here new photometric and spectroscopic data for IC~1182 gathered with the NOT (La Palma), the  1.54m  Danish telescope (La Silla), and the VLT Antu (U1, Paranal). These include broad and narrow band images under good to excellent seeing conditions, long slit spectra with three different slit orientations, and MOS observations of the detected knots. The new data reveal the presence of several knots in a highly obscured central region, a second, faint tail almost perpendicular to the main one, a complex rotation pattern of the ionized gas, and evidence for a generalized star formation process all over the galaxy. These peculiarities lead us to classify IC~1182 as a starburst galaxy, resulting from an ongoing merging process involving two spiral galaxies, one of them gas rich and of rather late type. ",
        "conclusions": "The global morphology of IC~1182 is reminiscent of an early type galaxy, but with important peculiarities. In the continuum light, IC~1182 shows an entangled structure, with $\\varepsilon$ = 0.25 and PA = 80 when the outermost isophotes are considered. A long tail, 63 kpc long, emerges from the central region to the E, at PA = 97, with a knotty structure. We  have detected a new, much fainter tail running almost perpendicular to the previous one, and stretching for 27 kpc from the center. It is continued by a diffuse, large S-shaped structure toward the NW, already visible in the photographic picture by Arp (1972). In the model subtracted image we have constructed, one sees a large, disk-like dust structure, already hinted in the B-band image, with a size of 5~kpc. Finally, to the SW of the dust structure, there is a residual structure that appears distinctly in the model-subtracted image. These characteristics, in particular the discovery of the Second Tail, strongly suggest that IC~1182 could be a merging system. Moreover, several knots along the main tail, detected in the broad band and H$\\alpha$ images, are resolved. They could be forming tidal dwarf galaxies as seen in the Antennae (Mirabel et al. 1992) or in Stephan's Quintet (Mendes de Oliveira et al. 2001), as recently suggested by van Driel et al. (2003) for the knot at the tip of MT.  The narrow band images centered in the main emission lines show a galactic size distribution of ionized gas, with plumes and filaments stretching several kpc from the central regions. The amount of ionized gas could be as high as 10$^6$ or   even 10$^7$ M$_{\\odot}$. It is organized in an inclined (i =  68; PA = 120) rotation disk with a major axis of about 12 kpc.  The velocity distribution we observed at two slit positions can be interpreted as produced by two distinct rotating systems. The main system would dominate the dynamics of most of the observed ionized gas, with a rotation amplitude of $\\approx$ 230 \\kms. The mass responsible for that rotation would amount to 2$\\times$10$^{10}$ M$_{\\odot}$. The G parameter amounts to $\\approx$ 300 \\kms kpc$^{-1}$. In that region, the stars don't show any clear rotation pattern, and have a velocity dispersion $\\sigma$ = 200 \\kms. The rotation amplitude and G value are typical of the bulge of a Sb galaxy (M\\'arquez  \\& Moles 1999). The second system would be less massive, with an observed rotation amplitude of 100 \\kms.  There are no clear signs of nuclear activity. The measured line ratios of the brightest knots are still compatible with photoionization by stars. We also notice that, given the complex structure of the central region, the disagreement about    the line ratios reported in earlier work, could simply reflect the differences in slit width and positioning.  Rafanelli et al. (1999) have reported rapid variability in X-rays, indicative of the presence of a very compact source. In principle this would support the idea of the existence of some kind of nuclear activity in IC~1182. This is however not confirmed by our spectroscopic analysis, that points out that all the optical properties can be explained in terms of a massive star formation process. Zezas et al. (2002) and Metz et al. (2003) have analyzed the X-ray flux distribution in the Antennae Galaxy and concluded that the X-ray emission is associated with young stellar clusters. IC~1182 could be, in that sense, similar to the Antennae Galaxy.  The evidence of the data lead us to consider IC~1182 as an ongoing merger. The Main Tail, and the Second Tail we report here for the first time, together with the the S-shaped material observed far from the body of the system and the presence of knots in the central region are clear signs of gravitational interaction. The data we present here indicate that that merging process is now  approaching the final steps. Morphologically the galaxy is close to a spheroidal system with peculiarities, but the two galaxies are still dynamically distinct. The observed rotation  amplitudes indicate that the involved galaxies are rather different, an early and a late type spiral. The difference in luminosity between the two tails would be an indication of the mass difference between the two galaxies. The organization of the ionized gas in a disk is in agreement with the predictions by the models developed by Barnes (2002). The low metallicity deduced for the gas in the two central knots indicates that important amounts of fresh  gas have been supplied by the late galaxy in the process of interaction and merging. The SFR deduced from the H$\\alpha$ luminosity is very high, amounting to 90~M$_{\\odot}$ per year, after correction for the high internal extinction, which places IC~1182 among the very bright starburst galaxies."
    },
    "0402/astro-ph0402443_arXiv.txt": {
        "abstract": "We present a new parallel PM N-body code named PMFAST that is freely available to the public.  PMFAST is based on a two-level mesh gravity solver where the gravitational forces are separated into long and short range components.  The decomposition scheme minimizes communication costs and allows tolerance for slow networks.  The code approaches optimality in several dimensions.  The force computations are local and exploit highly optimized vendor FFT libraries.  It features minimal memory overhead, with the particle positions and velocities being the main cost.   The code features support for distributed and shared memory parallelization through the use of MPI and OpenMP respectively.  The current release version uses two grid levels on a slab decomposition, with periodic boundary conditions for cosmological applications.  Open boundary conditions could be added with little computational overhead.  We present timing information and results from a recent cosmological production run of the code using a $3712^3$ mesh with $6.4\\times10^9$ particles.  PMFAST is cost-effective, memory-efficient, and is publicly available. ",
        "introduction": "N-body simulations are a key tool in astrophysics.  Applications range from cosmological problems involving dark matter to stellar systems and dynamics of galaxies.  In many astrophysical problems, $N$ can be very large.  For precision calibration of statistical weak lensing, large dynamic range is required and this provides a challenge to existing computational resources.  A recent development has been the move towards large massively parallel computers with cheap commodity components and relatively slow interconnects.  The burden of coding in the presence of a large memory hierarchy (commonly several layers of cache, local memory, remote memory, and secondary storage), and distributed message passing libraries is now placed on the scientist who wishes to utilize the large machines.  Our goal is to provide the community with a generic N-body code which runs close to optimally on inexpensive clusters. In this paper we describe the design and implementation of the algorithm, and performance numbers for cosmological applications.  Most real world applications on modern microprocessors achieve a small fraction of theoretical peak speed, often only a few percent. An order of magnitude in speedup is available through the use of assembly coded libraries.  These include routines such as FFT's that have been optimized to take advantage of the particular benefits that a given hardware manufacturer can offer in terms of instruction set and processor architecture developments.  A second limiting factor is the amount of physical memory.  Most N-body codes are not very efficient in memory use.  In principle, one only requires 6 numbers per particle to store the positions and velocities.  In practice, other data structures such as density fields and force fields dominate memory usage.  In this paper we present an algorithm that approaches minimal memory overhead, using only seven numbers per particle, plus temporary storage which is small.  The computation is off-loaded onto highly optimized FFT's, and the communication cost on parallel machines is mitigated by a two level mesh hierarchy. ",
        "conclusions": "We have presented a new freely available parallel particle-mesh N-body code that takes a significant step towards achieving optimality in CPU, communication and memory performance.  The only $O(N)$ memory required is six floating point and one integer per particle.  A two level force decomposition allows for the use of a short range force which minimizes communication.  It also eliminates the need to store a global fine grid density field.  CPU performance is optimized by the use of vendor optimized FFT libraries, which allows one to deploy very fine grids. The code is available for download at:  http://www.cita.utoronto.ca/webpages/code/pmfast/"
    },
    "0402/astro-ph0402169_arXiv.txt": {
        "abstract": "We investigate some observational constraints on decaying vacuum cosmologies based on the recently discovered old high redshift quasar APM 08279+5255. This object is located at $z = 3.91$ and has an estimated age of 2-3 Gyr. The class of $\\Lambda(t)$ cosmologies is characterized by a positive $\\beta$ parameter smaller than unity which quantifies the ratio between the vacuum and the total energy density. Assuming the lower limit age (2 Gyr) and that the cold dark matter contributes with $\\Omega_{\\rm M}=0.2$ we show that $\\beta$ is constrained to be $\\ge 0.07$ while for an age of 3 Gyr and $\\Omega_{\\rm M}=0.4$ the $\\beta$ parameter must be greater than $0.32$. Our analysis includes closed, flat and hyperbolic scenarios, and it strongly suggests that there is no age crisis for this kind of $\\Lambda(t)$ cosmologies. Lower limits to the redshift quasar formation are also briefly discussed to the flat case. For $\\Omega_{\\rm M}=0.4$ we found that the redshift formation is constrained by $z_{f}\\ge 8.0$. ",
        "introduction": "Recent observations from Supernovae (SNe) type Ia  strongly suggest that the bulk of energy in the Universe is repulsive and appears like a dark component; an unknown form of energy with negative pressure [in addition to the ordinary dark matter] which is probably of primordial origin\\cite{PR}. The most natural candidate for dark energy is the cosmological constant ($\\Lambda$), or equivalently, a perfect fluid obeying the equation of state, $p_v = - \\rho_v$, which is usually interpreted as the constant vacuum energy density of all fields existing in the Universe. The $\\Lambda$-term is the simplest but not the unique possibility. Other candidates appearing in the literature are: a relic scalar field component (SFC) which is slowly rolling down its potencial\\cite{ratra}, a decaying vacuum energy density, or a time varying $\\Lambda$-term\\cite{OT}, the so-called ``X-matter\", an extra component\\cite{turner} characterized by an equation of state $p_{\\rm x}=\\omega\\rho_{\\rm x}$ (XCDM), and the Chaplygin type gas whose equation of state is $p= -A/\\rho^{\\alpha}$, where $A$ and $\\alpha$ are positive constants\\cite{bert} (see also Lima\\cite{LIMAR} for a quick review). On the other hand, the existence of old high-redshift objects is one of the best methods for constraining the age of the Universe, as well the basic cosmological parameters\\cite{Krauss97}. Such objects also provide an important key for determining the first epoch of galaxy formation. In this connection, quasars are among the most luminous objects known in the universe and their prominent emission-lines contains valuable information to estimate their ages. The recently reported age estimates of the APM 08279+5255 quasar with a lower limit of 2-Gyr-old at redshift $z=3.91$ is therefore a particularly interesting event\\cite{KOM}. In this article we investigate some cosmological implications from the existence of this quasar to a large class of decaying vacuum cosmologies proposed by Lima and collaborators\\cite{LM}. Some constraints on the first epoch of quasar formation are also discussed. ",
        "conclusions": ""
    },
    "0402/astro-ph0402396_arXiv.txt": {
        "abstract": "We use the local curvature to investigate the possible existence of non-Gaussianity/asymmetry in the WMAP data. Considering the full sky we find results which are consistent with the Gaussian assumption. However, strong non-Gaussian features emerge when considering the northern and southern galactic hemisphere separately, particularly on scales between 1 and 5 degrees. Quite interestingly, the maximum non-Gaussianity is found for hemispheres centered near the ecliptic poles, which might suggest the presence of some systematic effect. The direction of the asymmetry seems consistent with the findings by Eriksen et al. 2004. ",
        "introduction": "\\label{sect:intro}  Many recent papers \\citep{komatsu,naselsky,naselsky2,chiang,park,eriksen,naselsky3,coles,copi,vielva,magueijo,eriksen2} have tested the Gaussianity of the cosmic microwave background (CMB) on the WMAP data. The most striking result from some of these papers is the possible existence of an asymmetry in the distribution of the CMB fluctuations on different hemispheres. For instance \\citep{eriksen2} find strong non-Gaussianity, as measured by the genus, on the northern galactic hemisphere; \\citep{vielva} detect, by using a wavelet analysis, a distinct non-Gaussian feature in the southern hemisphere; similarly results by \\citep{copi,eriksen,park} seem inconsistent with the assumption that CMB fluctuations are isotropic/Gaussian on various scales. The physical implications of these findings are still unclear. In this letter we aim at a further investigation on this issue by using a different method; namely the local curvature approach advocated by \\citep{dore}, which we extend from the flat sky approximation to the full sky setting \\citep{paolo}.\\\\ The plan of the paper is as follows; in section (\\ref{sect:curv}) we briefly describe the local curvature formalism on the sphere, section (\\ref{sect:method}) outlines our method and section (\\ref{sect:results}) collects the results, which we discuss in section (\\ref{sect:concl}). ",
        "conclusions": "\\label{sect:concl}  In this letter we have further investigated the possible existence of non-Gaussianity/asymmetry in the WMAP data. Considering the full sky data (using an extended Kp2 galaxy cut and point source mask) we find results which are consistent with the Gaussian assumption. However, strong non-Gaussian features (2-3$\\sigma$) emerge when considering the northern and southern galactic hemisphere separately, particularly on scales between 1 and 5 degrees where the proportions of hills and lakes are extermely high on one hemisphere and extremely low on the other. Quite interestingly, the maximum asymmetry is found for hemispheres centered close to the ecliptic poles, which could suggest the presence of a systematic effect. For the 5 to 10 degree FWHM smoothing scale, the southern hemisphere yields results so close to the Gaussian expected values to suspect overestimation of the variance in this part of the sky. The position of the non-Gaussianity on the sky seems very consistent with the asymmetric distribution of the power spectrum and the three-point correlation functions found by \\citep{eriksen}."
    },
    "0402/astro-ph0402675_arXiv.txt": {
        "abstract": "We report on the merger-induced generation of a shock-heated gas wind and formation of a remnant gas halo in simulations of colliding disk galaxies.  The simulations use cosmologically motivated initial conditions and include the effects of radiative cooling, star formation, stellar feedback and the non-adiabatic heating of gas.  The non-adiabatic heating, i.e. shocks, generated in the final merger forces gas out of the central region of the merger remnant and into the dark-matter halo.  We demonstrate that the amount of heating depends on the size of the progenitor disk galaxy as well as the initial orbit the galaxies are placed on.  Based upon these dependencies, we motivate a possible recipe for including this effect in semi-analytic models of galaxy formation. ",
        "introduction": "\\label{sec:intro}  Galactic winds are a common phenomenon associated with starbursting galaxies in the local \\citep{M99,Heck00,St02} and high redshift \\citep{Pet02,Adel03} universe.  It is generally assumed that feedback from massive stars, either stellar winds or supernova ejecta or both, power these outflows of hot gas.  The starburst origin of galactic winds in nearby galaxies is supported by the observed correlation between wind-induced mass loss and the star formation rate \\citep*{M99,Ru02}. In addition to transporting mass and energy from the central star-forming region, galactic winds also transport metals \\citep{Heck00}.  Depending on the energy in the winds and the depth of the halo potential well, outflows such as these may be responsible for heating and polluting the intergalactic medium.  Galactic winds likely play a significant role in galaxy formation and evolution.  As a source of energy, or {\\it feedback}, galactic winds regulate the conversion of gas into stars.  In recent semi-analytic models of galaxy formation, winds (in addition to a UV background) are required to remove gas and suppress star formation in small mass halos, making the theoretical halo mass function consistent with the observed galaxy luminosity function \\citep{S02,BenIII}.  Energy input due to galactic winds counteracts the efficient gas cooling in galactic halos - the so-called ``overcooling'' problem of theoretical models or simulations. Interestingly, not only do these models produce too many faint galaxies, but they typically produce too many bright galaxies as well.  Thus, overcooling is a problem in both low {\\it and} high mass halos.  While feedback from massive stars provides a physically reasonable and effective mechanism by which the low mass end of the halo mass function can be reconciled with the galaxy luminosity function, it is unclear that supernova-driven winds can provide enough energy to solve the overcooling problem for large halos.   \\citet{Ben03}, using semi-analytic models, were able to match the bright end of the luminosity function only by assuming extremely efficient thermal conduction or by including galactic winds whose energy greatly exceeded that produced by supernovae.  It is clear additional sources of energy are required during the process of galaxy formation and in this letter we suggest a possible candidate.  We demonstrate that simulations of disk-galaxy mergers generate shock-heated winds which can put gas with thermal energy up to 1 keV per baryon into the dark-matter halo.  These winds are {\\it not} induced by supernovae, but instead caused by orbital energy which is transferred to the diffuse gas.  Since major mergers are more common during the hierarchical formation of massive structures, energy input due to this process will preferentially act in more massive halos.  Further, the generation of hot gas due to the merging process suggests that we should be able to detect hot gas in nearby merging systems.  The structure of this letter is as follows: \\S \\ref{sec:sim} summarizes our simulation techniques and our initial conditions, \\S \\ref{sec:mergers} presents an overview of the galaxy merger process focusing on the conditions relevant to generating hot galaxy winds, \\S \\ref{sec:outfm} gives properties of the outflow material and presents a physically motivated fitting formula to capture our results which we discuss in \\S \\ref{sec:disc}. ",
        "conclusions": "\\label{sec:disc}  Hydrodynamic simulations of colliding disk galaxies demonstrate the ability to generate a significant amount of hot gas from shocks which abundantly occur in the central several kpc region during galaxy collisions.  As the centers of disk galaxies coalesce, collisional gas cannot inter-penetrate and gets heated while attempting to follow the potential well dominated by the collisionless dark matter and stellar components.  We note that any isothermal gas assumption used to simplify the treatment of gas \\citep{BH:1996,MH:1996} would not reproduce the process we discuss here.  In general, hot gas generated by galaxy major mergers could be an important source of energy during the process of galaxy formation.  Not only do radial orbits pump energy and entropy into the gas but they provide an effective mechanism to recycle moderately dense cold gas into diffuse hot gas, providing a feedback loop for energy, metals and mass that correlates with the hierarchical build-up of any galaxy. In reality the process we describe here may be augmented by contributions from stellar winds, supernovae, and AGN which will increase the total gas heating.  The remnants we report here have cooling times of several Gyr and could provide an enriched medium with which a gas disk could subsequently be formed.  Our results suggest that any galaxy system which has been involved in a major merger should have hot gas residing throughout its dark matter halo.  Additionally, the temperature of this gas, at a fixed time after the final merger, should correlate with the mass of the progenitor disk galaxies and their merger orbit.  Further work will investigate if the profiles in our merger remnants resemble the extended gas distributions found in relaxed elliptical dominated groups as found by \\citet{mz98} or if the outflowing gas is similar to the diffuse gas found in ongoing mergers such as Arp220 \\citep{McD03} and `The Antennae' \\citep{Fab01}."
    },
    "0402/astro-ph0402439_arXiv.txt": {
        "abstract": "The \\GEMS\\ project involves a multi-wavelength study of a sample of 60 galaxy groups, chosen to span a wide range of group properties. Substantial \\ROSAT\\ \\PSPC\\ observations, available for all of these groups, are used to characterise the state of the intergalactic medium in each.  We present the results of a uniform analysis of these \\ROSAT\\ data, and a statistical investigation of the relationship between X-ray and optical properties across the sample.  Our analysis improves in several respects on previous work: (a) we distinguish between systems in which the hot gas is a group-scale medium, and those in which it appears to be just a hot halo associated  a central galaxy, (b) we extrapolate X-ray luminosities to a fixed overdensity radius (\\rfh) using fitted surface brightness models, in order to avoid biases arising from the fact that cooler systems are detectable to smaller radii, and (c) optical properties have been rederived in a uniform manner from the \\NASA\\ Extragalactic Database, rather than relying on the data in the disparate collection of group catalogues from which our systems are drawn.  The steepening of the \\LX-\\TX\\ relation in the group regime reported previously is not seen in our sample, which fits well onto the cluster trend, albeit with large non-statistical scatter.  A number of biases affect the fitting of regression lines under these circumstances, and until the impact of these has been thoroughly investigated it seems best to regard the slope of the group \\LX-\\TX\\ relation as being poorly determined. A significant problem in comparing the properties of groups and clusters is the derivation of system radii, to allow different systems to be compared within regions having the same overdensity. We find evidence that group velocity dispersion (\\sigmav) provides a very unreliable measure of system mass (and hence radius), with a number of groups having remarkably low values of \\sigmav, given that they appear from their X-ray properties to be collapsed systems. We confirm that the surface brightness profiles of groups are significantly flatter than those of clusters -- the {\\it maximum} value of the \\betafit\\ parameter for our sample is 0.58, lower than the typical value of 0.67 seen in clusters -- however, we find no significant tendency within our sample for cooler groups to show flatter profiles. This result is inconsistent with simple universal preheating models. The morphology of the galaxies in the \\GEMS\\ groups is correlated to their X-ray properties in a number of ways: we confirm the very strong relationship between X-ray emission and a dominant early-type central galaxy which has been noted since the early X-ray studies of groups, and also find that spiral fraction is correlated with the temperature of the hot gas, and hence the depth of the gravitational potential. A class of spiral-rich groups with little or no X-ray emission, probably corresponds to groups which have not yet fully collapsed. ",
        "introduction": "\\label{sec_intro}  The principle that {\\it we are nowhere special}, which is fundamental to cosmology, also applies to galaxies.  The majority of galaxies are, like our own, located within bound systems, mostly containing just a handful of bright galaxies \\citep{tully87}. These are characterised as {\\it galaxy groups}, which are distinguished rather arbitrarily from richer and rarer {\\it galaxy clusters}. These systems are evolving, as they turn round from the Hubble expansion, virialise, and grow through mergers and accretion.  This dynamical evolution modifies the environment of their constituent galaxies, and can in turn have profound effects on the evolution of the galaxies themselves. On the other hand, energetic galaxy winds can have a substantial impact on the surrounding intergalactic medium (IGM) within groups and clusters \\citep*[e.g][]{ponman99}, so that there is a two-way interaction between the structure of galaxies and galaxy systems.  The picture which emerges is that galaxies and the systems in which most of them are located {\\it co-evolve}, and a full understanding of the evolution of either galaxies or galaxy clusters must take into account the two-way interactions which couple the development of both.  Galaxy groups have received far less attention from astronomers than either galaxies or galaxy clusters, and their properties are clearly very diverse, in terms of structure, dynamics and the types of galaxies they contain \\citep{hickson97,zabludoff98,mulchaey00}. Any meaningful study of the relationship between groups and galaxies needs to acknowledge this fact. We have therefore commenced a study of the properties of a substantial sample of 60 galaxy groups, and the galaxies they contain, with a view to clarifying the different stages of group evolution, and the ways in which this is related to galaxy properties. This \\GEMS\\ (Group Evolution Multi-wavelength Study) project involves optical photometry and spectroscopy to study the galaxies, radio observations to explore the HI content of galaxies and to look for cool intergalactic gas, and X-ray studies to probe the hot gas which dominates the baryonic content of at least some galaxy groups, and also provides a valuable indicator that a group is truly a dense system in 3-dimensions.  Given the value of X-ray data, all groups in our sample have been selected to have high quality \\ROSAT\\ observations available -- though we have {\\it not} selected only groups which are detected in the X-ray. The present paper describes the analysis of these \\ROSAT\\ \\PSPC\\ data, and the properties derived from them, and combines these with other properties of these systems and their galaxies drawn from the literature, and in particular from the \\NASA-\\IPAC\\ Extragalactic Database (\\NED).  There have been a number of previous studies of samples of galaxy groups based primarily on pointed \\ROSAT\\ observations (e.g. \\citealt*{pildis95}; \\citealt{mulchaey96,ponman96,mulchaey98,helsdon00a,helsdon00b,mulchaey03}).  The present work improves on these in a number of respects:  \\begin{itemize} \\item it is one of the largest samples for which the X-ray data have been analysed in a uniform manner, \\item it includes systems with low X-ray luminosity, and some which are entirely undetected in the X-ray, \\item systems showing intergalactic X-ray emission have been distinguished from those in which the X-ray emitting gas appears to constitute only hot halo associated with the central galaxy, \\item galaxy membership and internal velocity dispersion of the groups have been rederived in a consistent way, using NED data and a sigma-clipping approach, within a projected overdensity radius, \\item fitted models have been used to extrapolate X-ray luminosity to a fixed overdensity radius, to compensate for systematic trends with temperature in the radial extent to which X-ray data are available. \\end{itemize}  The only other studies which share some (but not all) of these features, are those of \\citet{helsdon00a,helsdon00b} and \\citet{mulchaey03}, with which we make a number of comparisons below.  Throughout this paper we use \\Hzero\\ = 70 \\kmpspMpc, and all errors correspond to 1$\\sigma$. ",
        "conclusions": ""
    },
    "0402/astro-ph0402113_arXiv.txt": {
        "abstract": "The Hubble classification scheme of galaxies is based on blue-light appearance. Atlases reveal the rich variety of responses of the Population I component (`the mask') of gas and dust to the underlying, older, stellar population. However, the Population I component may only constitute 5 percent of the dynamical mass of the galaxy; furthermore, dusty masks are highly effective in hiding bars. We firstly discuss the rich duality in spiral structure, and highlight a near-infrared classification scheme for spiral galaxies. We next show that images secured with SALTICAM will be ideally suited to probe key questions such as whether the optical light in the gaseous Population I component is the result of Kolmogorov turbulence, cascasding from the largest of scales down to the Nyquist limit. If so, the optical emission in galaxies will be  organized in a global fractal pattern with an intrinsic 1D power spectrum having a slope of -5/3, or -8/3 in 2D. ",
        "introduction": "In Roget's Thesaurus, we find the following:  {\\bf Mask:} [noun] screen, cloak, shroud. [verb] to camouflage, to make opaque, to disguise.  Optically thick dusty domains in galactic disks can completely camouflage or disguise underlying stellar structures. {\\it Cosmic dust grains act as masks}. The dust masks obscure whether or not the dust lies in an actual screen or is well intermixed with the stars. The presence of dust and the morphology of a galaxy are inextricably intertwined: indeed, the morphology of a galaxy can completely change once the Population I disks of galaxies -- the masks -- are dust penetrated (e.g., Block and Wainscoat 1991; Block et al., 1994, 2000).  \\begin{figure} \\vspace{12cm} \\caption{The bar in the Large Magellanic Cloud is beautifully portrayed in this naked eye drawing by Sir John Herschel in 1847. Two dimensional power spectra of HI emission in the LMC, spatially spanning three orders of magnitude, betrays the presence of Kolmogorov turbulence (Elmegreen, Kim and Staveley-Smith, 2001), with a power law slope of -8/3. SALTICAM will provide unique opportunities to examine Kolmogorov turbulence in galaxies beyond the Magellanic Clouds. Reproduced from de Vaucouleurs \\& Freeman (1972).} \\end{figure}  The classification of galaxies has traditionally been inferred from photographs or CCD imaging shortward of the 1$\\mu m$ window, where stellar Population II disks are not yet dust-penetrated. Images through an $I$ (0.8 $\\mu m$) filter can still suffer from attenuations by dust at the 50\\% level. The NICMOS and other near-infrared camera arrays offer unparalleled opportunities for deconvolving the Population I and II morphologies, because the opacity at $K$ -- be it due to absorption or scattering -- is always low. The extinction (absorption+scattering) optical depth at $K$ is only 10\\% of that in the V-band.  Many years before the advent of large format near-infrared camera arrays, it became increasingly obvious from rotation curve analyses that optical Hubble type is not correlated with the evolved Population II morphology. This was already evident in the pioneering work of Zwicky (1957) when he published his famous photographs showing the `smooth red arms' in M51. In the {\\it Hubble Atlas} and other atlases showing optical images of galaxies, we are looking at masks: at the gas, not the stars, to which the properties of rotation curves are inextricably tied. ",
        "conclusions": ""
    },
    "0402/astro-ph0402325_arXiv.txt": {
        "abstract": "We present the discovery of seven new T dwarfs identified in the Two Micron All Sky Survey. Low-resolution (R$\\sim$150) 0.8--2.5 $\\micron$ spectroscopy obtained with the IRTF SpeX instrument reveal the characteristic H$_2$O and CH$_4$ bands in the spectra of these brown dwarfs. Comparison to spectral standards observed with the same instrument enable us to derive classifications of T3 to T7 for the objects in this sample.  Moderate-resolution (R$\\sim$1200) near-infrared spectroscopy for a subset of these discoveries reveal \\ion{K}{1} line strengths consistent with previously observed trends with spectral type. Follow-up imaging observations provide proper motion measurements for these sources, ranging from $<$ 0$\\farcs$1 to 1$\\farcs$55 yr$^{-1}$. One object, 2MASS 0034+0523, has a spectrophotometric distance placing it within 10 pc of the Sun.  This source also exhibits a depressed K-band peak reminiscent of the peculiar T dwarf 2MASS 0937+2931, and may be a metal-poor or old, high-mass brown dwarf.  We also present low resolution SpeX data for a set of M and L-type dwarf, subdwarf, and giant comparison stars used to classify 59 additional candidates identified as background stars. These are primarily M5-M8.5 dwarfs, many exhibiting \\ion{H}{1} Paschen $\\gamma$, but include three candidate ultracool M subdwarfs and one possible early-type L subdwarf. ",
        "introduction": "T dwarfs are a spectral class of brown dwarfs distinguished by the presence of CH$_4$, H$_2$O, and H$_2$ collision-induced absorption (CIA) in the near-infrared \\citep{kir99,me02a,geb02}, and heavily pressure-broadened \\ion{K}{1} and \\ion{Na}{1} absorption at optical wavelengths \\citep{tsu99,bur00,me03d}.  These objects have effective temperatures (T$_{eff}$s) ranging from $\\sim$1300 K at the transition between L and T dwarfs \\citep{kir00,stp01,dah02,vrb04} to $\\sim$750 K for the latest-type T dwarf 2MASS 0415$-$0935 \\citep{me02a,vrb04}. T dwarfs therefore comprise the coldest and intrinsically faintest brown dwarfs currently known, and as such are key objects for testing brown dwarf and extrasolar giant planet atmosphere models \\citep{bar03}, probing the extreme low-mass end of the initial mass function \\citep{all04,me04a}, and expanding the census of the Sun's nearest neighbors.  For the past two years, we have been conducting a wide-field (74\\% of the sky) search for T dwarfs in the Two Micron All Sky Survey \\citep[hereafter 2MASS]{skr97,cut03}.  This three-band, near-infrared ($JHK_s$) imaging survey samples the peak of the T dwarf spectral energy distribution, and is therefore the most sensitive wide-field sky survey currently available for identifying these cold brown dwarfs.  Our results to date include the discovery of the bright, and therefore potentially very close ($d$ $\\approx$ 8 pc) T dwarf 2MASS 1503+2525 \\citep[hereafter Paper I]{me03a} and three new T dwarfs identified in the Southern Hemisphere \\citep[hereafter Paper II]{me03e}. Here we present the discovery of seven new T dwarfs in both Northern and Southern Hemispheres, all of which were verified by low resolution (R$\\sim$150) spectroscopic observations obtained with the IRTF 3.0m SpeX instument \\citep{ray03}.  In $\\S$ 2 we describe near-infrared imaging and spectroscopic observations of T dwarf candidates and comparison stars made using SpeX and other imaging instruments. In $\\S$ 3 we analyze these data, classifying both the new T dwarfs and background stars, including four potential ultracool (spectral types later than sdM7) subdwarfs, using the low-resolution spectra and spectral comparison stars. We also report line strengths for the 1.243/1.252 $\\micron$ \\ion{K}{1} doublet in six T dwarfs observed at moderate resolution (R$\\sim$1200) with SpeX, and proper motions for all of the T dwarf discoveries.  In $\\S$ 4 we discuss our results, including distance and tangential velocity estimates, signatures of gravity and/or metallicity in the near-infrared spectrum of 2MASS 0034+0523, and prospects for future discoveries. Results are summarized in $\\S$ 5. ",
        "conclusions": "\\subsection{Spectrophotometric Distances}  Using the derived classifications and 2MASS photometry, it is possible to estimate the spectrophotometric distances of our T dwarf discoveries. We employed the polynomial $J$- and $K_s$-band absolute magnitude/spectral type relations from \\citet{tin03}, based on parallax measurements from their own program and from \\citet{dah02}.  Distance estimates for each of the newly discovered T dwarfs were calculated in both bands (with the exception of 2MASS 0034+0523 which was not detected at $K_s$ by 2MASS) and for $\\pm$0.5 subclasses about the nominal classification.  Final distances and uncertainties were determined by the mean and standard deviation of these estimates, respectively, and are listed in Table 5.  Assuming that they are single sources, all of the objects listed are within 25 pc of the Sun, with the most distant object being the earliest-type T dwarf 2MASS 1209$-$1004. Three objects, 2MASS 0034+0523, 2MASS 1231+0847, and 2MASS 1828$-$4849, are at or within 10 pc from the Sun within the reported uncertainties, with the first (and latest-type) object having an estimated distance of only 8.2$\\pm$1.5 pc.  It should be noted that the uncertainties in these distance estimates do not take into account systematic deviations in the absolute magnitude/spectral type relation \\citep{tin03} nor possible duplicity, and should be confirmed by parallax measurement.  Combining the distance estimates with the proper motion determinations from $\\S$ 3.3, we calculated tangential velocities for these T dwarfs, listed in Table 8.  The mean $V_{tan}$ of those T dwarfs with detected motion is 43 km s$^{-1}$ with a standard deviation of 28 km s$^{-1}$; including the upper limits yields a somewhat lower mean of 33 km s$^{-1}$.  This value is consistent with the mean $V_{tan}$ for disk dwarfs (39 km s$^{-1}$; Reid \\& Hawley 2000) but is somewhat higher than that for field late-type M and L dwarfs (22 km s$^{-1}$; Gizis et al.\\ 2000), suggesting that the T dwarfs in this sample may be drawn from a somewhat older population. A similar difference in the $V_{tan}$ distribution between field L and T dwarfs has also been noted by \\citet{vrb04}, and a mean age difference between these classes is predicted in field substellar mass function simulations \\citep{me04a}. This possible age segregation is consistent with the evolution of brown dwarfs, as an object of a given mass will evolve from warm (L dwarf) to cold (T dwarf) as it ages; however, a larger sample must be considered before drawing any firm conclusions about the relative ages of the L and T dwarf field populations.  We note that the spectrophotometric distance of 2MASS 1231+0847 is consistent within its uncertainties to the nearby (13.4$\\pm$0.3 pc; HIPPARCHOS; Perryman et al.\\ 1997) K7V star Gliese 471, located roughly 8$\\farcm$1 (6500 AU) to the northwest. Indeed, 2MASS 1231+0847 was also identified in a parallel search for wide brown dwarf companions to nearby stars currently being conducted by J.\\ D.\\ Kirkpatrick.  However, while their direction of motion is nearly identical ($\\sim$230$\\degr$), 2MASS 1231+0847 has a proper motion nearly twice as large as Gliese 471, and is therefore not a bound companion.   \\subsection{Gravity/Metallicity Signatures at K-band}  The T7 2MASS 0034+0523 has the bluest $J-K_s$ color amongst the T dwarf discoveries, and it may be even bluer as 2MASS photometry provides only an upper limit. Examination of near-infrared spectral data confirms this color, as 2MASS 0034+0523 exhibits a fairly suppressed $K$-band peak in comparison to the rest of the T6--T8 dwarfs observed.  This is quantified in Figure 9, which plots the logarithm of the spectral ratio \\begin{equation} K/J = \\frac{\\int{F_{2.06-2.10}}}{\\int{F_{1.25-1.29}}} \\end{equation} as a function of spectral type for all of the T dwarfs observed.  SpeX prism data are ideally suited for this measurement, as the full near-infrared spectral range is sampled in a single order and no correction is required to match the relative scalings between the $J$- and $K$-bands. The ratio exhibits a tight linear trend across the full spectral type range of T dwarfs, although there is a somewhat greater spread in values amongst the T5--T6 dwarfs.  2MASS 0034+0523 stands well above this trend, however, having the smallest value of $K/J$, and hence the most suppressed $K$-band peak, in the entire sample.  The spectral properties of this source are reminiscent of the peculiar T6 2MASS 0937+2931 \\citep{me02a}, which not only has a very blue near-infrared color ($J-K_s = -0.62{\\pm}0.14$; Paper I), likely due to enhanced CIA H$_2$ absorption, but also strong absorption from the pressure-broadened 0.77 $\\micron$ \\ion{K}{1} resonance doublet and the 0.99 $\\micron$ FeH band \\citep{me03d}.  The enhanced pressure-sensitive features are symptomatic of a high pressure photosphere, which can exist on a brown dwarf with a high surface gravity (i.e., old and massive) and/or a metal deficient atmosphere \\citep{bur02}. Indeed, strong FeH and CIA H$_2$ absorption are hallmarks of cool halo subdwarf spectra (Figure 1), which are themselves typically older and metal-poor, and 2MASS 0937+2931 has been interpreted as a possible thick disk or halo brown dwarf \\citep{bur02,me03d}.  2MASS 0034+0523, like 2MASS 0937+2931, also has exceedingly weak 1.243/1.252 $\\micron$ \\ion{K}{1} lines that may be due to reduced metallicity or higher surface gravity \\citep{bur02}, although its late spectral type (and hence cool temperature) may be the dominant factor. 2MASS 0034+0523 does not have a large $V_{tan}$ as might be expected for a halo star, although on an individual basis this does not rule out its membership in an older kinematic population.  Clearly, a more detailed study of both of these peculiar T dwarfs is needed to assess metallicity and/or gravity effects in cool brown dwarf spectra.  \\subsection{Prospects for Further Discoveries}  To date, we have observed roughly 70\\% of our 2MASS search sample, identifying 31 T dwarfs, 8 of which have estimated or measured \\citep{dah02,tin03,vrb04} distances within 10 pc of the Sun.  This is roughly consistent with the predicted numbers from Paper I ($\\sim$35--45 T dwarfs), although we have uncovered less than half of the $\\sim$20 T dwarfs predicted to have distances less than 10 pc.  Many of the latter are probably late-type T dwarfs, T$_{eff}$ $\\lesssim$ 1000 K, too faint to be detected by 2MASS beyond a few parsecs. In addition to these very low-luminosity sources, our search has identified one {\\em bona-fide} ultracool halo subdwarf, the late-type sdL 2MASS 0532+8246 \\citep{me03f}, and now four candidate late-type subdwarfs requiring further verification. In retrospect, the search criteria employed are well-suited for identifying these metal-deficient objects, which have red optical/near-infrared colors, peak in flux at $J$-band, and exhibit relatively blue $J-K_s$ colors due to enhanced CIA H$_2$ absorption \\citep{leg00}. Such serendipitous discoveries provide a new opportunity for exploring the physical properties, particularly metallicity and age diagnostics, of very cool stars and brown dwarfs."
    },
    "0402/astro-ph0402055_arXiv.txt": {
        "abstract": "The most commonly used definition of halo formation is the time when a halo's most massive progenitor first contains at least half the final mass of its parent. Reasonably accurate formulae for the distribution of formation times of haloes of fixed mass have been available for some time. We use numerical simulations of hierarchical gravitational clustering to test the accuracy of formulae for the mass at formation. We also derive and test a formula for the joint distribution of formation masses and times. The structure of a halo is expected to be related to its accretion history.  Our tests show that our formulae for formation masses and times are reasonably accurate, so we expect that they will aid future analytic studies of halo structure. ",
        "introduction": "There is a simple analytic approximation for the distribution of halo formation times, when formation is defined as the time when the most massive progenitor first contains at least half the mass of the final object (Lacey \\& Cole 1993, 1994).  (Throughout, we will use the word parent to denote the final object, and the word progenitor to denote the smaller pieces which made up the mass of the parent at some earlier time.)  This formula provides a good description of what is seen in numerical simulations of gravitational clustering from Gaussian initial conditions, although recent work indicates that the agreement is not perfect (e.g., Wu 2001; Lin, Jing \\& Lin 2003). The sense of the discrepancy is that haloes in simulations appear to form slightly earlier than predicted, in qualitative agreement with previous work by Tormen (1998).  A related question is, what is the distribution of the mass of a halo at formation?  Absent other information, natural assumptions about this distribution are (i) that it is a delta function centered at one-half, or (ii) that the formation mass is uniformly distributed between one-half and unity. The second assumption is motivated by the fact that halo formation is expected to be a stochastic process; haloes of the same mass may have had different formation histories. The main purpose of the present paper is to derive and test a formula for the joint distribution of formation times and masses. Section~\\ref{massform} studies the distribution of formation masses whatever the formation time.  It shows that the distribution of masses just prior to, and just after formation, measured in simulations are both significantly different from delta functions, or from a uniform distribution, but are rather similar to simple formulae for these quantities derived by Nusser \\& Sheth (1999). Section~\\ref{joint} studies the conditional distribution of the formation mass, when the formation time is known.  This distribution is much better fit by a formula we derive here, than by a delta function or a uniform distribution. A final section summarizes our findings, and discusses possible applications. ",
        "conclusions": "We presented evidence that formulae for the distribution of formation masses (equations~\\ref{pmformgt} and~\\ref{pmformlt}), were reasonably accurate (Figure~\\ref{fmass}). These formulae do not depend on the shape of the underlying power spectrum, so they are simple to use. We then derived an expression for the conditional distribution of formation masses if the formation time is known (equation~\\ref{pmz}), and showed that it was also in quite good agreement with measurements made in simulations (Fig.~\\ref{mgivenz}).  Application of Bayes' rule then gives the joint distribution of formation mass and time.  Our results indicate that haloes which form at abnormally early times are more likely to have formation masses of order one-half that of the final mass of the parent, whereas haloes which form at abnormally late times are more likely to have formation masses which are closer to that of the parent.  One consequence of this is that haloes which form late are more likely to have experienced a recent major merger. We argued that this was a generic consequence of hierarchical formation.  Our formulae for the joint distribution of formation masses and times will find use in studies which attempt to relate the structure of a halo to its formation history (e.g. Tormen 1997, 1998; Tormen, Diaferio \\& Syer 1998; van den Bosch 2002; Wechsler et al. 2002; Zhao et al. 2003). For instance, haloes which formed recently with large formation masses are almost certainly further from equilibrium than haloes which formed at higher redshift with formation masses of order fifty-percent.  Such haloes (i.e. ones which have suffered major-mergers recently) may plausibly be less centrally concentrated than haloes of the same mass which had more quiescent accretion histories.  Addressing such issues is the subject of on-going work. If these formulae do prove to be useful, it will become necessary to modify them slightly so that they are more fully consistent with the parent halo mass function described by Sheth \\& Tormen (1999).  \\bigskip  We would like to thank the Aspen Center for Physics for support, and for providing the stimulating environment in which this work was completed. We would also like to thank the Virgo consortium for making the simulation data used here publically available at {\\tt http://www.mpa-garching.mpg.de/Virgo}. RKS was supported by the DOE and NASA grant NAG 5-7092 at Fermilab when work on this project began, and acknowledges support from NSF grant AST-0307747."
    },
    "0402/astro-ph0402263_arXiv.txt": {
        "abstract": "The energy, or mass scale $M_{\\rm SUSY}$, of the supersymmetry (SUSY) phase transition is, as yet, unknown. If it is very high (i.e., $\\gg 10^{3}\\, {\\rm GeV}$), terrestrial accelerators will not be able to measure it. We determine $M_{\\rm SUSY}$ here by combining theory with the cosmic microwave background (CMB) data. Starobinsky suggested an inflationary cosmological scenario in which inflation is driven by quantum corrections to the vacuum Einstein's equation. The modified Starobinsky model (MSM) is a natural extension of this. In the MSM, the quantum corrections are the quantum fluctuations of the supersymetric (SUSY) particles, whose particle content creates inflation and whose masses terminate it. Since the MSM is difficult to solve until the end of the inflation period, we assume here that an effective inflaton potential (EIP) that reproduces the time dependence of the cosmological scale factor of the MSM can be used to make predictions for the MSM. We predict the SUSY mass scale to be $M_{\\rm SUSY}\\simeq 10^{15}{\\rm GeV}$, thus satisfying the requirement that the predicted density fluctuations of the MSM be in agreement with the observed CMB data. ",
        "introduction": "Cosmological models based on the anomaly-induced effective action of gravity take into account the vacuum quantum effects of particles in the early universe. These models naturally lead to inflation, which is a direct consequence of the assumption that, at high energies, all particles can be described by massless, conformally invariant fields with negligible interaction between them. \\cite{fhh}\\cdash\\cite{wave}  In the modified version of the original Starobinsky model \\cite{star1} [modified Starobinsky model (MSM)], it was assumed that some of the scalar and fermion fields are massive. \\cite{anju}\\cdash\\cite{asta} If we also assume a supersymmetric particle content, inflation is, then, stable during the period which starts at a sub-Planck scale, continuing almost until the end of inflation, when most of the sparticles decouple. \\cite{asta}\\cdash\\cite{gorb} Subsequently, the universe enters into an unstable regime, eventually undergoing a transition to the FRW evolution. It is to be noted that this occurs without any need for fine-tuning, regardless of the values of the cosmological constant $\\Lambda$ and the curvature parameter $k$.  In fact both the cosmological constant $\\Lambda$ and the curvature $k$ are very small in the context of the high-T early universe. Present observations indicate that the curvature $k$-term is less than $10\\%$ of the energy density term in the Friedmann equation and that the $\\Lambda$ term is comparable to the energy density term in the $\\Lambda$CDM model. Extrapolating back to the high-T early universe, the $k$-term is then $\\lesssim 10^{-50}\\%$ and the $\\Lambda$ term $\\lesssim 10^{-100}\\%$ of the energy density term. The $k$-term is thus not included in our equations and the $\\Lambda$-term is included up to Eq.(\\ref{parabola}), just for completeness. We set $\\Lambda=0$ in the equations after Eq.(\\ref{parabola}).   It is not possible to solve analytically the equation for the cosmological scale factor $a(t)$ in the MSM. However an approximate solution for $a(t)$ during inflation was obtained and confirmed by numerical analysis with very high precision. \\cite{asta}  We assume here that an effective inflaton potential (EIP) that reproduces the time dependence of $a(t)$ of the MSM during the inflationary period can be used to make predictions of the MSM throughout this period, including its end. We use a reverse engineering method, discussed in Ref.~\\refcite{ellis1}, to derive the EIP from $a(t)$.  In $\\S$ 2, we give a brief review of anomaly-induced inflation of the MSM and the approximate time dependence of $a(t)$. The effective potential and value for $M_{\\rm SUSY}$ are derived in $\\S$ 3. Our conclusions are presented in $\\S$ 4. ",
        "conclusions": "A scalar potential was constructed using the approximate solution for the time dependence of the cosmological scale factor of the MSM during inflation. The potential was normalized at a time $\\sim 60$ $e$-folds before the end of inflation in order to obtain the observed level of density fluctuations in the CMB, $\\delta\\rho/\\rho\\sim 10^{-5}$. The mass (energy) scale of the MSM at the end of the inflation, $M_{\\ast}\\simeq 10^{12}{\\rm GeV}$, which we identify with the Hubble rate when the massive particles decouple, predicts a SUSY scale $M_{\\rm SUSY}$, consistent with the GUT scale $M_{\\rm SUSY}\\simeq 10^{15} {\\rm GeV}$."
    },
    "0402/astro-ph0402277_arXiv.txt": {
        "abstract": "% Extensive time-resolved observations of Kuiper Belt object 2001 QG$_{298}$ show a lightcurve with a peak-to-peak variation of $1.14 \\pm 0.04$ magnitudes and single-peaked period of $6.8872 \\pm 0.0002$ hr.  The mean absolute magnitude is 6.85 magnitudes which corresponds to a mean effective radius of 122 (77) km if an albedo of 0.04 (0.10) is assumed. This is the first known Kuiper Belt object and only the third minor planet with a radius $>$ 25 km to display a lightcurve with a range in excess of 1 magnitude.  We find the colors to be typical for a Kuiper Belt object ($B-V = 1.00\\pm 0.04$, $V-R = 0.60\\pm 0.02$) with no variation in color between minimum and maximum light. The large light variation, relatively long double-peaked period and absence of rotational color change argue against explanations due to albedo markings or elongation due to high angular momentum.  Instead, we suggest that 2001 QG$_{298}$ may be a very close or contact binary similar in structure to what has been independently proposed for the Trojan asteroid 624 Hektor.  If so, its rotational period would be twice the lightcurve period or $13.7744 \\pm 0.0004$ hr.  By correcting for the effects of projection, we estimate that the fraction of similar objects in the Kuiper Belt is at least $\\sim$10\\% to 20\\% with the true fraction probably much higher.  A high abundance of close and contact binaries is expected in some scenarios for the evolution of binary Kuiper Belt objects. ",
        "introduction": "The Kuiper Belt is a long-lived region of the Solar System just beyond Neptune where the planetisimals have not coalesced into a planet.  It contains about 80,000 objects with radii greater than 50 km (Trujillo, Jewitt \\& Luu 2001) which have been collisionally processed and gravitationally perturbed throughout the age of the Solar System.  The short-period comets and Centaurs are believed to originate from the Kuiper Belt (Fernandez 1980; Duncan, Quinn \\& Tremaine 1988).  Physically, the Kuiper Belt objects (KBOs) show a large diversity of colors from slightly blue to ultra red ($V-R \\sim 0.3$ to $V-R \\sim 0.8$, Luu and Jewitt 1996) and may show correlations between colors, inclination and/or perihelion distance (Jewitt \\& Luu 2001; Trujillo \\& Brown 2002; Doressoundiram et al. 2002; Tegler \\& Romanishin 2003). Spectra of KBOs are mostly featureless with a few showing hints of water ice (Brown, Cruikshank \\& Pendleton 1999; Jewitt \\& Luu 2001; Lazzarin et al. 2003).  The range of KBO geometric albedos is still poorly sampled but the larger ones likely have values between 0.04 to 0.10 (Jewitt, Aussel \\& Evans 2001; Altenhoff, Bertoldi \\& Menten 2004).  Time-resolved observations of KBOs show that $\\sim 32 \\%$ vary by $\\ge 0.15$ magnitudes, $18 \\%$ by $\\ge 0.40$ magnitudes and $12 \\%$ by $\\ge 0.60$ magnitudes (Sheppard \\& Jewitt 2002; Ortiz et al. 2003; Lacerda \\& Luu 2003; Sheppard \\& Jewitt 2004). One object, (20000) Varuna, displays a large photometric range and fast rotation which is best interpreted as a structurally weak object elongated by its own rotational angular momentum (Jewitt \\& Sheppard 2002).  A significant fraction of KBOs appear to be more elongated than main-belt asteroids of similar size (Sheppard \\& Jewitt 2002). The KBO phase functions are steep, with a median of $0.16$ magnitudes per degree between phase angles of 0 and 2 degrees (Sheppard \\& Jewitt 2002; Schaefer \\& Rabinowitz 2002; Sheppard \\& Jewitt 2004).  About $4\\% \\pm 2\\%$ of the KBOs are binaries with separations $\\ge 0.15 \\arcsec$ (Noll et al. 2002) while binaries with separations $\\ge 0.1 \\arcsec$ may constitute about $15 \\%$ of the population (Trujillo 2003, private communication).  All the binary KBOs found to date appear to have mass ratios near unity, though this may be an observational selection effect.  The mechanism responsible for creating KBO binaries is not clear.  Formation through collisions is unlikely (Stern 2002).  Weidenschilling (2002) has proposed formation of such binaries through complex three-body interactions which would only occur efficiently in a much higher population of large KBOs than can currently be accounted for.  Goldreich, Lithwick \\& Sari (2002) have proposed that KBO binaries could have formed when two bodies approach each other and energy is extracted either by dynamical friction from the surrounding sea of smaller KBOs or by a close third body.  This process also requires that the density of KBOs was $\\sim10^2$ to $10^3$ times greater than now.  They predict that closer binaries should be more abundant in the Kuiper Belt while Weidenschilling's mechanism predicts the opposite.  The present paper is the fourth in a series resulting from the Hawaii Kuiper Belt variability project (HKBVP, see Jewitt \\& Sheppard 2002; Sheppard \\& Jewitt 2002; Sheppard \\& Jewitt 2004).  The practical aim of the project is to determine the rotational characteristics (principally period and shape) of bright KBOs ($m_{R} \\le 22$) in order to learn about the distributions of rotation period and shape in these objects.  In the course of this survey we found that 2001 QG$_{298}$ had an extremely large light variation and a relatively long period.  We have obtained optical observations of 2001 QG$_{298}$ in order to accurately determine the rotational lightcurve and constrain its possible causes.  2001 QG$_{298}$ has a typical Plutino orbit in 3:2 mean-motion resonance with Neptune, semi-major axis at 39.2 AU, eccentricity of 0.19 and inclination of 6.5 degrees. ",
        "conclusions": "Kuiper Belt Object 2001 QG$_{298}$ has the most extreme lightcurve of any of the 34 objects so far observed in the Hawaii Kuiper Belt Variability Project.  1.  The double-peaked lightcurve period is $13.7744 \\pm 0.0004$ hr and peak-to-peak range is $1.14 \\pm 0.04$ mag.  Only two other minor planets with radii $\\ge$ 25 km (624 Hektor and 216 Kleopatra) and one planetary satellite (Iapetus) are known to show rotational photometric variation greater than 1 mag.  2.  The absolute red magnitude is $m_R$(1,1,0) = 6.28 at maximum light and 7.42 mag. at minimum light.  With an assumed geometric albedo of 0.04 (0.10) we derive effective circular radii at maximum and minimum light of 158 (100) and 94 (59) km, respectively.  3.  No variation in the BVR colors between maximum and minimum light was detected to within photometric uncertainties of a few percent.  4. The large photometric range, differences in the lightcurve minima, and long period of 2001 QG$_{298}$ are consistent with and strongly suggest that this object is a contact or nearly contact binary, viewed equatorially.  5.  If 2001 QG$_{298}$ is a contact binary with similarly sized components, then we conclude that such objects constitute at least 10\\% to 20\\% of the Kuiper Belt population at large sizes."
    },
    "0402/astro-ph0402088_arXiv.txt": {
        "abstract": "We study several effects that influence the strength of the proton cyclotron and atomic features in the thermal spectra of magnetic neutron stars. We show that it is possible for vacuum polarization to strongly suppress the spectral lines when the magnetic field $B\\ga 10^{14}$~G. For weaker fields ($B\\la 7\\times 10^{13}$~G), the surface spectrum is unaffected by vacuum polarization; thus the proton cyclotron absorption line can have a large equivalent width. Using an approximate calculation of the synthetic spectra, we show that variation of magnetic fields over the neutron star surface leads to broadening of the line. The recently detected absorption features from the isolated neutron stars RX~J$1308.6+2127$, RX~J$1605.3+3249$, and RX~J$0720.4-3125$ can plausibly be explained by the proton cyclotron resonance, with possible blending due to atomic transitions, in the atmosphere of the star. ",
        "introduction": "\\label{sec:intro}  In the last few years, considerable observational resources (e.g., {\\it Chandra} and {\\it XMM-Newton} telescopes) have been devoted to the study of thermal emission from isolated neutron stars (NSs) and, in particular, to the search for spectral features in the radiation. For many NSs, the spectra are found to be featureless and often well fit by a blackbody (e.g., Marshall \\& Schulz~2002; Burwitz \\etal~2003; see Pavlov, Zavlin, \\& Sanwal~2002 for a review).  Recently, absorption features have been found in the thermal emission of several isolated NSs. For example, the spectrum of the young NS 1E~$1207.4-5209$ (with surface temperature $T_{\\rm eff}\\simeq 2$~MK) shows features at 0.7 and 1.4~keV (Sanwal \\etal~2002; Mereghetti \\etal~2002; Hailey \\& Mori 2002) and possibly at 2.1 and 2.8 keV (Bignami \\etal~2003). Two NSs, RX~J$1308.6+2127$ (= RBS~1223) and RX~J$1605.3+3249$, belonging to the class of dim, radio-quiet isolated NSs (see Haberl~2003), have been observed to possess broad absorption features in their spectra: at $\\simeq 0.2-0.3$~keV for RX~J$1308.6+2127$ (Haberl \\etal~2003a) and at $\\simeq 0.45$~keV for RX~J$1605.3+3249$ (van Kerkwijk \\etal~2004). These are in contrast to two similar dim NSs, RX~J$1856.5-3754$ (Pons \\etal~2002; Burwitz \\etal~2003) and RX~J$0720.4-3125$ (Paerels \\etal~2001), which show featureless, blackbody-like spectra (however, see Haberl \\etal~2003b for the detection of a line at 0.27~keV in RX~J$0720.4-3125$). It is particularly striking that, although RX~J$0720.4-3125$, RX~J$1605.3+3249$, and RX~J$1308.6+2127$ have similar effective temperatures ($\\simeq 1$~MK), the equivalent widths (EW) of their lines are very different (van Kerkwijk~2003), from weakest in RX~J$0720.4-3125$ (EW $\\approx 40$~eV) to stronger in RX~J$1605.3+3249$ (EW $\\approx 80$~eV) to the strongest in RX~J$1308.6+2127$ (EW $\\approx 150$~eV). Motivated by these observations, we study in this paper the strength of the proton cyclotron line (and atomic lines) in the atmosphere emission of NSs as a function of magnetic field strength. We show that the absorption features observed in the three NSs (RX~J$1308.6+2127$, RX~J$1605.3+3249$, RX~J$0720.4-3125$) can possibly be explained by the proton cyclotron resonance at (unredshifted) energy \\be \\Ebp={\\hbar eB\\over m_pc}=0.63\\,B_{14}~{\\rm keV}, \\ee where $B=10^{14}B_{14}$~G is the magnetic field strength. (There may be blending of absorption features from radiative transitions in neutral H atoms). We also show that, owing to the vacuum polarization effect, the EW of the line can decrease dramatically as the magnetic field enters the magnetar regime ($B\\ga 10^{14}$~G).  In \\S\\ref{sec:vp}, we review and elaborate upon the effect of vacuum polarization on radiative transfer in NS atmospheres. In \\S\\ref{sec:spectra}, we present our numerical results for the atmosphere spectra with different magnetic field strengths; these results illustrate the suppression of atmosphere line features as the field strength increases. We consider synthetic spectra from NS atmospheres in \\S\\ref{sec:synth} and discuss our results in \\S\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  In several recent papers (Lai \\& Ho~2002, 2003; Ho \\& Lai~2003), we have shown that at superstrong magnetic fields, $B\\ga 10^{14}$~G, vacuum polarization can significantly affect the radiation spectrum from magnetized neutron star atmospheres: it softens the high-energy tail of the spectrum and suppresses the proton cyclotron feature (above $\\sim 0.6$~keV, unredshifted) and other features (see also Ho \\etal~2003).  The latter could provide an explanation for the non-detection thus far of lines in the observed thermal spectra of several anomalous X-ray pulsars (Patel \\etal~2001, 2003; Juett \\etal~2002; Tiengo \\etal~2002; Morii \\etal~2003) and soft gamma-ray repeaters (Kulkarni \\etal~2003), which are thought to possess $B\\ga 10^{14}$~G.  In this paper, we have studied the dependence of the neutron star atmosphere spectrum on the magnetic field strength. As we explain qualitatively and show numerically, at normal neutron star magnetic fields, $B\\la 10^{14}$~G, vacuum polarization has little effect on the atmosphere emission spectra. Therefore, strong proton cyclotron or other atomic features may be present in the thermal spectrum.  Our calculations of neutron star synthetic spectra, taking into account the line broadening effect due to magnetic field variation over the neutron star surface, show that the recently observed broad absorption features in several dim isolated neutron stars could be explained naturally as the proton cyclotron line (with possible blending from atomic lines of neutral hydrogen) from neutron star atmospheres with $B\\la 10^{14}$~G (see Fig.~\\ref{fig:synth}). The variation in the strength of the observed spectral features in these sources is then due to different fractions of the surface with $B\\la 10^{14}$~G.  For RX~J$0720.4-3125$, the weakness (phase-averaged EW $\\approx 40$~eV) of the line at 0.27~keV suggests that most of the observable surface of this neutron star has $B\\ga 10^{14}$~G, and the line is from the magnetic equatorial region of the star (where $B$ is weaker). The observed line energy does not change with phase but has a larger EW at pulse decline/minimum (Haberl \\etal~2003b), which also indicates the line is produced at a small region near the magnetic equator. Note, however, that Kaplan \\etal~(2002) and Zane \\etal~(2002) place an upper limit of $\\sim\\mbox{a few}\\times 10^{-13}\\mbox{s s$^{-1}$}$ on the period derivative of RX~J$0720.4-3125$, which, given its period of 8.39~s, implies a dipole field with $B_{\\rm dipole}<5\\times 10^{13}$~G; thus our inference of the surface field based on the thermal spectrum implies an appreciable non-dipolar magnetic field on the neutron star.  For RX~J$1308.6+2127$, the stronger (EW $\\approx 155$~eV) line at 0.3~keV suggests that most of the observable surface of this neutron star has $B\\la 10^{14}$~G. This neutron star also has a double peaked pulse profile (Haberl \\etal~2003a), as compared to the very sinusoidal single peaked pulse profile of RX~J$0720.4-3125$ (Haberl \\etal~2003b). The difference in pulse profiles may be due to differences in the viewing geometry and angle between the magnetic and rotation axes: in the case of RX~J$1308.6+2127$, the observer sees the equatorial region most of the time (thus the stronger line and slightly larger pulse fraction) and each magnetic pole, while only one pole is visible in the case of RX~J$0720.4-3125$.  Our calculations show that, for a given line-emitting area, the 0.45~keV line of RX~J$1605.3+3249$ should have a smaller EW than the 0.27~keV line of RX~J$0720.4-3125$. However, the line is stronger (EW $\\approx 80$~eV) in RX~J$1605.3+3249$. One possible way to reconcile the observed smaller EW of the 0.27~keV line in RX~J$0720.4-3125$\\footnote{de Vries \\etal~(2004) find long-term changes in the spectra of RX~J$0720.4-3125$; it is not clear whether this variability can affect the EW of the line.} is for the line-emitting area (with $B\\la 10^{14}$~G) to be a small fraction of the observed surface, while the larger fraction of the surface has a much higher magnetic field ($B\\ga 10^{14}$~G). Alternatively, we have only considered hydrogen, and some of the lines may be due to other elements, such as helium.  As discussed in \\S\\ref{sec:nc}, the numerical treatment of the vacuum polarization effect in the neutron star atmosphere models is approximate, i.e., it does not properly account for the partial mode conversion associated with the vacuum resonance. Future work in this direction is necessary (see \\laihotwo), as well as a more accurate and comprehensive study of partially ionized atmosphere models (see Ho \\etal~2003). Furthermore, we have only examined thermal emission; there have been cyclotron features seen in the non-thermal emission from magnetars (e.g., the 5.0~keV feature from SGR~$1806-20$ during outburst, Ibrahim \\etal~2002, 2003; the 8.1~keV feature from AXP~1RXS~J$170849-400910$, Rea \\etal~2003). Such theoretical/numerical studies, when combined with observational data, should provide useful constraints on the nature of various types of neutron stars, including the dim isolated neutron stars."
    },
    "0402/astro-ph0402041_arXiv.txt": {
        "abstract": "We present a stroboscopic system developed for optical observations of pulsars and its application in the CLYPOS survey. The stroboscopic device is connected to a GPS clock and provides absolute timing to the stroboscopic shutter relative to the pulsar's radio ephemerides. By changing the phase we can examine the pulsar's light curve. The precisely timed stroboscope in front of the CCD camera can perform highly accurate time resolved pulsar photometry and offers the advantages of CCD cameras, which are high quantum efficiency as well as relatively large field of view, which is important for flux calibrations.  CLYPOS (Cananea Ljubljana Young Pulsar Optical Survey) is an extensive search for optical counterparts of about 30 northern hemisphere radio pulsars. It is a collaboration between the INAOE, Mexico and the Faculty of Mathematics and Physics of the University of Ljubljana. Stroboscopic observations were done between December 1998 and November 2000 at the 2.12 m telescope of the Observatory Guillermo Haro in Cananea, Sonora. The first results of the survey are presented. Analyzed data indicate that there is no optical counterpart brighter than $V\\sim 22$. ",
        "introduction": "Today more than 1000 radio pulsars are known. The interest for detection of their optical counterparts has increased in the last decade, mainly due to bigger and better telescopes and due to improved detectors with high temporal resolution (Fordham et al. 2000, Dhillon et al. 2001, Moon et al. 2001, Straubmeier et al. 2001). Despite all efforts the number of detected optical counterparts has enlarged only to nine (Mignani et al. 2000, Mignani et al. 2002). Almost all of them are weak optical sources, with a $V$ magnitude of $\\geq 24$. Moreover, until now there has been no systematic search for optical counterparts of young radio pulsars. The Clypos survey, a mexican-slovene collaboration project, was an attempt to systematically search for optical counterparts of $\\sim 30$ radio pulsars. The goal of this survey was to detect new optical pulsars or at least to determine magnitude limits for their detection and thus set stronger observational constraints on theoretical high-energy emission models for pulsars. ",
        "conclusions": "With the analysis described above no optical counterpart with an average magnitude of 22-24 on the sample of 12 radio pulsars was detected. Optical counterparts already discovered are very faint. The limiting magnitude that we reached with the Clypos survey is most probably not high enough for any new detection. However, the stroboscopic method is a useful observational technique since it enables not only detection but also periodicity identification of optical counterparts. Once the absolute phase of the pulsar is known one can synchronize the stroboscope and thus fully use the reduction of the sky brightness contribution resulting from magnitude enhancement. For example, with the VLT using active optics and in stroboscopic mode we would expect that 0.01 photometry of a 25 average magnitude pulsar would be allowed in 10 minutes."
    },
    "0402/astro-ph0402331_arXiv.txt": {
        "abstract": "{ We present a new catalog of absorption-line systems identified in the quasar spectra. It contains data on 821 QSOs and 8558 absorption systems comprising 16~139 absorption lines with measured redshifts in the QSO spectra. The catalog includes absorption-line systems consisting of lines of heavy elements, lines of neutral hydrogen, Lyman limit systems, damped Ly$\\alpha$ absorption systems, and broad absorption-line systems. The catalog is available in electronic form at the CDS via anonymous ftp to {\\it cdsarc.u-strasbg.fr} (130.79.128.5) or via http://{\\it cdsweb.u-strasbg.fr/cgi-bin/qcat?J/A+A/412/707} and at {\\it www.ioffe.ru/astro/QC}. Using the data of the present catalog we also discuss redshift distributions of absorption-line systems. ",
        "introduction": "\\label{sect-intro} Absorption lines and absorption-line systems (ALSs) observed in the spectra of QSOs contain fundamental information on the distribution of matter between the observer and the QSO, and on physical processes in the Universe in different epochs of the cosmological evolution. To date, thousands of ALSs have been identified and their number grows persistently, scattered over numerous sources. This stimulates the creation of catalogs of ALSs comprising the most complete data on the absorption lines and their systems.  Catalogs of ALSs have been compiled many times. We mention the early catalogs of Perry et al.\\ (\\cite{pbb78}) and Ellis \\& Phillips\\ (\\cite{ep78}), and the later vast QSO catalogs of Hewitt \\& Burbidge (\\cite{hb80},  \\cite{hb87}, \\cite{hb89}, and \\cite{hb93}) which include also data on the ALSs detected in the QSO spectra. Junkkarinen et al.\\ (\\cite{jhb91}) and York et al.\\ (\\cite{yyccgm91}) created special ALS catalogs most complete for that time. The new generation of telescopes (Keck, VLT, etc.) has yielded a great amount of new spectroscopic data. Some have been collected in special catalogs including either the results of certain spectral investigations or the definite types of ALSs (e.g., Lyman limit systems -- LLS, damped Ly$\\alpha$ absorption systems  -- DLA, broad absorption-line systems -- BAL, etc.). For example, the catalog by Outram et al.\\ (\\cite{ossbclm01}) of the ALSs detected in the 2dF QSO Redshift Survey or the catalogue of DLAs compiled by Curran et al.\\ (\\cite{cwmbcf02}). However, as far as we know, there are no modern catalogs comprising comprehensive data on the ALSs registered to date.  Our new catalog is an attempt to collect the basic information on the ALSs in QSO spectra. The data are taken from publications available up to January 2002. The catalog includes, in particular, all the data of the catalogs of Junkkarinen et al.\\ (\\cite{jhb91}) and York et al.\\ (\\cite{yyccgm91}). The catalog consists of introduction (ReadMe), Tables 1 and 2, and list of references, which are available in electronic form at the CDS and at {\\it www.ioffe.ru/astro/QC}. ",
        "conclusions": ""
    },
    "0402/astro-ph0402107_arXiv.txt": {
        "abstract": "{ Autocorrelation functions (ACFs) are studied for a sample of 16 long gamma-ray bursts (GRBs) with known redshift~$z$, that were observed by the BATSE and Konus experiments. When corrected for cosmic time dilation, the ACF shows a bimodal distribution. A {\\it narrow} width class (11 bursts) has at half-maximum a mean width $\\tau'_o=1.6$~s with a relative dispersion of $\\sim 32$\\%, while a {\\it broad} width class (5 bursts) has $\\tau'_o=7.5$~s with a $\\sim 4$\\% dispersion. The separation between the two mean values is highly significant ($>7\\sigma$). This temporal property could be used on the large existing database of GRBs with unknown redshift. The broad width set shows a very good linear correlation between width at half-maximum and $(1+z)$, with a correlation coefficient $R=0.995$ and a probability of chance alignment $<0.0004$.  The potential application of this correlation to cosmology studies is discussed, using it in combination with recently proposed luminosity indicators. ",
        "introduction": "\\label{intro}  The knowledge of time scales and source distances are essential for the physical understanding of astronomical phenomena. From the first detections in 1969 by {\\it Vela} satellites \\citep{KSO73}, until the launch of {\\it BeppoSAX} in 1997, the distance scale of gamma-ray bursts (GRBs) remained unsettled. This mission provided arc-minute localization, leading to the discovery of a fading emission towards lower energy bands, the so-called afterglows. Thereafter, burst redshifts $z$ have been determined from spectroscopic analysis of the afterglows or, in some cases, of their associated host galaxies, proving that at least long-duration bursts are at cosmological distances. So far, the redshift of no short-duration burst  has been clearly determined \\citep[although see][]{kul02}. In this paper only the class of long GRBs will be considered (i.e., those with time duration $> 2$ s).   Up to date, more than 30 burst redshifts have been spectroscopically measured thanks to immediate follow-up observations. On the other hand, there is a wealth of data from thousands of GRBs for which the redshift is unknown. Most of these were detected by the Burst and Transient Source Experiment (BATSE). Other important motivations to find a redshift estimator based only on the gamma-ray prompt emission are the lack of optical counterparts in some cases (the so-called {\\it dark} afterglows), and the difficulty of spectroscopically determine redshifts beyond $z=5$ due to the Lyman alpha absorption. In recent years, two empirical relations have been discovered to estimate the luminosity distance exclusively from the analysis of the gamma emission. One relates the isotropic luminosity with the time lag between different energy channels \\citep{NMB00}, and the other with a variability parameter of the light curve \\citep{rei01}. Both luminosity correlations can be used to derive luminosity distances and, assuming some specific cosmology, the corresponding redshifts. Thus, from these correlations it has been possible to estimate GRB luminosity functions and demographic distributions \\citep[see, e.g.,][]{Nor02,LFR02}. These first estimations indicate that the GRB population may peak at redshift $z \\sim 10$, being then ideal probes of the early universe. However, the luminosity functions derived in these works predict source counts $N(>P)$, as a function of photon flux $P$, that differ significantly from the observed one \\citep{schmidt03}. Much better calibration of these empirical relations is needed, and that will only be possible with a much larger number of independent redshift determinations covering a broader $z$ range.   Individual power density spectra (PDS) of GRB are in general very diverse, but the longest bursts show power-law spectra extended over two frequency decades. Shorter bursts also display this property by averaging the PDSs of a large sample \\citep{BSS98,BSS00}. This underlying power-law behavior indicates the absence of any preferred time scale.  The autocorrelation function (ACF) is the Fourier transform of the PDS, therefore it contains in principle the same information that can be visualized in a different way. The ACF gives a measure of the correlation between different points in the light curve that are separated by a given time lag.  Various efforts have been made using these data analysis tools to find a temporal characteristic that might correlate with the redshift and, e.g., \\citet{CYC02} have found a weak correlation between the power-law index of the PDS and~$z$. See also \\citet{att03} for a proposed redshift indicator.  In this paper it will be shown that the ACF can be used to define characteristic times that strongly correlate with the redshift. In \\S~\\ref{methods} the data selection and the use of the ACF are described. Next in \\S~\\ref{results} it is shown that the ACF corrected for time dilation effects has a bimodal distribution, and that this property could be used to construct an empirical relation to estimate $z$. Finally, the results and their possible applications are discussed in \\S~\\ref{discussion}.       \\begin{figure*} \\centering \\includegraphics[width=17cm]{fig1.ps} \\caption{Comparison of the ACFs of 6 GRBs obtained using data from two different experiments. {\\it Solid lines:} Konus 64 ms data in the 50--200 keV energy band; {\\it gray lines:} BATSE 64 ms data in the 55--320 keV energy range. There is sufficiently good agreement for the bright bursts, when the results are not very sensitive to the background estimation. GRB 971214 is considerably dimmer than the others (see the text for discussion). \\label{combo}} \\end{figure*} ",
        "conclusions": "\\label{discussion}   The average PDS of bursts shows an overall power-law behavior, indicative of a self-similar underlying process where there are no preferred timescales. If this is the case, then the width of the ACF is related with the low-frequency cut-off of the PDS, which is due to the finite duration of the burst \\citep{BSS00}. As it was mentioned in \\S~\\ref{intro}, in principle the information given by the ACF and the PDS is the same. In practice, since they express this information differently, they are affected by noise and window effects in different ways. It would be difficult to make a good estimation of the low-frequency cut-off in the PDS due to the large statistical fluctuations. On the other hand, the width at half-maximum of the ACF gives a robust measure.  In this analysis the ACFs of bursts were only corrected for the cosmic time dilation. However, since the detectors are sensitive over a finite energy band, effects due to the shift in energy should also be present. Studing a set of 45 bright long bursts, \\citet{Fen95} found that the full-width $W$ of the average ACF (at the $e^{0.5}$ level) depends on the energy $E$ as $W(E) \\propto E^{-0.4}$.  This narrowing of the ACF should partially counteract the time stretching since for large redshifts the energy window of the instrument will see photons emitted at higher mean energies.  Furthermore, due to the trigger threshold bursts detected at high redshifts are more luminous. There are indications that the pulse width, and therefore the ACF width, correlates with the luminosity \\citep{LBP00}. One should consider also that because of the energy shift bursts at high redshifts are detected at earlier stages. If the local ACF has an approximately constant width $\\tau'_o$ (for each width class) these effects should produce a deviation from linearity in Fig.~\\ref{corr}.  Since no important deviation is observed, the net combined effect must be small. To explore how sensitive our results are to such effects we will assume now that the local width of the ACF is given by $\\tau'_{o} =\\tau_{o}/(1+z)^{1+a}$, where the index $a$ takes into account additional redshift dependencies. Figure~\\ref{index_a} shows the relative dispersion of the width $\\tau'_{o}$ for each set as a function of $a$. The dispersion minima occur at small $a$ index values in both cases, with $a_{\\rm min}^{(b)}=-0.05\\pm0.05$ and $a_{\\rm min}^{(n)}=-0.13\\pm0.23$ for the broad and narrow width sets respectively. The difference between the mean values of each set versus $a$ is also shown in Fig.~\\ref{index_a}. It peaks at $a=-0.05$ with $8.8\\sigma$, where now $\\sigma$ is the total standard deviation. Noteworthy, the {\\it gap} between sets remains larger than $3\\sigma$ over a large range ($-0.4<a<0.3$), indicating how robust the bimodality result is to any additional correction.    \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{fig6.ps}} \\caption{Variation of the dispersion of the ACF width with additional redshift dependencies. The local values are calculated assuming $\\tau'_{o} =\\tau_{o}/(1+z)^{1+a}$, where $a$ is an unknown correction index. The left ordinate shows the relative dispersion of the broad ({\\it dotted line}) and narrow ({\\it dashed line}) sets. Dispersion minima occur at $a \\ll 1$ in both cases. The right ordinate shows for the {\\it solid line} the difference between the mean values of each set ($\\mu_b$, $\\mu_n$) in number of total standard deviations ($\\sigma_b+\\sigma_n$). The {\\it gap} between sets remains above $3\\sigma$ over a large range ($-0.4<a<0.3$), indicating the robustness of the bimodality result to additional corrections. \\label{index_a}} \\end{figure}  The practical use of the proposed empirical relations requires a criterion to decide to what width class a burst belongs. For a width $\\tau_o \\lesssim 7.5$ s the two classes do not overlap (see Fig.~\\ref{corr}). On the other hand, if extrapolations are valid, $\\tau_o \\gg 20$ s implies for the narrow width class unrealistically large $z$.  No burst spectral characteristic was found to correlate with the width class. The same was true for the luminosity and total energy release, but again larger samples should be studied to be conclusive. The mentioned luminosity correlations could give a first $z$ estimation to determine the width class, and then it will be possible using the ACF to obtain a second better and independent estimation. Since the two classes are separated by more than a factor 5, the selection should not depend in practice on the assumed cosmological parameters, and therefore the ACF width-redshift correlation could be used in addition to constrain them.  To investigate the combined use of these correlations, the ACFs of a BATSE burst sample studied by \\cite{FRR00} were calculated. They estimated the redshifts of 220 GRBs using a power-law fit of the luminosity $L$ versus the variability $V$ based on only 7 bursts. Therefore they obtained the best-fit parameters for their model $L \\propto V^{\\alpha}$ with large uncertainties.  Furthermore, the best-fit parameters vary considerably if GRB 980425 is included in the calibration, and this burst was excluded in the derivation of the published redshifts. Figure~\\ref{histo} shows the logarithmic distribution of the observed ACF widths $\\tau_o$ derived for this sample ({\\it dashed-line}) and the distribution of the widths $\\tau'_{o}$ corrected for time dilation ({\\it solid-line}). In Fig.~\\ref{histo} those cases where the iterative method to determine $z$ diverged have been excluded, as well as bursts with observed ACF widths $\\tau_{o}<2.4$ s, but our conclusions do not depend on this selection.  While the uncorrected distribution seems unimodal, the corrected one appears bimodal, in qualitative agreement with the results of \\S~\\ref{results}. The probability $p$ that a statistical fluctuation could produce this feature was estimated.  Assuming an underlying log-normal distribution, a conservative estimation gives $p<0.02$. The $\\tau'_{o}$ distribution is considerably broad, but that is to be expected given the large spread found in the luminosity-variability correlation. The median values for the two width subsets are $\\sim$$0.7$ s and $\\sim$$4.0$ s, both approximately a factor 2 smaller than the mean values of the distributions shown in Fig.~\\ref{averages}. However, this discrepancy can be accounted considering the uncertainties.  Based on the analysis of 20 bursts with known redshifts \\cite{rei01} estimated $\\alpha=3.3^{+1.1}_{-0.9}$.  In particular, the exclusion of GRB 980425 will overestimate the exponent $\\alpha$ and consequently the redshifts, given smaller $\\tau'_{o}$ corrected widths.  The ACF width-redshift correlations described inhere will need to be confirmed by a larger statistical sample. Additionally, the lag and variability luminosity correlations need to be known for a larger redshift range to avoid uncertain extrapolations.  The close agreement between the ACFs using data from the past mission BATSE and the presently operating Konus and INTEGRAL will allowed to continue improving the statistic of this work. In combination with the luminosity correlations we should be able to construct a GRB-based Hubble diagram (i.e., a luminosity distance versus $z$ plot) for high $z$, following a procedure similar to that of \\citet{sch03}. To obtain such diagram would have important implications in cosmology studies. Ongoing efforts in this direction will be presented in the near future.   \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{fig7.ps}} \\caption{Distribution of the ACF width at half-maximum for a sample of 170 BATSE GRBs. {\\it Dashed-line:} logarithmic histogram of the width $\\tau_{o}$ (in seconds). The distribution seems unimodal. {\\it Solid-line:} The width corrected for cosmic time dilation $\\tau'_{o}$, using estimated redshifts derived from the luminosity-variability correlation \\citep{FRR00}. The distribution appears now bimodal, with median values approximately $0.7$ s and $4.0$ s for each subset. \\label{histo}} \\end{figure}"
    },
    "0402/astro-ph0402661_arXiv.txt": {
        "abstract": "{Gravitational radiation that propagates through an inhomogeneous mass distribution is subject to random gravitational lensing, or scattering, causing variations in the wave amplitude and temporal smearing of the signal.  A statistical theory is constructed to treat these effects. The statistical properties of the wave amplitude variations are a direct probe of the power spectrum of the mass distribution through which the waves propagate.  Scattering temporally smears any intensity variations intrinsic to a source emitting gravitational radiation, rendering variability on time scales shorter than the temporal smearing time scale unobservable, and potentially making the radiation much harder to detect.  Gravitational radiation must propagate out through the mass distribution of its host galaxy before it can be detected at the Earth.  Plausible models for the distribution of matter in an $L_*$ host galaxy suggest that the temporal smearing time scale is at least several milliseconds due to the gas content alone, and may be as large as a second if dark matter also scatters the radiation.  The smearing time due to scattering by any galaxy interposed along the line of sight is a factor $\\sim 10^5$ times larger.  Gravitational scattering is an excellent probe of matter on parsec and sub-parsec scales, and has the potential to elucidate the nature of dark matter. ",
        "introduction": "Current and planned gravitational wave detectors have the potential to identify sources of gravitational radiation out to extragalactic distances and, for some types of sources, out to cosmological distances (Cutler \\& Thorne 2002) over a wavelength range spanning more than 8 orders of magnitude. Instruments such as the Laser Interferometer Gravitational Wave Observatory (LIGO) are sensitive to radiation in the range $10-10^4$\\,Hz, while the planned Laser Interferometer Space Antenna (LISA) experiment will be sensitive to radiation from frequencies $10^{-1} -10^{-4}$\\,Hz.  Other methods of detection, such as those due to pulsar timing observations, are sensitive to nano-Hertz gravitational radiation (Sazhin 1978, Detweiler 1979).  As the detection of gravitational radiation requires extremely sensitive apparatus, consideration has been devoted as to whether gravitational lensing by intervening massive objects might focus the radiation and enhance its detection likelihood (Ruffa 1999).  These lensing calculations (Ohanian 1974, Bliokh \\& Minakov 1975, Bontz \\& Haugan 1981, Thorne 1983, Deguchi \\& Watson 1986a) are performed in the regime of physical optics. Diffractive effects are important because the wavelength of gravitational radiation is large compared to the Schwarzschild radius of any lensing object. For a wave of frequency $\\nu$, diffractive effects are important for lens masses $\\la 10^5 \\nu^{-1} \\,$M$_\\odot$ (Takahashi \\& Nakamura 2003).  Thus lensing by almost any intervening single object must be treated using physical optics.  Moreover, even for lens masses in which geometric optics is valid, it is preferable to work in the framework of physical optics because it allows one to compute certain quantities related to the wave amplitude that are in principle measurable with gravitational wave detectors but are not calculable under the framework of geometric optics.  Foregoing lensing treatments have largely been limited to the effects of a single massive object on the wave amplitude.  The approach has hitherto been from a deterministic rather than a statistical point of view, in that the mass profile of the lensing object is well specified.  However, a statistical approach is more appropriate under many circumstances.  The argument for a statistical approach stems from the large number of objects that are likely to contribute to the wave amplitude in many astrophysically interesting situations.  There are two senses in which a statistical approach is applicable to the gravitational lensing of gravitational radiation.  One is when a large number of lensing objects can drift in front of the line of sight to a source, so that an average over time would involve a sufficiently large number of objects to merit an average over an ensemble of lensing objects.  A second, more restrictive, sense applies when a large number of objects contributes to the observed wave amplitude at any instant.  This latter sense is likely to apply in practice.  The wave amplitude of a scattered signal is determined by the phase fluctuations caused by lensing objects in a region at least as large as the first Fresnel zone (see, e.g., Narayan 1992).  For gravitational lensing occurring at distance $D$ from the observer, the Fresnel radius $r_{\\rm F}  \\sim \\sqrt{c D/2 \\pi \\nu}$, is of order a parsec for lensing of $\\nu \\sim 1\\,$Hz radiation at cosmological distances.  Thus the area defined by the first Fresnel zone encompasses contributions from a large number of objects.  Consider, for example, the lensing influence of stars on the gravitational radiation passing through the spiral arm of a galaxy.  Even the volume of a line of sight orthogonal to the disk axis of a spiral galaxy with radius $r_{\\rm F}\\approx 1\\,$pc would typically encompass many thousands of stars.  The argument for a statistical approach is strengthened when one considers the possible contribution of gas and of dark matter, both of which may be inhomogeneous on sub-parsec scales.   Scattering gives rise to a number of important effects that influence the detectability of signal.  The two that motivate the present treatment are (i) the focusing and defocusing of radiation by phase fluctuations, leading to temporal variability in the observed wave amplitude, and (ii) the temporal smearing of intrinsic intensity\\footnote{Throughout the text we refer to the quantity defined by $\\tilde \\phi(\\nu)$ in Eq. (\\ref{Fresnel}) as the (complex) wave amplitude at frequency $\\nu$, while the quantity $| \\tilde \\phi(\\nu) |^2(t)$ is the intensity of the radiation at frequency $\\nu$ and time $t$.  Thus intensity variations refer to changes in the quantity $|\\tilde \\phi (\\nu) |(t)$ with time.} fluctuations in the source.  Temporal fluctuations in the wave field occur whenever there is movement of the lensing system relative to the line of sight to the source of gravitational radiation.  Fluctuations in the gravitational potential drifting transverse to the line of sight induce both phase and amplitude fluctuations in the gravitational radiation, leading to temporal decorrelation of the wave field on sufficiently long time scales.  This effect is important if the decorrelation time scale associated with gravitational scattering is short compared to the time scale on which the wave field would oscillate in the absence of scattering.  In this paper we concentrate on fluctuations in the wave amplitude itself, while in a second paper we consider fluctuations in the intensity.  Temporal smearing is due to the fact that the radiation reaching an observer can arrive from a number of different angles.  This effect is significant when the phase fluctuations induced by the lensing are sufficiently large that radiation from outside the Fresnel radius is scattered toward the observer.  The radiation then arrives from a number of different regions in the scattering region, in an effect known as multipath propagation. There is a range of propagation times associated with the range of angles over which the radiation arrives. The effect is important because it potentially smears out the intrinsic intensity variations of a source.  For instance, time variations of the source would be undetectable if the intrinsic variability time scale were significantly shorter than the scatter broadening time.  The scattering of gravitational radiation potentially constitutes an exceptional probe of the distribution of matter out to the distances to which sources can be detected.  In this context it is useful to compare gravitational scattering to the interstellar scattering (ISS) of radio wavelength radiation by the diffuse ionized component of our Galaxy's interstellar medium.  Both scattering phenomena require a statistical approach and occur in a regime in which physical optics is important.  ISS yields information on the power spectrum of turbulent fluctuations in the ionized plasma distribution of our own Galaxy over a range of scales spanning some five orders of magnitude in wavenumber (Armstrong, Rickett \\& Spangler 1995).  This information is gleaned primarily from pulsars which are effectively point sources, and thus constitute excellent probes of fluctuations in the interstellar medium.  Other radio sources are less subject to ISS, and their larger angular diameters severely truncate the range of wavenumbers over which the power spectrum can be measured.  Gravitational wave sources are the analogues of pulsars in the regime of gravitational scattering. They emit coherently and possess small angular diameters so the source structure does not influence the scattering characteristics, unlike most sources of electromagnetic radiation that are subject to gravitational lensing.  Thus the scattering of gravitational radiation is a fine-scale probe of the distribution of dark matter in the local universe.  The close relationship between gravitational and interstellar scattering implies that many phenomena identified in this paper are analogous to effects previously identified in the context of interstellar scintillation.  For example, the temporal broadening of gravitational radiation is analogous to the temporal broadening of pulsar radiation, and it poses similar limitations on source detectability (see the review by Rickett 1977 and references therein).  The outline of the paper is as follows.  In Sect.\\,2 we review the effect of an arbitrary mass distribution on the amplitude of a gravitational wave.  In  Sect.\\,3 we compute the temporal variability of the wave amplitude by considering the power spectrum of the wave amplitude fluctuations.  The effect of temporal broadening on the scattered signal is considered in Sect.\\,4.  Estimates of the power spectrum of mass fluctuations due to the various constituents of a galaxy are presented in Sect.\\,5.  In Sect.\\,6 we estimate the magnitude of scattering effects and discuss their relevance to the observability of gravitational waves.  The conclusions are presented in Sect.\\,7. ",
        "conclusions": "The gravitational potentials of mass fluctuations encountered by gravitational radiation as it propagates toward Earth perturb its wave front, causing an observer to perceive variability in the wave amplitude and temporal smearing of the signal.  We employ a statistical approach to relate these effects to the underlying power spectrum of mass density perturbations that drive the scattering.  This approach is motivated by the fact that gravitational radiation produced by extragalactic sources is likely to encounter the gravitational potentials of many objects along its path toward Earth.  Mass fluctuations along the line of sight to a source of gravitational radiation induce variations in its wave amplitude as it propagates toward Earth.  The characteristic fluctuation time scale is the time on which the wave amplitude fluctuations decorrelate.   This occurs over the interval over which the root mean square phase difference along the line of sight to the source changes by one radian. This time scale is frequency dependent because the phase induced by gravitational perturbations is linearly proportional to frequency.  Wave amplitude variability therefore occurs more rapidly at high frequencies. We consider lensing by gas and dark matter, and estimate that decorrelation occurs on a time scale $\\ga 10^3 \\nu^{-1}\\,$days.  However, even when temporal variations are too slow to be discerned, the random amplification of the radiation due to focusing and defocusing of the radiation still alters the detectability of the signal.  The most important scattering effect relates to the temporal smearing of scattered radiation, as it has a direct bearing on the detectability of any source whose intensity, the square of the wave amplitude, varies with time.  The lightcurve of a temporally smeared source in a stochastic medium is the intrinsic intensity lightcurve of the source convolved with a function whose width is the temporal broadening time scale, $t_D$.  Temporal smearing arises in a scattering medium because radiation scattered through an angle and deflected back towards towards the line of sight takes longer to reach an observer than radiation propagating directly along the line of sight.  The magnitude of this effect depends on the power spectrum of the mass fluctuations and on the lensing geometry.  The power spectrum depends on the particular distribution and density of stars, gas and dark matter along the line of sight.  Gravitational radiation is always subject to scattering, as it must escape through the mass distribution of its host galaxy and propagate through the mass distribution of the Milky Way in order to be detected.   We consider the consequences of lensing due to a power law distribution of mass fluctuations, due to either gas or dark matter; we also consider a simple model due to lensing by a collection of stars. These estimates suggest that temporal smearing due to the host galaxy is dominated by the dark matter distribution, and is important for radiation whose intensity varies on a time scale less than a second.  The temporal smearing time is independent of the observing frequency.  Temporal smearing effects scale linearly with the effective distance to the scattering material, $D_{\\rm eff}=D_L D_{LS}/D_S$, so the smearing time due to lensing by an external galaxy is exceptionally large.  For lensing at cosmological distances, $D_{\\rm eff} \\sim 1\\,$Gpc, the temporal smearing time is estimated to be of order $t_D\\sim 10^5\\,$s.  As the distribution of dark matter on small scales relevant to temporal broadening is ill-constrained, we caution that the foregoing estimates are highly uncertain. When dark matter is absent, even the gas content of a galaxy can appreciably broaden the intrinsic intensity variations of a source.  The broadening time associated with the gas content of the host galaxy is of order milliseconds, while the time associated with scattering due an intervening system is of order $10^3$\\,s.   The circumstances in which temporal smearing poses a serious limitation on the detectability of intrinsic source variations depends on the distribution of scattering material.  The description of temporal smearing presented here is valid when one is justified in treating the phase fluctuations using a statistical approach.  However, it fails when only a few inhomogeneities contribute to the scattering at any one instant, at which point the observed intensity lightcurve becomes the sum of a small number of delayed copies of the intrinsic intensity lightcurve.  Thus its validity varies according to the line of sight and the scale on which matter distributed along it is inhomogeneous.  Clearly, the scale on which dark matter is inhomogeneous is critically important.  If it is inhomogeneous on scales smaller than the Fresnel scale, as is expected for CDM fluctuations, then the statistical approach adopted for temporal broadening here is well justified.  Thus, the scale on which dark matter is distributed is of critical importance in determining the detectability of gravitational radiation.  Conversely, temporal smearing measurements of gravitational radiation place strong constraints on the distribution and hence the nature of dark matter along any line of sight through which a source of gravitational radiation is detected."
    },
    "0402/astro-ph0402382_arXiv.txt": {
        "abstract": "We present computations of the metals photoionization structures of pressure supported gas clouds in gravitationally dominant dark-matter mini-halos in the extended Galactic halo or Local Group environment. We consider low-metallicity ($0.1-0.3$ solar) clouds that are photoionized by the present-day metagalactic radiation field, and which are also pressure-confined by a hot external medium. We study the properties of $\\Lambda$CDM Burkert (and also NFW) mini-halos with characteristic circular velocities from 12 to 30~km~s$^{-1}$ ($10^8$ -- $2\\times 10^9$ $M_\\odot$). We present results for the volume densities and projected column densities of low-ionization metal atoms and ions, \\ion{C}{2}, \\ion{N}{1}, \\ion{O}{1}, \\ion{Si}{2}, and \\ion{S}{2}, produced in the neutral cloud cores, and high-ions, \\ion{C}{4}, \\ion{N}{5}, \\ion{O}{6}, \\ion{Si}{3}, \\ion{Si}{4}, and \\ion{S}{3} produced in the ionized shielding envelopes. We examine the possible relationships between the compact \\ion{H}{1} high-velocity clouds (CHVCs), dwarf galaxies, and high-velocity ionized \\ion{C}{4} absorbers, within the context of photoionized mini-halo models for such objects. In pressure-confined ($P/k \\sim 50$ cm$^{-3}$ K) mini-halo models for the CHVCs the photoionization states are much lower than in the \\ion{C}{4} absorbers. However, for lower bounding pressures $\\lesssim 1$ cm$^{-3}$ K, in more massive $\\sim 2\\times 10^9$ $M_\\odot$ dwarf galaxy scale halos, the photoionization states in the outer envelopes resemble those observed in the \\ion{C}{4} HVCs. An important exception is \\ion{O}{6} for which our photoionization models underpredict the relative abundances by about an order-of-magnitude, suggesting that additional processes are at play as concluded by Sembach et al.~(2000; 2003). For the cloud size-scales expected in median $\\Lambda$CDM Burkert halos, a gas metallicity of $0.3$ solar is required to produce absorption-line strengths comparable to those in the \\ion{C}{4} HVCs. We argue that fully ionized and starless ``dark galaxies'' could be detectable in the local universe as UV metal line absorbers with ionization states similar to the \\ion{C}{4} absorbers. ",
        "introduction": "\\label{introduction}  Recent ultraviolet (UV) absorption-line observations have revealed the existence of highly ionized ``high-velocity'' gas clouds along lines of sight to extragalactic UV continuum sources (Sembach et al.~1999, 2000, 2002, 2003; Murphy et al.~2000; Wakker et al.~2003). These gas clouds may be the ionized components, or counterparts, of the well studied neutral hydrogen (\\ion{H}{1}) high-velocity clouds (HVCs) (Muller, Oort \\& Raimond 1963; Wakker \\& van Woerden 1997; Putman et al.~2002; de Heij, Braun \\& Burton 2002), objects which are usually defined as \\ion{H}{1} systems with radial velocities inconsistent with differential rotation in the Galactic disk. The origin and nature of the HVCs remains controversial. One attractive possibility is that at least some of the high velocity gas, both neutral and ionized, represents remnant ``infalling'' material associated with galaxy formation processes (Oort 1970; Blitz et al.~1999; Braun \\& Burton 1999; Sternberg, McKee \\& Wolfire 2002).  In the UV studies, many high-ionization metal absorption lines have been detected, including \\ion{Si}{3} $\\lambda~1206.5$, \\ion{Si}{4} $\\lambda\\lambda~1393.8, 1402.8$, and \\ion{C}{4} $\\lambda\\lambda~1548.2, 1550.8$, observed with GHRS and STIS on board HST, and more recently \\ion{O}{6} $\\lambda 1031.9$ absorption observed with FUSE. Some of the ionized UV absorbers can, in fact, be identified with the large ($\\sim 100$ deg$^2$) and nearby ($\\lesssim 50$ kpc) \\ion{H}{1} HVC ``complexes'', such as the Magellanic Stream, or Complex C (Sembach et al.~2003). However, some absorbers, such as those observed towards Mrk 509 and PKS 2155-304 (Sembach et al.~1999; Collins, Shull \\& Giroux~2004) do not appear directly linked to such structures, and could be clouds at much larger distances, perhaps distributed within the Local Group. The ``high-ion absorbers'' may be the local analogs of metal-line absorbers observed in damped Ly$\\alpha$ systems at high redshifts (Lu et al.~1996; Prochaska \\& Wolfe~1997; Boksenberg, Sargent \\& Rauch~2003; Maller et al. 2003).  Sembach et al.~(1999) suggested that photoionization by metagalactic radiation in low density clouds may be responsible for the \\ion{C}{4} observed towards Mrk~509 and PKS~2155-304. However, the detection of highly ionized \\ion{O}{6} along these sight-lines strongly suggests that additional processes, such as collisional ionization in conductive interfaces or in thermally unstable cooling gas, are important (Heckman et al.~2002; Fox et al.~2003).  In this paper we reexamine the properties of purely photoionized clouds, within the context of dark-matter ``mini-halo'' models for the compact \\ion{H}{1} high-velocity clouds (CHVCs) and dwarf galaxies. The CHVCs are of particular interest. They are kinematically distinct objects, with typical angular diameters  $\\lesssim 1^\\circ$ as determined by the recent HIPASS and Leiden-Dwingeloo 21 cm surveys (Putman et al.~2002; de Heij et al.~2002), and they may be at larger characteristic distances from the Galaxy disc than are the large HVC complexes. In Sternberg, McKee \\& Wolfire (2002; hereafter SMW) we explored the hypothesis (Blitz et al.~1999; Braun \\& Burton 1999) that the population of CHVCs traces dark-matter (DM) substructure (Moore et al.~1999; Klypin et al.~1999) in the Galactic halo or in the more extended Local Group environment. In this picture the CHVCs are essentially the lower-mass ``cousins'' of the Local Group dwarf galaxies, and represent objects in which gas is confined to low-mass dark-matter ``mini-halos''. In SMW we focused on the phase properties of hydrogen gas, ionized/neutral/warm/cold, of hydrostatic gas clouds in gravitationally dominant DM halos, including the effects of photoionization and heating by the metagalactic radiation field, and pressure confinement by an external hot medium. Explicit models were constructed for the observed distributions of warm \\ion{H}{1} in the Local Group galaxies Leo~A and Sag~DIG. The \\ion{H}{1} boundaries in these objects are consistent with a cutoff due to photoionization by the metagalactic field.  The aim of this paper is to study the associated metals photoionization structures and resulting absorption-line properties of such ($10^8 - 2\\times 10^9$ $M_\\odot$) mini-halo clouds. We consider photoionization by the present-day metagalactic field, which is the dominant source of radiation at distances $\\gtorder~150$~kpc from the Galactic disk (SMW). The calculations and models we construct are similar to those presented by Kepner et al.~(1999) in their mini-halo models for the QSO absorption-line systems. Our models are more directly applicable to CHVCs in the outer Galactic halo and dwarf galaxies in the Local Group. In particular, we determine the conditions required for the formation of highly ionized absorbers similar to the \\ion{C}{4} HVCs.   In \\S2 we discuss the basic parameters of our mini-halo models, and describe the methods we used to compute the metals photoionization structure. In \\S3 we present auxiliary computations of the abundances of highly ionized metal species in uniform density and optically thin clouds that are photoionized by the present-day metagalactic radiation field. In \\S4 we present photoionization calculations for mini-halo clouds for a range of dark-matter and gas masses, metallicities, radiation fields and bounding gas pressures. We compare our results to the \\ion{C}{4} absorbers towards Mrk~509 and PKS~2155-304. We summarize in \\S5. ",
        "conclusions": "In this paper we study the distributions and relative abundances of metal ions in gas clouds that are confined in low-mass dark-matter mini-halos in the extended Galactic halo or Local Group, and which are photoionized by the present-day metagalactic radiation field. We consider spherical and hydrostatic clouds, with gas temperatures $\\sim 10^4$ K, which are also subjected to external bounding pressures provided by a hot medium. We compute the distributions of low-ions, \\ion{C}{2}, \\ion{N}{1}, \\ion{O}{1}, \\ion{Si}{2}, and \\ion{S}{2}, produced primarily in the optically thick WNM cores, and high-ions, \\ion{C}{4}, \\ion{N}{5}, \\ion{O}{6}, \\ion{Si}{3}, \\ion{Si}{4}, and \\ion{S}{3}, produced in the outer WIM shielding envelopes of the clouds. We study the ionization properties of the compact \\ion{H}{1} high-velocity clouds, as photoionized pressure-confined gas clouds in $\\sim 10^8$ M$_\\odot$ mini-halos. We then focus on the metal ion distributions in more massive $\\sim 2\\times 10^9$ M$_\\odot$ halos at the dwarf-galaxy scale. In particular, we determine conditions for the formation of highly-ionized components, with ionization states similar to the high-velocity \\ion{C}{4} absorbers observed by Sembach et al.~(1999) and Collins et al.~(2004) towards Mrk~509 and PKS~2155-304.  First, we address the question of whether the \\ion{C}{4} HVCs could be the ionized counterparts or envelopes of pressure-confined \\ion{H}{1} CHVCs, which are photoionized by metagalactic radiation, where external bounding pressures $P/k \\sim 50$~cm$^{-3}$~K are required to bind, compress, and neutralize the \\ion{H}{1} clouds. We find that in such objects the ionization parameter is too low to allow efficient production of \\ion{C}{4} and other high-ions. The dominant ionization stages in these objects remain \\ion{C}{2}, \\ion{N}{1}, \\ion{O}{1}, \\ion{Si}{2}, and \\ion{S}{2}. If the CHVCs are pressure-confined circumgalactic objects (as favored by SMW) then significant columns of high-ions require alternative ionization mechanisms, as suggested by Fox et al.~(2003) for the relatively nearby Complex C. One attractive possibility to be considered is the role of non-equilibrium ionization in the conductive interfaces that must exist between the WIM envelopes of mini-halo clouds and surrounding hot gas. Our computations, and the recent \\ion{O}{6} detections, confirm the conclusions of Sembach et al.~(2000; 2003) that additional processes are likely affecting the WIM ionization properties in the Mrk~509 and PKS~2155-304 absorbers. In the WNM cores, the column densities of the low-ions and atoms are predicted to be equal to the metal abundance times the \\ion{H}{1} column along the line-of-sight. UV absorption-line observations of such low-ions can therefore provide useful and direct measures of the CHVC metallicities, if photoionization dominates in the cores.  Second, we study the properties of more extended and tenuous photoionized WIM envelopes around more massive, dwarf galaxy scale, mini-halos. In such systems, gravitational confinement of the gas is more effective, and the bounding pressures can be taken to much lower values. The ionization parameters in the WIM may therefore be increased to values that yield large abundances of \\ion{C}{4}, \\ion{S}{3}, \\ion{Si}{4}, and \\ion{Si}{3}, and simultaneously, reduced abundances of low-ions.  We find that ionization states comparable to those in the \\ion{C}{4} HVCs are produced in the WIM envelopes of dwarf galaxies for bounding pressures $P/k \\lesssim 1$~cm$^{-3}$~K. An important exception is \\ion{O}{6}, which is inefficiently produced by photoionization.  In our models, the predicted metal column densities are proportional to the cloud sizes and associated path-lengths. For a given gas velocity dispersion, the cloud size-scale is fixed by the dark-matter gravitational potential. For $\\Lambda$CDM Burkert halos with median dark-matter concentration parameters, we find that the WIM envelopes extend out to 5-10~kpc from the cloud centers, depending on the value of the bounding pressure. \\ion{C}{4} columns of up to $\\sim 10^{14}$~cm$^{-2}$, as observed by Sembach et al.~(1999), would then require metallicities of at least 0.3 times solar for absorption occurring in a single WIM envelope.  We conclude that the WIM envelopes of dwarf galaxies are natural sites for the formation of \\ion{C}{4}, \\ion{S}{3}, \\ion{Si}{4}, and \\ion{Si}{3} via metagalactic photoionization provided $P/k \\lesssim 1$~cm$^{-3}$~K in the surrounding IGM. This is consistent with the finding of Sembach et al. (1999) that such low pressures favor the production of \\ion{C}{4}. Such low pressures appear to be indicated by recent X-ray absorption-line observations of very highly-ionized \\ion{O}{7} and \\ion{O}{8} that may be probing the hot IGM medium in the Local Group (Nicastro et al.~2002; Rasmussen, Kahn \\& Paerels 2003).  We also present results for fully ionized dwarf galaxies, in which the gas mass confined to the mini-halos remains optically thin. In such objects WNM cores do not exist and the \\ion{H}{1} column densities are correspondingly low. Such objects are also expected to be starless given that photoionization was more effective at earlier epochs, at least since reionization (Efstathiou 1992; Kepner, Babul \\& Spergel 1997; Barkana \\& Loeb 1999).  Our computations show that such ``dark galaxies'' could be detectable as UV metal-line absorbers, with ionization states similar to those observed in the \\ion{C}{4} absorbers. However, given the relatively high metallicities required, the gas that was trapped in such halos must have been previously enriched.  The origin, nature, and location of the \\ion{C}{4} HVCs remain uncertain. Sembach et al.~(1999) suggested that the absorbers towards Mrk~509 and PKS~2155-304 may be related to the group of \\ion{H}{1} HVCs observed in the ``Galactic center negative'' (GCN) direction (Wakker \\& van Woerden 1991). Some of these GCN clouds have velocities close to the \\ion{C}{4} HVCs, and are located within 5$^\\circ$ of them. The GCN clouds may be tracing a system of low-surface brightness galaxies (Mirabel \\& Cohen 1979), although an origin in the Magellanic Stream has also been suggested (Wakker \\& van Woerden 1991). Our models suggest that a system of ionized dwarf galaxies, perhaps associated with the GCN \\ion{H}{1} clouds viewed as low-surface brightness galaxies, could, with the exception of \\ion{O}{6}, give rise to UV metal absorption features similar to those observed in the \\ion{C}{4} HVCs."
    },
    "0402/astro-ph0402457_arXiv.txt": {
        "abstract": "The polarization of gamma ray emission scattered off the baryon rich material that collimates a GRB fireball and the coasting speed of the irradiated matter are calculated numerically for different geometries of the radiation source and the collimating wall.  It is shown that when the scattering material is Compton sailing, the direction of maximum polarization is quite generally well within the beaming cone of scattered radiation. As a result, the probability for observing bright highly polarized GRB's increases considerably, provided the Lorentz factor of the coasting matter is not well below 30, and the scattered radiation is highly polarized even when the beam thickness is large compared to $1/\\Gamma$. It is suggested that correlation between polarization and intensity could provide clues as to whether energy flows from matter to photons or the reverse. ",
        "introduction": "Gamma ray bursts (GRB's) are believed to contain about $10^{51}$ ergs in the form of  several $ \\times 10^{57}$ photons averaging  several hundred KeV in energy. On the other hand, they are believed to have Lorentz factors of at least $10^2$, and, hence, at most $\\sim 10^{51}$ protons. For each proton, there must be at least $10^5$ photons/baryon emitted.  The question arises: what determines the number of photons emitted per baryon and  why should the average photon energy in a GRB be consistently close to  the electron rest mass?  A popular model posits that these photons are emitted as synchrotron radiation in an optically thin environment, typically at distances $10^{13}$ to $10^{15}$ cm. However, this begs the question of why a) the synchrotron frequency in the lab frame, which depends on the magnetic field, b) the comoving Lorentz factor $\\gamma'$ of the emitting electrons, c) the bulk Lorentz factor $\\Gamma$, and d) the viewing angle together so frequently conspire to put the peak luminosity near 300 KeV. The fact that this peak energy has been observed to vary from $\\sim 2$ MeV all the way down to $\\sim 10$ KeV notwithstanding, we contend that the typical value of several hundred KeV is remarkable. For example, the strong linear correlation in the HETE II data set (Atteia et al. 2003) between the square root of the isotropic equivalent fluence $E_{iso}$ and the spectral peak $h\\nu_{peak}$, $h\\nu_{peak}/100 KeV \\sim [E_{iso}/10^{52} \\rm{erg}]^{1/2}$, suggests that the softer GRB's are intrinsically weaker and quite possibly merely  the peripheries of GRB jets. The bright bursts show  remarkably little variation in the location of the spectral peak.  A  detailed analysis by Preece et al. (2002) shows that even if the spectral peak distribution could be arranged within the optically thin synchrotron model, the spectra at low energy, which typically show $dn_{\\gamma}/dE \\propto E^{-1}$, are inconsistent with optically thin, efficiently cooling synchrotron emission which would give $dn_{\\gamma}/dE \\propto E^{-3/2}$. (It could be argued that the bulk Lorentz factor $\\Gamma$ is frequently high enough that the synchrotron cooling is inefficient, which would give a harder spectrum, but then it becomes hard to understand how the emission, which depends on a high power of $\\Gamma$  in this regime, is frequently so efficient.) The recent claim of high polarization in GRB's (Coburn and Boggs 2003) has nevertheless strengthened the views of some that the emission mechanism is optically thin synchrotron  (e.g. Lyutikov, Parlev and Blandford 2003). Others have argued that the claimed degree of polarization, 70 to 80 percent (but with considerable experimental uncertainty), is too high to be accommodated by this mechanism, which limits the polarization to at most 56 percent for the case of  a wide beam, $\\Delta\\theta \\gg1/\\Gamma$.  At the time of this writing, the observational situation is, in the opinion of the authors, unresolved on several fronts. The significance of the X-ray flashes is not yet fully understood. The degree of polarization estimated by Coburn and Boggs may be considered uncertain by several factors: e.g. large error bars, the chance proximity of the modulation frequency of the RHESSI detector and the intrinsic variability of GRB021206, including  2s variability in the hardness (Hajdas et al. 2003), and the absence of other bursts to date that could establish how typical this burst was. Much of the motivation for the model discussed in this paper is thus tentative, but it should be clear that working out the predictions theoretically is part of the process of resolving the issues observationally.  An alternative model has been developed in a series of papers (Levinson and Eichler, 1993, 2000, 2003; Eichler and Levinson 1999, 2000, 2003; Eichler  1994, 2003; van Putten and Levinson 2002, 2003) originally and still motivated by the above questions. Its basic assumption is that the extreme baryon purity that must exist in the high $\\Gamma$  GRB outflows is enforced by the  event horizon of a central black hole, so that the original baryonic content in the outflow must vanish. The photosphere near the axis of the outflow is therefore at about the annihilation radius of the pairs, which until then constitute the main component of the fireball plasma.  The fireball's ability to manufacture more photons drops precipitously beyond the e$^{\\pm}$ annihilation radius, so the total photon entropy of the GRB is frozen near its value at pair annihilation, when the average energy per photon is of the order of magnitude of $m_ec^2$. [Actually, the comoving temperature should be about an order of magnitude  (the log of the compactness parameter)  less than $m_ec^2$ if there is perfect thermal equilibrium, but the dissipation introduced by the friction of the fireball with the surrounding walls probably introduces some non-equilibrium signature on the spectrum, and the bulk Lorentz factor is probably at least a few even in the presence of such friction.]  In this class of models, the average photon energy in a GRB, though somewhat dependent upon the geometry of the walls and how the photons bounce off them,  is naturally of order $m_ec^2$ without fine tuning of the parameters. As long as photons can bounce off the walls at large angles and make pairs with other outgoing photons, the dissipation continues, the expansion is accompanied by additional photon production, and  the bulk Lorentz factor remains moderate. As soon as the energy density is below that needed to maintain pairs, transparency sets in and the photon entropy remains nearly constant thereafter.  Yet another class of models posits that GRB's are photons that are upscattered optical/UV photons by relativistic electrons with a Lorentz factor of order 300 (Shaviv and Dar 1995a, b; Lazzati et al. 2000; Dado, Dar, and De Rujula 2003).  In these models, the number of photons is conserved and the energy of the outgoing photons, which is provided by the Compton scattering material, has been increased by a factor of $10^5$ to $10^6$. Thus, the once scattered photons constitute the primary GRB beam. Because the number of photons is conserved, the isotropic equivalent luminosity of the Compton scattered photons is  of order $\\Gamma^4$ times that in the optical/UV photons. If the latter is of order the Eddington luminosity for a $10 M_{\\odot}$ star, then the observed isotropic luminosities for bright bursts, $10^{13}$ to $10^{15}L_{edd}$ require Lorentz factors exceeding $10^3$. Shaviv and Dar (1995) predicted in their model that inverse Compton emission is expected to emerge polarized with a correlation between polarization and other quantities associated with the angle of observation.  This should include overall fluence and peak luminosity, which are down from the peak values by factors of 8 and 16 respectively when the polarization is at a maximum. The high polarization is possible only when the angular width of the beam is narrower than $1/\\Gamma$, and this constrains the solid angle of the beam to be extremely small if high net polarization is to be observed.  The Compton upscattering model for polarization (Dar and Shaviv 1995) should be compared and  contrasted to the ensheathment model of Eichler and Levinson (1999). Both models predicted linear polarization before the announced discovery (Coburn and Boggs 2003) In the ensheathment model, the energy of the photon is not necessarily increased by the scattering and certainly not by a dramatic factor. The photons can thus originate in a much smaller region without violating the black body limit, and short timescales are, in principle, possible even for the once scattered radiation (the criticism of Lazzati et al. on this point is therefore not justified). The   scattering material in the ensheathment model may have a relatively moderate bulk Lorentz factor and the solid angle of the scattered light can subtend a much larger angle than the high $\\Gamma$ scattering material invoked by the Compton upscattering models. It is therefore possible to combine a large (lab frame) scattering angle and high polarization.  It is also easy in the ensheathment model to have a geometry in which only scattered radiation is seen by the observer. As proposed by Eichler and Levinson (1999), the large reduction of the \"on-axis\" luminosity that was apparently achieved by GRB980425/SN1998bw, perhaps 4 or 5 orders of magnitude, comes with only modest reduction of frequency.  It can result from a small fraction of photons being scattered, by a non-relativistic or modestly relativistic scatterer, into such a large angle as to be observable by us. The smooth light curve that was observed for GRB980425, rare among GRB's and far removed from the relation of Atteia et al. (2003), is consistent with the light having reached us via many different paths.  In recent years, the view that GRB's are  collimated (e.g. Levinson and Eichler 1993)has been almost universally accepted, and the collimating matter is probably a host star or a wind. This confirms the assumption that GRB fireballs pass through some sort of surrounding wall in their early stages. This has particular significance for models in which the early fireball is baryon-free, because it provides a source of baryon seeding that could account for afterglow-generating plasma escaping intact to infinity.  In the first place, some baryons are introduced downstream as entrained matter or pick-up neutrons that diffuse into the outflow from surrounding walls of baryonic matter. This matter can be a host star or simply a baryonic wind driven from the accretion disk of the black hole. The baryonic density in this picture has a complicated transverse structure (Levinson and Eichler 2003), the density increases outward, and the photosphere can be close to the pair annihilation radius near the axis, about $10^{10}$ cm, and near the walls at the edge (see Fig. 1).  In this paper, we  consider another source of baryons - the collimating walls themselves  - and note that, due to the impact of the GRB photons  on the walls, relativistic motion of the walls can ensue. This matter can then be considered part of the GRB. It can cause highly polarized gamma ray emission in certain viewing directions and can cause afterglow. It might even be the case that there is a smooth transition from the supernova ejecta to  the GRB fireball. This could account, for example, for the peculiar features of SN 1998bw, which has a huge expansion velocity (c/6) and peculiar abundances.  This was considered by Eichler (2003, hereafter E03) under simplifying geometric assumptions: that the photons emerged from a small, effective point source.   It was argued that the photon pressure would accelerate the  wall material to a Lorentz factor of $1/sin\\chi$, where $\\chi$ is the angle between the photon flow and the wall, and that the polarization of scattered photons off the wall would be greater than 0.6 with considerable probability. Here we relax the point source assumption. In section 2 we consider the terminal Lorentz factor in the presence of a photon field from a finite source, whose shape is characterized by two parameters.  In section 3, we consider the effect of finite source size on the expected polarization. Finally, we consider multiple scattering within a localized neighborhood of the point of photon impact on the sheath.  We demonstrate that in the present model, large polarization (of order 40 percent) is typical even when the beam \"acausal\", i.e. when it has a much larger opening angle than $1/\\Gamma$. This is not the case in most models where the intensity along the line of sight typically receives contributions from a surrounding neighborhood of velocity vectors within the beam, whereby the polarizations of the various contributing beams cancel.  In the limit of axisymmetric contribution around the line of sight, this cancellation would be complete.  In the present model, on the other hand, there is a lower degree of symmetry (hence the calculations are unfortunately more complicated), because the incident radiation arrives from a direction much different (even in the comoving frame) from the velocity vector. ",
        "conclusions": "We have shown, under less simplified assumptions than made in Eichler (2003), that when matter is Compton dragged into a \"sailing\" equilibrium by photons, the polarization can be quite high for  singly scattered photons.  The direction of maximum polarization coincides with the maximum intensity of the polarized component. Finite source size is found to  not qualitatively alter this conclusion. Finite source angular width $\\delta \\theta$, which can nearly eliminate net polarization in most models, is found in the present model to reduce it only somewhat, and to leave about 40 percent polarization for most relevant viewing angles.  In the  picture presented here,  singly scattered photons may in principle compete with unscattered ones, because the photon energy as seen in the lab frame  is not dramatically altered. [This is in contrast to the scenario of Shaviv and Dar, 1995, Lazzati et al. 2000, where the photon energy is dramatically raised by the scattering, so that virtually all the gamma rays are attributed to such scattering.] However, the scattered and unscattered photons do not necessarily arrive at the same time, are not concentrated in the same directions,  and do not necessarily have exactly the same spectrum. [In  some cases, the line of sight does not necessarily admit unscattered photons, e.g. as was argued (Nakamura 1998, Eichler and Levinson 1999) for the case of GRB980425, however, this GRB may be exceptional.] It is thus important that future GRB polarization experiments be sensitive to the polarization as a function of time and energy range.  If the photons are scattered (undergoing \"diffuse reflection\" in the terminology of Chandrasekhar, 1960) off an optically thick sheath, then the maximum polarization corresponds to 47 percent along the direction of maximum intensity of the polarized component. Typically, polarizations of 30 to 47 percent are expected in this situation among the scattered photons. Consistently observed values in this range of polarizations, with a sharp cutoff at 47 percent, might thus be a signature of diffuse reflection.  The situation in which the optical depth of the scatterers is slightly greater than unity remains to be studied in detail. We believe that the results will be not too much different from the case of single scattering, because the path of maximal polarization corresponds to minimum optical depth in a relativistically expanding flow. In any case, they should lie between the two extremes considered in this paper. It might be claimed that an optical depth of order unity is \"fine tuned\". However, in the present picture, the sheath tapers to infinity, so most of its area is in fact comprised of material with an optical depth of order unity. The path of least optical depth is along the direction of motion, because then the relative velocity between the photon and the material, as seen by the lab frame is the smallest (Levinson and Eichler, 2003).  If future polarization studies eventually show some indication of scattering of GRB photons, correlation between maximum polarization and maximum intensity might provide clues as to whether energy is indeed flowing from the scattering material to the photons or whether most of the scattering occurs after Compton equilibrium has been reached, which would suggest that matter is being dragged by the photons rather than upscattering them.  This research was supported by the Israel-U.S. Binational Science Foudnation, the Arnow Chair of Theoretical Astrophysics and Ben gurion University, and by  an ISF grant for a Israeli Center for High Energy Astrophysics."
    },
    "0402/astro-ph0402511_arXiv.txt": {
        "abstract": "Before the launch of the Rossi X-ray Timing Explorer ({\\em RXTE}) it was recognized that neutron star accretion disks could extend inward to very near the neutron star surface, and thus be governed by millisecond timescales. Previous missions lacked the sensitivity to detect them. The kilohertz quasi-periodic oscillations (QPO) that {\\em RXTE} discovered are often, but not always, evident in the X-ray flux. In 8 years {\\em RXTE} has found kilohertz signals in about a fourth of 100 low-mass X-ray binaries (LMXB) containing neutron stars.  The observed power spectra have simple dominant features, the two kilohertz oscillations, a low frequency oscillation, and band-limited white noise. They vary systematically with changes in other source properties and offer the possibility of comparison with model predictions. New information from the millisecond pulsars resolves some questions about the relations of the QPO and the spin.  Coherence, energy spectrum and time lag measurements have indicated systematic behaviors, which should constrain mechanisms. ",
        "introduction": "Soon after the discoveries of Sco X--1 and Cyg X--2, it was realized that accretion onto a neutron star in a binary was a likely source of the X-ray emission. But while clear pulsations were seen in the flux from Hex X--1, these sources exhibited no periodic signal. The possibility was raised that accretion over a long lifetime had spun up the neutron star to frequencies higher than could have been measured in the early observations \\citep{alpar82}. Successive missions strove to increase their sensitivity to higher frequencies. {\\em EXOSAT} and {\\em Ginga} pushed the frontier to about 200 Hz.  The world before {\\em RXTE} is below 200 Hz.  {\\em EXOSAT} discovered timing signals, but quasi-periodic signals rather than the coherent clock of the neutron star. QPO were found in many X-ray sources in the galactic bulge. The frequencies varied in the 1-50 Hz range. {\\em EXOSAT} proportional counter data provided spectral information at the same time. \\citet{HasvdK89} showed that the spectral variations fell in two categories, denoted ``Z'' and ``Atoll'' and that the QPO frequencies depended on the source's position in a plot of hard versus soft ``colors'' or energy ratios. The widths and amplitudes varied also in systematic ways. \\citet{Wijnands01} compiled the {\\em RXTE} version of a figure summarizing the properties of both Z and Atoll sources.  At first the frequency appeared to be positively correlated with the X-ray luminosity and a simple explanation  was attractive, the magnetic beat frequency model \\citep{alpar85}. The accretion rate through the disc, should be stopped eventually, by a magnetosphere due to the neutron star, but closer to the neutron star because the magnetic field was much weaker. The Kepler period of gas in the disk would beat with the spin frequency of the neutron star to cause brightness oscillations. Changing the accretion rate would change the magnetospheric radius, the Kepler frequency at the boundary, and thus the beat frequency. It implied spin rates of 50-350 Hz in several cases \\citep{GL92}.  However the model was not a satisfactory fit to the data from several sources and there was evidence that the luminosity was not a good measure of the accretion rate. The character of bursts and their recurrence rate changed in 4U 1636--53 as it moved through the ``Atoll'' pattern \\citep{vdK90}, while the luminosity did not increase smoothly. In Cyg X--2 \\citep{Has90,Vrtilek90} and Sco X--1 \\citep{Vrtilek91} UV emission decreased as the X-ray flux increased, while the magnetospheric beat frequency model implied it should increase \\citep{Has90}. Nevertheless coherent oscillations were sought \\citep{Vaughan94} and upper limits of less than 0.5 \\% were achieved for frequencies below 200 Hz.  The idea that the magnetic fields of the neutrons stars are $10^{2} - 10^{4}$ lower than the $10^{12}$ G of ``classical'' pulsars was advanced to explain the failure to detect strongly channeled accretion flow that should show up as pulsations and the higher frequencies.     \\begin{figure}  \\resizebox{1.0\\textwidth}{!}{ \\begin{tabular}{lr} \\epsfig{file=swankj_f1a_c.ps,width=2.9in} & \\epsfig{file=swankj_f1b_c.ps,width=2.9in} \\\\ \\end{tabular} }  \\caption{Dynamical Power Spectra of 4U 1728-34 and of Sco X-1. For each the spectra are shown for three consecutive observation intervals. At the top is an averaged Fourier power spectrum for all the intervals. For each source on the right is the light curve. The intervals included two bursts from 4U~1728--34. } \\label{fig:1} \\end{figure}   Kluzniak and Wagoner saw that accretion disks around low magnetic field neutron stars could be very interesting  if the equation of state of nuclear matter meant that neutron stars were inside the innermost stable orbit of orbiting material. The accretion disk could extend down to the innermost stable orbit and be truncated there rather than at the magnetosphere. A signal might even indicate the Kepler frequency of the inner most stable orbit \\citep{Kluz85,Kluz90}. These papers foresaw that signals bearing the imprint of General Relativity could come from these sources.   In 1996, {\\em RXTE} began observing and the first observations of the Atoll source 4U 1728--34 by \\citet{Stroh96} and the Z source Sco X-1 by \\citet{vdK96} showed signals with frequencies in the range that orbits close to neutron stars would have. Figure 1 shows several of the important aspects of the kilohertz QPO discovered in the flux from 21 LMXB. As the count rates vary, the QPO center frequencies vary significantly compared to the widths of the features. The features in the Atoll and the Z source are very similar. The phenomena and the physical models that have been explored during the 8 yr since the discovery have been described in several review articles \\citep[See][]{vdK00,Wijnands01}.    Looking back at why {\\em RXTE} could detect the signals while previous missions did not, the increase in sensitivity came from several factors. The number of sigmas of the detection of a QPO feature can be expressed as  \\begin{equation} n_{\\sigma} = (1/2) S^{2}/(S+B)(rms/S)^{2} \\surd (T/\\Delta \\nu) \\end{equation}  \\noindent Here, S is the source count rate, B the background rate, rms the root mean square variance in S, T the duration of the observation and $\\Delta \\nu$ the width of the QPO feature. This scales as the detector area. The PCA has observed with a maximum of 6250 cm$^2$, compared to {\\em EXOSAT}'s 1600 cm$^2$, but {\\em Ginga} had 4000 cm$^2$, and did not detect kiloHertz oscillations because of insufficient time resolution. Other factors - background, noise, dead time, low duty cycle of observations - have influenced the sensitivity to these phenomena. So far, {\\em RXTE} has been the only instrument to detect them. ",
        "conclusions": "In the complex of QPO observed in data from the LMXB containing neutron stars {\\em RXTE} has found very clearly signatures which are generally characteristic of the sources. The signatures are related to those seen in black hole sources and perhaps to phenomena seen in a few white dwarf systems. But the neutron star systems have exhibited definitive characteristics that allow very quantitative study of the dependences on parameters. Timing is a tool for obtaining spatial resolution and understanding of dynamics. Even without final detailed interpretation, the QPO are clearly entirely consistent with the neutron stars being inside their ISCOs and with the accretion disks penetrating so close to the neutron star as to be affected by General Relativity. The set of signatures is very simple compared to, say, atomic spectra, but intricate enough that the correct explanation of the details will be interesting and important.     \\begin{theacknowledgments} I  thank the many {\\em RXTE} users whose results I discuss here, the {\\em RXTE} Science Operations Center members for the roles they play in making the observations successful, and NASA for its support of all aspects of {\\em RXTE}. \\end{theacknowledgments}"
    },
    "0402/astro-ph0402604_arXiv.txt": {
        "abstract": "We present deep images of dust continuum emission at 450, 800, and 850~$\\mu$m of the dark cloud LDN~1689N which harbors the low-mass young stellar objects (YSOs) IRAS~16293$-$2422A and B (I16293A and I16293B) and the cold prestellar object I16293E. Toward the positions of I16293A and E we also obtained spectra of CO-isotopomers and deep submillimeter observations of chemically related molecules with high critical densities (HCO$^+$, H$^{13}$CO$^+$, DCO$^+$, H$_2$O, HDO, H$_2$D$^+$). To I16293A we report the detection of the HDO $1_{01}$--$0_{00}$ and H$_2$O $1_{10}$--$1_{01}$ ground-state transitions as broad self-reversed emission profiles with narrow absorption, and a tentative detection of H$_2$D$^+$ $1_{10}$--$1_{11}$. To I16293E we detect weak emission of subthermally excited HDO $1_{01}$--$0_{00}$. Based on this set of submillimeter continuum and line data we model the envelopes around I16293A and E. The density and velocity structure of I16293A is fit by an inside-out collapse model, yielding a sound speed of $a$=0.7~km~s$^{-1}$, an age of $t$=(0.6--2.5)$\\times 10^4$~yr, and a mass of 6.1~M$_\\sun$. The density in the envelope of I16293E is fit by a radial power law with index $-1.0\\pm 0.2$, a mass of 4.4~M$_\\sun$, and a constant temperature of 16~K. These respective models are used to study the chemistry of the envelopes of these pre- and protostellar objects.  We made a large, fully sampled CO $J$=2--1 map of LDN~1689N which clearly shows the two outflows from I16293A and B, and the interaction of one of the flows with I16293E. An outflow from I16293E as reported elsewhere is not confirmed. Instead, we find that the motions around I16293E identified from small maps are part of a larger scale fossil flow from I16293B. Modeling of the I16293A outflow shows that the broad HDO, water ground-state, and CO $J$=6--5 and 7--6 emission lines originate in this flow, while the HDO and H$_2$O line cores originate in the envelope. The narrow absorption feature in the ground-state water line is due to cold gas in the outer envelope. The derived H$_2$O abundance is $3 \\times 10^{-9}$ in the cold regions of the envelope of I16293A ($T_{\\rm kin}<14$~K), $2\\times 10^{-7}$ is warmer regions of the envelope ($>14$~K), and $10^{-8}$ in the outflow. The HDO abundance is constant at a few $\\times 10^{-10}$ throughout the envelopes of I16293A and E. Because the derived H$_2$O and HDO abundances in the two objects can be understood through shock chemistry in the outflow and ion-molecule chemistry in the envelopes, we argue that both objects are related in chemical evolution. The [HDO]/[H$_2$O] abundance ratio in the warm inner envelope of I16293A of a few times 10$^{-4}$ is comparable to that measured in comets. This supports the idea that the [HDO]/[H$_2$O] ratio is determined in the cold prestellar core phase and conserved throughout the formation process of low-mass stars and planets. ",
        "introduction": "}  IRAS~16293$-$2422 is one of the best studied low-mass young stellar objects (YSOs) \\citep{bla94,evd95,cec00}. It is deeply embedded in the LDN~1689N cloud in Ophiuchus (distance 160~pc) and classified as an extreme Class I/Class 0 YSO \\citep{Lad91}. It is a low-luminosity (27~$L_\\sun$, \\citealt{Walk86}) binary system (hereafter, I16293A and I16293B) with a separation of 840~AU ($5''$) and individual stellar masses of 0.5~M$_\\sun$ \\citep{Mun92}. Two powerful outflows emanate from the system.  At least three physically and chemically different regions can be recognised in a $20''$ (3000~AU) beam \\citep{evd95}: (i) a cold, relatively low-density outer envelope ($T_{\\rm kin} \\simeq $10--20~K, $n$(H$_2)\\simeq 10^4 $--$10^5$ cm$^{-3}$) which marks the transition into the extended parental cloud LDN~1689N; (ii) a warmer circumbinary envelope of size $\\sim 2000$~AU ($T_{\\rm kin}\\simeq 40$~K, $n$(H$_2)\\simeq 10^6$--$10^7$~cm$^{-3}$); and (iii) a warm, dense core of about 500--1500~AU ($T_{\\rm kin} \\ge 80$ K, $n$(H$_2)\\simeq 10^7$ cm$^{-3}$) which traces the interaction of the outflow and the stellar radiation with the inner part of the circumbinary envelope. This picture emerged from excitation analysis of species tracing these distinct regions and conditions, assuming uniform temperatures and densities for each. Continuous descriptions of the density and temperature as function of radius in the collapsing envelope were developed by \\citet{cec00} and \\citet{Scho02} on scales from 30 to 5000~AU. In this Paper, we present new observations and modeling, that allows us to independently derive the density and temperature structure in the extended envelope around the (pre-) protostellar cores through continuum photometry and imaging of the dust, and heterodyne spectroscopy of the molecular gas. The resulting density and temperature distribution ranges from the warm inner envelope region (100~AU) to the coldest (12~K) outer envelope regions (7300~AU) and serves as a basis for excitation and abundance analysis of the line data.  Another component of the IRAS~16293$-$2422 system is located 80$''$ east and 60$''$ south of I16293A. This small clump, I16293E, has strong DCO$^+$, NH$_3$, and NH$_2$D emission \\citep{Woo87,Mun90,Sha01} connected to the I16293A core, and significant submillimeter continuum emission from dust \\citep{San94}. Single-dish observations of NH$_3$ in a $40\\arcsec$ beam indicate a temperature $T_{\\rm kin} = 12$~K \\citep{Miz90}. Millimeter-interferometric data, sampling smaller scales and less sensitive to the warmer extended envelope, yield $T_{\\rm kin} \\simeq 8$~K and $n({\\rm H}_2)\\simeq 1\\times 10^5$ cm$^{-3}$ if homogeneous conditions are assumed \\citep{Sha01}. Because no far-infrared (FIR) or submillimeter point source has been found embedded in I16293E, it is classified as a pre-(proto)stellar core. In this Paper, we present submillimeter images that constrain the density structure of I16293E.  It is difficult to study chemical evolution in the early protostellar stages, because during collpase the temperature stays low at $\\sim 10$~K while the density increases to $n($H$_2)\\ge 10^8$ cm$^{-3}$. Under these conditions the line emission of many molecules is dominated by a warmer outer envelope around the prestellar core. Also, many molecules in the cold core will condense on dust grains. Eventually, only H$_2$, H$_3^+$, and their isotopomers HD and H$_2$D$^+$ will remain in the gas phase. This leads to a significant enhancement of H$_2$D$^+$ through the deuterium exchange reaction \\begin{equation}               % {\\rm H}_3^+ +{\\rm  HD} \\rightleftharpoons {\\rm H}_2{\\rm D}^+ + {\\rm H}_2 + \\Delta E, \\label{e1} \\end{equation} since the backward reaction becomes negligible at low temperatures \\citep{Smi82,Her82}. H$_2$D$^+$ is thought to play a pivotal role in the deuteration of molecules \\citep{Wat76,Mil89} and the detection of H$_2$D$^+$ in the Class~0 YSO NGC~1333 IRAS~4A \\citep{Sta99} is an important confirmation of the cold gas-phase chemical networks. The H$_2$D$^+$ enhancement is reflected in the high abundance ratios that are generally observed of, e.g., [DCO$^+$]/[HCO$^+$], [NH$_2$D]/[NH$_3$], [DCN]/[HCN], and [N$_2$D$^+$]/[N$_2$H$^+$] in cold prestellar cores and Class~0 YSOs (e.g., \\citealt{Gue82, Olb85, Woo87a, But95, Wil98, Ber02}). Models predict that in the collapse stage the abundances of many deuterated radicals and molecules increase sharply after their non-deuterated versions start to get heavily depleted on dust grains, then reach a peak and start to decrease as the deuterated molecules too condense on the dust (e.g., \\citealt{Bro89, Rod96}). Reactions with atomic deuterium on the grain surface may further enhance the abundance of deuterated species in the ice \\citep{Tie89,Bro89}. After the formation of a YSO, the condensed molecules may evaporate from grains, resulting in a temporary increase in the gas-phase deuteration fraction. The deuteration may decline after $\\sim 10^4$ years \\citep{Rod96} when higher temperature chemistry becomes effective. In addition, in such hot ($T_{\\rm kin} \\simeq 100$--200 K) regions the reversal of reaction (\\ref{e1}) becomes dominant and very little new fractionation is expected to occur. We report in this Paper observations of several key deuterated species in I16293A and I16293E.  Water is an important species in the early stages of star formation. H$_2$O is the most important hydride molecule in oxygen chemistry and models predict that H$_2$O plays a dominant role in the energy balance and chemical evolution during star formation. Water will have the largest abundance and excitation contrast between the protostellar source and the surrounding cloud (e.g., \\citealt{Cec96, Dot97}). Measurements of the rotational and ro-vibrational transitions of H$_2$O in star-forming regions have recently become available from the Infared Space Observatory (ISO; e.g., \\citealt{Hel96a, cec00}) and the Submillimeter Wave Astronomical Satellite (SWAS; e.g., \\citealt{Ash00, Sne00, Neu00}). In particular the heterodyne velocity-resolved measurements of the ground-state transition of ortho-H$_2$O with SWAS allow a direct determination of the water abundance throughout the envelope. The HDO $1_{01}$--$0_{00}$ ground-state transition at 465~GHz ($E_u= 23.2$ K, $n_{\\rm crit} \\ge 10^8$ cm$^{-3}$ for $T_{\\rm kin} \\le 50$ K, \\citealt{Gre89}) is an excellent tracer of high density gas at low temperatures where the abundance of HDO relative to water is enhanced.  Ground-based observations of the HDO ground-state transition are still very sparse. To date, this transition has only been observed in the high-mass star-forming regions Orion-KL \\citep{Sch91}, W3(OH), and W3(H$_2$O) \\citep{Hel96b}. The latter study indicates that the [HDO]/[H$_2$O] abundance ratio is comparable to that found in hot cores and is not a sensitive indicator of the evolutionary stage in high-mass star formation. This could mean either that the W3 cloud always stayed warm or that the [HDO]/[H$_2$O] ratio retrusn to thermal equilibrium faster than ratios like [DCN]/[HCN] \\citep{Hel96b}. For a low-mass Class~0 YSO like I16293A the situation may be different, because the time scale to reach steady-state may be much longer. In this Paper, we present observations of HDO and H$_2$O of I16293A and E.   In this paper we present a detailed study of the physical and chemical structure of the low-mass star-forming cloud LDN~1689N with particular emphasis on the H$_2$O/HDO chemistry and deuterium fractionation.  We present extensive submillimeter continuum photometry and maps revealing the warm and cold dust around I16293A, B and I16293E (\\S \\ref{Obs_cont}). We report the detection of the ground-state lines of (deuterated) water and a tentative detection of H$_2$D$^+$ toward I16293A as well as weak HDO emission in the prestellar core I16293E (\\S \\ref{Obs_line}). These data are combined with pointed observations of CO and HCO$^+$ isotopomers to determine the temperature and density structure throughout the envelope (\\S \\ref{Results}). We address the two outflows and their origin in \\S \\ref{CO_outflows}. Our deep observations of the HDO and H$_2$O ground-state transitions are used to study the [HDO]/[H$_2$O] ratio throughout the envelope (\\S \\ref{hdo_h2o}).  A simple chemistry network is used to model the deuterium chemistry and in particular the water fractionation (\\S \\ref{Simple_chem_model}).  Finally, we discuss the evolutionary difference between I16293A, I16293B, and I16293E (\\S \\ref{Nature_of}), and summarize the conclusions in \\S \\ref{conclusion}. ",
        "conclusions": "}  We have analyzed our submillimeter dust and continuum data in terms of an excitatiom analysis using a model for the temperature and density structure of the disk, envelope, and outflow of I16293A, B, and E. The spectral lines of H$_2$O, HDO, and H$_2$D$^+$ were subsequently compared with predictions from a simple chemical network.  Figure \\ref{f17} shows a schematic overview of the IRAS 16293-2422 region showing its distinct physical regions (disk, warm dense core, warm envelope, cold envelope, outflow regions, etc.). For each region we indicate the derived physical parameters, and their molecular abundances. For the submillimeter disk of I16293A we derive a dust temperture $T_{\\rm d}=40$ K, a source size of about 800 AU, a total mass of 1.8\\Msun\\ and a luminosity of 16.5 \\Lsun\\ which corresponds to a gas density $n({\\rm H}_2)\\ge 10^9$ cm$^{-3}$.  The envelope of I16293A can be well described by a collapsing envelope model with parameters age $t=(0.6$--$2.5) \\times 10^4$ yr and sound speed $a=0.7$ km s$^{-1}$. In fact the center of the expansion wave, $r_{\\rm CEW}= a t$, is 890--3700~AU from the centre which corresponds to 6--23$''$ at the distance of 160 pc, comparable to the beam size of the SCUBA continuum and submillimeter line emission of the species we used in the analysis.  The outer radius of the envelope is about 7300 AU and the temparature ranges from 115 K in the inner region to 12 K in the outer region. The envelope mass is $M_{\\rm env}({\\rm H}_2)=6.1$ \\Msun. We find that for a YSO such as I16293A the line emission is beter suited than the continuum emission to constrain the infall parameters. The excitation analysis of the envelope yields a constant HDO abundance of $3\\times 10^{-10}$ throuhout the envelope, and an ortho-H$_2$O abundance of $2\\times 10^{-7}$ for the regions with $T_{\\rm kin}>14$ K, while the abundance is lowered by a factor of 150 in the colder regions. From the C$^{17}$O analysis we find that CO is depleted by a facot of 50 for the envelope regions where $T_{\\rm kin}<20$ K.  The line wings of the CO, HDO, and H$_2$O emission can qualitatively be explained by a single outflow model with a turbulent velocity of 2.3 km~s$^{-1}$ and where the maximum outflow velocity is about 6 km~s$^{-1}$.  The excitation analysis of the outflow model indicates that the ortho-H$_2$O abundance in the outflow is about $1 \\times 10^{-8}$ and that the HDO abundance in the outflow is in the range of $3\\times 10^{-9}$ to $3 \\times 10^{-10}$.  A more quantitative comparison requires a multiple outflow model.  We find that only two outflows are present, a NE--SW outflow with a dynamical time scale of about $(3.0$--$3.5) \\times 10^3$ yr which is powered by I16293A a Class~0 source, and an E--W outflow which is a fossil flow with a dynamical time scale of $(6.4$--$7.6) \\times 10^3$ yr driven by I16293B which looks like a young low-luminosity T~Tauri star.  The prestellar object I16293E is well described with an isothermal core of $T_{\\rm kin}=16$ K, a core radius of 8000 AU and a power law density structure of the form $n(r)=n_0 (r/1000~{\\rm AU})^{-p}$ where $n_0=n(R=1000$ AU$)=1.6 \\times 10^6$ cm$^{-3}$. The total mass mass of this core is $M_{\\rm core}({\\rm H}_2)=4.35$ \\Msun.  Our excitation analysis of the core indicate a depletion of CO and HCO$^+$ by a factor of ten relative to the `standard abundances' and that the DCO$^+$ abundance is larger than that of H$^{13}$CO$^+$. Our tentative detection of HDO indicates the presence of a cold condensation in the heart of this core where the HDO abundance is about $2\\times 10^{-10}$.  Our detailed CO mapping shows that I16293E has no outflow, so its evolutionary stage is younger than Class~0. We argue that this core may hide a first hydrostatic core, a so called Class~$-$I object which is deeply embedded in a largely unaffected cold core.  The reason for the extreme deuteration in the L16293N cloud and in particular in I16293A, and I16293E may lie in the fact that I16293AB is a binary system where the YSOs and the outflows differ in age causing a firecracker efffect of the deuteration. In this scenario the older E--W ouflow from I16293B is responsible for the formation of I16293E out of a dense dark core through a slow shock where only deuterated species were desorbed from the grains which, in combination with the low temperatures, amplifies the deuterium enrichment in the gas phase before a possible re-absorption on grains. This strong deuterated core may have been shocked again an a later time by the younger NW--SW outflow from I16293A further boosting the gas-phase deuterium chemistry to extreme values.  It is remarkable that the abundances of many molecules in the prestellar core I16293E and in the warm part of the envelope of I16293A are comparable. In particular, the HDO abundances are similar which indicates that the warm gas in I16293A may have gone through the same cold pre-collapse phase that I16293E is currently in.  A simple chemical gas-phase modeling of the deuterium chemistry requires that CO is depleted by a facor of 50 for $T_{\\rm kin}<22$ K in order to reproduce the observed H$_2$D$^+$ emission from our chemical network applied to the envelope $(T,\\,n)$ structure, and that an H$_2$O abundance of $3 \\times 10^{-7}$ for $T_{\\rm kin}>14$ K and $4\\times 10^{-9}$ for $T_{\\rm kin}<14$ K is required to reproduce the observed spectrum from the chemical model. The [HDO]/[H$_2$O] abundance in the warm inner-envelope of I16293A ($2\\times 10^{-4}$) is close to the value that has been measured in the solar system and supports this. Our chemical modeling indicates that the [HDO]/[H$_2$O] ratio is determined in the cold gas phase prior to star formation, and that this ratio is conserved in the early stages of low-mass star through freeze out on dust grains formation and that the molecules are released unprocessed to the gas phase after the formation of a low-mass YSO. If this warm deuterated gas is subsequently locked-up in a proto-planetary disk before the high temperature gas chemistry causes a de-deuteration, the [D]/[H] ratio of HDO in low-mass pre-solar systems would remain conserved.  To investigate this scenario in detail, sensitive follow-up observations of the HDO and H$_2$O ground-state transitions in a large sample of low-mass YSOs are needed at high spatial and spectral resolution (i.e., interferometry from high altitudes [HDO] and from space [H$_2$O]) and allow a study of the [HDO]/[H$_2$O] abundance evolution from prestellar cores to planets.  The H$_2^{18}$O ground-state transition may be more suitable than the ortho-H$_2$O transition since the ground-state of the latter molecule is at $\\Delta E/k=34.3$ K above ground, while that of H$_2^{18}$O is at 0~K, and the main isotopic form of water is expected to be mostly optically thick in its ortho and para ground-state lines."
    },
    "0402/astro-ph0402118_arXiv.txt": {
        "abstract": "A search of more than 3,000 square degrees of high latitude sky by the Sloan Digital Sky Survey has yielded 251 faint high-latitude carbon stars (FHLCs), the large majority previously uncataloged. We present homogeneous spectroscopy, photometry, and astrometry for the sample. The objects lie in the $15.6 < {\\it r} < 20.8$ range, and exhibit a wide variety of apparent photospheric temperatures, ranging from spectral types near M to as early as F. Proper motion measurements for 222 of the objects show that at least $50\\%$, and quite probably more than $60\\%$, of these objects are actually low luminosity dwarf carbon (dC) stars, in agreement with a variety of recent, more limited investigations which show that such objects are the numerically dominant type of star with C$_{2}$ in the spectrum. This SDSS homogeneous sample of $\\sim110$ dC stars now constitutes $90\\%$ of all known carbon dwarfs, and will grow by another factor of 2-3 by the completion of the Survey. As the spectra of the dC and the faint halo giant C stars are very similar (at least at spectral resolution of $10^3$) despite a difference of 10 mag in luminosity, it is imperative that simple luminosity discriminants other than proper motion be developed. We use our enlarged sample of FHLCs to examine a variety of possible luminosity criteria, including many previously suggested, and find that, with certain important caveats, {\\it JHK} photometry may segregate dwarfs and giants. ",
        "introduction": "Carbon stars (objects with prominent C$_{2}$ in their spectra) have been studied for more than a century \\citep{dun84}, although faint (R $> 13$) high-latitude carbon stars (FHLCs), which presently number in the hundreds, were not easily found until recently.  Most FHLCs are thought to be distant giants, as there is no obvious way for C$_{2}$ to reach the photosphere prior to the red giant phase. While relatively rare, these objects are interesting because, for example, of their utility as a halo velocity tracers \\citep{msg85,bem91}. However, a small number of FHLCs display parallaxes and/or large proper motions, implying main-sequence luminosities (M$_{V} \\sim10$), and have been designated dwarf Carbon (dC) stars. It now appears that a significant fraction of FHLCs are in fact not distant giants, but are nearby dwarfs \\citep[hereafter Paper I]{mar02}.  Based on data obtained during the Sloan Digital Sky Survey \\citep[SDSS;][]{york00} commissioning period, we claimed in Paper I that at least $40-50\\%$ of all FHLCs are dwarfs.  Here we report the results of an expanded sample of SDSS FHLCs.  In Section~2, we discuss the observations and selection criteria for our sample, while Section~3 describes the classification of the objects as dwarfs or giants.  A key problem remains the derivation of simple luminosity discriminants other than proper motion, as the low resolution ($\\sim1000$) spectra of the giants and dwarfs are so similar.  Section~4 discusses photometric and spectroscopic luminosity indicators, while our conclusions are given in Section~5. ",
        "conclusions": "As previously noted, at least $50\\%$ the FHLCs are dwarfs.  Therefore the development of a simple observational luminosity discriminant is imperative if we are to, for example, use the giants as halo velocity tracers.  The fact that the dwarfs and giants differ in intrinsic luminosity by $\\sim10$ mag gives hope that there would be such discriminants.  In fact, several luminosity discriminants have been proposed, although the small number of objects known to date has made it difficult to assess the validity and range of applicability of these discriminants.  With our large sample of objects, an objective analysis of the proposed discriminants, as well as the possibility of uncovering others, is now possible.  \\subsection{Previously proposed photometric luminosity discriminants}  Numerous authors \\citep[e.g.,][]{gma92,joyce98,tiw00,lkr03} have proposed near-IR colors as a luminosity discriminant, while \\citet{mac03} showed a tight sequence for N-type stars for his survey objects in the {\\it (H-K)}, {\\it (J-H)} diagram. While we will discuss the luminosity aspect below, we present in Figure~\\ref{fig5} the 2MASS color-color diagram for all stars ($\\sim100$ stars) in our survey, in the carbon star catalog ($\\vert b \\vert > 10^\\circ$; $\\sim300$ stars) of \\citet{alk01}, and all previously reported FHLC dwarfs (11 stars), summarized by \\citet{lkr03}.  While the N-type stars continue to show the trend seen by \\citet{mac03}, the R-type stars are far more scattered, and the ``F/G Carbon stars'' appear to form a nice continuation of the C stars sequence.  The one N-type star in our survey fits nicely in with those from \\citet{alk01}, while the one R-type star from \\citet{alk01} at {\\it (H-K)}=1.0 is classified as R-N, so is likely an N-type; the R-type star at {\\it (H-K)}=1.6 is the bright Mira variable RU~Vir.  Although near-IR colors have been proposed as a luminosity discriminant, \\citet{wj96} and \\citet{jbh98} argue on purely physical grounds that {\\it JHK} colors alone should not be an unambiguous indicator. Based on 4 objects from Paper~I, we found that 2 objects were consistent with the {\\it JHK} discriminant, while 2 objects were not consistent.  In Figure~\\ref{fig6}, we show an enlargement of the near-IR color-color diagram, with just the SDSS stars, as well as the previously known dwarfs, displayed.  For {\\it (H-K)} colors bluer than about 0.3, there is no separation between giants and dwarfs.  However, for redder colors, the bulk of the objects are dwarfs.  For the (somewhat arbitrary) line shown in the figure, 20 ($74\\%$) of the objects are dwarfs, 1 ($4\\%$) is a giant, and 6 ($22\\%$) are uncertain; if only the confirmed giants and dwarfs are counted, then $95\\%$ of the objects are dwarfs.  A slight shift of the line to the blue (or red) does not significantly change the giant to dwarf ratio. An examination of Figure~\\ref{fig5} shows that most of the previously known dwarfs do indeed appear to be offset from the bulk of the R-type stars from the \\citet{alk01} catalog, but with the addition of the SDSS FHLCs, it now appears that this discriminant can be refined and confirmed.  \\citet{lkr03} proposed that dwarfs and giants could be separated in the {\\it (J-K)}, {\\it (R-J)} diagram.  While we do not have R magnitudes for our SDSS sample, the SDSS {\\it r} band should be a reasonable proxy to the Johnson R, so we have examined this discriminant with our sample.  Figure~\\ref{fig7} shows that, as with the near-IR colors, this discriminant can identify dwarfs, but is somewhat less effective (there are 5 giants in the dwarf region) than using only 2MASS colors.  In Paper~I, we commented that photometric variability may prove to be a simple luminosity criterion, since only giants should show consistent, chaotic variations associated with mass loss; absence of variability, however, does not imply the object is a dwarf.  Many of the SDSS objects have multiple photometric measurements in the imaging database, so we have performed a preliminary check for variability. Given the faintness of many of the objects (particularly in {\\it u}), we define an object as variable if at least 2 of the {\\it g, r}, and {\\it i} magnitudes changed by at least 0.10 mag.  With that criteria, only 1 of the 49 FHLCs with multiple measurements is variable, and this object is listed in Table~\\ref{tbl-2}.  That object has an uncertain luminosity class, and thus is a strong candidate for a giant.  Note, however, that the object could be similar to the dC PG~0824+289 \\citep{hbj93}, which shows variability due to heating effects from a hot companion, and so multiple measurements to look for periodic variability would be interesting.   \\subsection{New photometric luminosity discriminants}  In Figures 1-3, we plot the 3 SDSS color-color diagrams. An examination of those figures shows that, while the ``F/G Carbon stars'' are clearly separated from the R-type stars, there is no indication of a luminosity discriminant. We also investigated other color-color plots, none of which presented any valid discriminant.  The model for a dwarf C star is that the C$_{2}$ was deposited in a previous mass-transfer episode from a companion that is now a faint white dwarf.  In some cases (e.g., PG~0824+289 \\citep{hbj93}; SBS 1517+5017 \\citep{lsl94}), the white dwarf is hot enough to be detectable in the visible.  However, in most cases, the white dwarf is too cool to be visible in the optical.  We examined those dwarfs in our sample that have the bluest {\\it u-g} colors, to look for any spectroscopic evidence of a white dwarf, but none was found. On the other hand, it may be possible to detect the white dwarf in the ultraviolet, where the C star is faint and the contrast may be optimized.  To determine how faint (i.e., cool) the white dwarf would typically have to be to be undetected either spectroscopically or photometrically, we selected one blue and one red giant FHLC from our sample, and added a blackbody and/or real white dwarf spectrum of various temperatures.  We find that a 10,000~K blackbody is barely detectable in the red giant, while an 11,000~K blackbody is barely detectable in the blue giant.  However, in the satellite ultraviolet, these cool white dwarfs could be detectable (e.g., with the ACS SBC camera on {\\it HST}).  Therefore, UV colors may in fact be a useful photometric luminosity discriminant.  \\subsection{Previously proposed spectroscopic luminosity discriminants}  Numerous spectroscopic luminosity discriminants have been previously proposed.  \\citet{gma92} proposed that the appearance of a strong C$_{2}$ band head at $\\lambda6191$ indicated the object was a dwarf, while \\citet{mar02} suggested the feature was both temperature and luminosity dependent.  However, both studies noted that while the absence of $\\lambda6191$ could not be used to firmly classify the object as a giant, the presence of $\\lambda6191$ was a good indicator that the object is a dwarf.  In Figure~\\ref{fig8} we show a histogram of the strength of $\\lambda6191$, which demonstrates that this feature is in fact probably not a good luminosity discriminant.  When $\\lambda6191$ is weak (1.0-1.4), about $50\\%$ of the objects are confirmed dwarfs, while at its strongest ($>1.6$) about a third of the objects are confirmed dwarfs.  Only in the 1.4-1.6 bin do the dwarfs dominate ($70\\%$).  In total, $50\\%$ of the objects with $\\lambda6191$ are confirmed dwarfs, so this feature can no longer be considered a luminosity discriminant.  In Paper~I we suggested that the presence of H$\\alpha$ emission could be an indicator that the object is a giant, since dwarfs are unlikely to process active chromospheres or undergo active mass loss.  However, a contrary case (PG~0824+289) was noted, where heating of the dwarf by a hot DA companion caused the emission.  In that case, the white dwarf was visible in the optical spectrum, so H$\\alpha$ emission in dCs where there is no indication of a hot white dwarf could still be a luminosity discriminant.  Of the 14 objects in our sample with H$\\alpha$ emission, 3 are giants, 3 are dwarfs, and the remaining 8 have are uncertain luminosity.  Thus, the use of H$\\alpha$ emission as a luminosity discriminant is speculative.  The dCs with H$\\alpha$ emission are rare, and thus should be further observed.  The bands of CaH at $\\lambda\\lambda6382,~6750$ are normally strong only in K and M dwarfs \\citep[e.g.,][]{khm91}, which led us to suggest in Paper~I that the presence of these features was a luminosity indicator in FHLCs.  In Figure~\\ref{fig9} we show plot of the equivalent width of $\\lambda6750$ versus that of NaD.  While it is true that there are no giants whose spectra show CaH, there are 10 objects (out of 16 total) that have an uncertain luminosity.  So, while promising, this discriminant must still be consider unproven.  It is interesting to note that sharp turn-on of the CaH feature at W$_{\\lambda}$(NaD)~=~10~\\AA, although there are two objects (one with weaker NaD (hotter) which shows CaH, one with stronger NaD (cooler) that does not) that are exceptions.  To estimate the temperature of this turn-on point, the temperature index \\citep{coh79} of every star that showed NaD was measured, and a fit of NaD as a function of Temperature Index was made.  The Temperature index was converted to a temperature \\citep{yam67}, with the result that the CaH turn-on occurs at T=2900~K.  \\citet{gma92} suggested that dC's have strong $\\lambda6191$ band heads given their relatively weak CN band strength.  In Figure~\\ref{fig10}, we plot the strength of $\\lambda6191$ versus the strength of the CN band at $\\lambda7900$.  Unfortunately, there are so few confirmed giants which show CN that a definitive statement about this criterion is not possible.  However, given the range of $\\lambda6191$ strengths seen for the dwarfs alone, this discriminant can be considered doubtful. On the other hand, the strength of the CN band alone, particularly at the weakest and stronger levels, may be a discriminant (see below).  \\subsection{New spectroscopic luminosity discriminants}  With our large sample of FHLCs, we have searched for possible new spectroscopic discriminants.  \\citet{cgw01} found a correlation between the strengths of the C$_{2}$ bands at $\\lambda5165$ and $\\lambda4737$, which they used for detection of carbon stars.  Our data also shows this correlation, but there is no separation between dwarfs, giants, and the stars of uncertain luminosity.  In a similar vein, we show, in Figure~\\ref{fig11}, the equivalent relation for the strongest C$_{2}$ bands.  Again, there is a correlation, but C$_{2}$ band strength does not appear to be a luminosity discriminant.  As previously noted, the strength of the CN band at $\\lambda7900$ may be a luminosity indicator.  In Figure~\\ref{fig12}, we plot a histogram of the strength of this feature.  As can be seen, all objects with band strengths greater than 1.6 are either giants (6) or of uncertain luminosity (1; the object with the strongest CN).  Similarly, most of the objects with weak CN are dwarfs (19); there are 2 giants and 1 object of uncertain luminosity.  Thus, extremely strong CN seems to imply a giant luminosity, while extremely weak (but present) CN implies a dwarf luminosity."
    },
    "0402/astro-ph0402432_arXiv.txt": {
        "abstract": "{ In this paper we present a three-dimensional numerical model for the radio emission of Magnetic Chemically Peculiar stars, on the hypothesis that energetic electrons emit by the gyrosynchrotron mechanism. For this class of radio stars, characterized by a mainly dipolar magnetic field whose axis is tilted with respect to the rotational axis, the geometry of the magnetosphere and its deformation due to the stellar rotation are determined. The radio emitting region is determined by the physical conditions of the magnetosphere and of the stellar wind. Free-free absorption by the thermal plasma trapped in the inner magnetosphere is also considered. Several free parameters are involved in the model, such as the size of the emitting region, the energy spectrum and the number density of the emitting electrons, and the characteristics of the plasma in the inner magnetosphere. By solving the equation of radiative transfer, along a path parallel to the line of sight, the radio brightness distribution and the total flux density as a function of stellar rotation are computed. As the model is applied to simulate the observed 5\\,GHz lightcurves of \\object{HD\\,37479} and \\object{HD\\,37017}, several possible magnetosphere configurations are found. After simulations at other frequencies, in spite of the large number of parameters involved in the modeling, two solutions in the case of \\object{HD\\,37479} and only one solution in the case of \\object{HD\\,37017} match the observed spectral indices. The results of our simulations agree with the magnetically confined wind-shock model in a rotating magnetosphere. The X-ray emission from the inner magnetosphere is also computed, and found to be consistent with the observations. ",
        "introduction": "Magnetic Chemically Peculiar (MCP) stars are commonly characterized by strong ($B \\sim 10^3$ -- $10^4$~Gauss) and periodically-variable surface magnetic fields. The observed magnetic variability is related to the global magnetic field topology of these stars which is mainly dipolar, with the magnetic axis tilted with respect to the rotational axis. The observed variability is a simple consequence of stellar rotation (Babcock \\cite{Babcock49}).  These stars show evidence of anisotropic winds, as indicated by spectral observations of UV lines (Shore et al. \\cite{shore_et_al}, Shore \\& Brown \\cite{sho_e_bro}). Theoretical studies have demonstrated that a stellar wind in the presence of a dipolar magnetic field can freely flow only from the polar regions, where the magnetic field lines are open (Shore \\cite{shore}). Therefore in MCP stars the wind forms two polar jets, whereas at the latitudes near the magnetic equator the wind is inhibited and the matter is trapped, forming the so called ``dead zone\". The presence of jets and circumstellar matter can explain the emission features observed in the UV spectra and $H_\\alpha$ wings (Walborn \\cite{walborn}). In some cases the mass loss rate ($\\dot{M}=10^{-9}$ -- $10^{-10}~M_{\\sun}\\mathrm{yr}^{-1}$) and the outflow terminal speed ($v_{\\infty}\\simeq 600$ km s$^{-1}$) was derived (Shore et al. \\cite{shore_et_al}, Groote \\& Hunger \\cite{gro_e_hun}).  About 25\\% of MCP stars also show evidence of non-thermal radio continuum emission (Drake et al. \\cite{drake_et_al}, Linsky et al. \\cite{linsky_et_al}, Leone et al. \\cite{leone_et_al}). It has been suggested that the physical processes responsible for the observed radio emission are strictly related to the interaction between wind and magnetic field (Linsky et al. \\cite{linsky_et_al}, Babel \\& Montmerle \\cite{bab_e_mon}, hereafter BM97). The flat spectral index and the observed degree of circular polarization have been interpreted in terms of gyrosynchrotron emission from continuously injected mildly relativistic electrons trapped in the stellar magnetosphere.  The radio emission of MCP stars is variable with the same period as the magnetic field variability, as shown by Leone (\\cite{leone}) and Leone \\& Umana (\\cite{leo_e_uma}) for \\object{HD\\,37479} and \\object{HD\\,37017}, and by Lim et al. (\\cite{lim}) for \\object{HR\\,5624}. In particular, for \\object{HD\\,37479} and \\object{HD\\,37017}, the minimum of the radio emission coincides approximately with the zero of the longitudinal magnetic field, whereas the maxima coincide with the extremes of the magnetic field curves (Leone \\& Umana \\cite{leo_e_uma}). The observed modulation suggests that the radio emission arises from a stable corotating magnetosphere. The temporal variability of the radio flux is probably related to the change of the orientation of the emitting region in the space, due to the misalignment of magnetic and rotational axes.  In this paper we present a three-dimensional gyrosynchrotron model developed with the aim of investigating the nature of the rotational modulation of the radio emission. This kind of study can be used to test the physical scenario proposed to explain the origin of radio emission from MCP stars. ",
        "conclusions": "The study of radio emission from the MCP stars and, in particular, the analysis of the observed modulations offer a unique opportunity to determine the physical properties of the stellar magnetosphere.  The numerical model developed in this paper is used to reproduce the 5 GHz radio lightcurves of two well known MCP stars: \\object{HD\\,37479} and \\object{HD\\,37017}.  The capability of our numerical model to give us an estimation of the physical condition of the magnetosphere is limited by the lack of multifrequency radio lightcurves. However, the spectral information available up to now allows us to focus on a small range of possible solutions, all of them indicating a hard energetic population of non-thermal-emitting electrons, and an inner magnetosphere filled by a thermal plasma consistent with a wind-shock model that provides also X-ray emission.  The possibility to test the radio curves at more than one frequency would allow us to put more stringent constraints on the model.  Our simulations provide useful information that may be summarized as follows: \\begin{itemize} \\item[ ] we confirm the qualitative model proposed by several authors to explain the origin of non-thermal electrons as responsible for the observed radio emission from MCP stars;\\\\ \\item[ ] we point out the importance of the thermal electrons trapped in the inner magnetosphere for reproducing the rotational modulation of the measured radio flux density and the radio spectra of the MCP; this plasma can prroduce X-ray emission;\\\\ \\item[ ] the length of the current sheets, where the electrons are accelerated up to relativistic energies, is about one half of the Alfv\\'en radius; the acceleration process has an efficiency of about $10^{-4}$, as only this small fraction of the electrons in the current sheets is accelerated ; those non thermal electrons have an hard energetic spectrum ($N(\\gamma)\\propto (\\gamma-1)^{-\\delta}$), with $\\delta\\approx 2$; the inner magnetosphere is filled by thermal plasma, whose temperature increases outward from $10^5$ up to $10^6$ K, consistent with a rotating magnetosphere. \\end{itemize}  \\noindent We emphasize the importance of multifrequency radio lightcurves and contemporaneous X-ray observations for a further test our model. This will be a powerful investigation tool for studying the physics of the stellar magnetosphere.   \\begin{acknowledgement} We thank Carlo Nocita for his help in making some figures. We particularly thank the referee Dr. T. Montmerle for his constructive criticism which enabled us to strongly improve this paper. \\end{acknowledgement}   \\appendix"
    },
    "0402/astro-ph0402268_arXiv.txt": {
        "abstract": "We report the results of an XMM-Newton observation of the Narrow-Line Seyfert 1 galaxy Mkn 1239. This optically highly polarized AGN has one of the steepest X-ray spectra found in AGN with \\ax=+3.0 based on ROSAT PSPC data. The XMM-Newton EPIC PN and MOS data confirm this steep X-ray spectrum. The PN data are best-fit by a powerlaw with a partial covering absorption model suggesting two light paths between the continuum source and the observer, one indirect scattered one which is less absorbed and a highly absorbed direct light path. This result agrees with the wavelength dependent degree of polarization in the optical/UV band. Residuals in the X-ray spectra of all three XMM-Newton EPIC detectors around 0.9 keV suggest the presence of an emission line feature, most likely the Ne IX triplet. The detection of NeIX and the non-detection of OVII/OVIII suggest a super-solar Ne/O ratio. ",
        "introduction": "With the launch of the X-ray satellite ROSAT \\citep{tru83} the X-ray energy range down to 0.1 keV became accessible for the first time. During the half year ROSAT All-Sky Survey (RASS, \\citet{vog99}) a large number of sources with steep X-ray spectra were detected (\\citet{tho98, beu99, schwo00}). About one third to one half of these sources are AGN. \\citet{gru96} and \\citet{gru98a, gru03a} found that about 50\\% of bright soft X-ray selected AGN are Narrow-Line Seyfert 1 galaxies (NLS1s, \\citet{oster85, good89}). They turned out to be the class of AGN with the steepest X-ray spectra (e.g. \\citet{bol96, gru96, gru98a, gru01, vau01, gru03a, wil02}). NLS1s are AGN with extreme properties which seem to be linked to each other: An increase in their X-ray spectral index \\ax~correlates with the strength of the optical FeII emission and anti-correlates with the widths of the Broad Line Region (BLR) Balmer lines and the strength of the Narrow-Line Region (NLR) forbidden lines (e.g. \\citet{gru96, gru99, gru03b, laor94, laor97, sul00}). All these relationships are governed by one fundamental underlying parameter, usually called the \\citet{bor92} 'Eigenvector-1' relation in AGN. The most accepted explanation for Eigenvector 1 is the Eddington ratio $L/L_{Edd}$ \\citep{bor02, sul00, gru03b, yuan03} in which NLS1s are AGN with the highest Eddington ratios.   In a spectropolarimetry study of 18 NLS1s \\citet{good89} found three sources to show significant polarization that allowed a detailed spectropolarimetric analysis: Mkn 766, Mkn 1239, and IRAS 1509--211. All these sources show an increase of the degree of polarization towards the blue. In the Broad Line Region (BLR) Balmer lines the degree of polarization is larger then in the continuum, while is is less in the Narrow Line Region (NLR) forbidden lines, suggesting the scattering medium is somewhere located between the BLR and NLR (e.g. \\citet{wills92}). All three sources were observed with the Position Sensitive Proportional Counter (PSPC, \\citet{pfe86}) on board ROSAT (\\citet{rush96a, gru98b, pfe01}). The NLS1 Mkn 1239 ($\\alpha_{2000}$=09h 52m 19.1s; $\\delta_{2000}=-01^\\circ~36^{'}~43^{''}$; z=0.019) had an unusually steep  X-ray spectral index \\ax=2.9 \\citep{bol96} during the ROSAT pointed observation. Here we present the results of a serendipitous observation of Mkn 1239 with XMM-Newton \\citep{jan01}.  The outline of this paper is as follows: in \\S\\,\\ref{observe} we describe the XMM and ROSAT observations and the data reduction, in \\S\\,\\ref{results} we present the results of the ROSAT and XMM data analysis, and in \\S\\,\\ref{discuss} we discuss the results. Throughout the paper spectral indexes are denoted as energy spectral indexes with $F_{\\nu} \\propto \\nu^{-\\alpha}$. Luminosities are calculated assuming a Hubble constant of $H_0$ =75 \\kms Mpc$^{-1}$ and a deceleration parameter of $q_0$ = 0.0. ",
        "conclusions": "We have studied the X-ray spectra of Mkn 1239 in a serendipitous observation by XMM and our three main results are: \\begin{enumerate} \\item The 0.2-2.0 keV soft X-ray spectrum of the source is very steep with \\ax$\\approx$3.0 confirming previous results from ROSAT. \\item The 0.2-12 keV EPIC PN X-ray continuum spectrum is best-fitted by a powerlaw with \\ax=2.92 and two intrinsic absorbers with $N_{\\rm H}=6\\times10^{20}$\\cm and 3$\\times10^{23}$\\cm, supporting the optical spectropolarimetry results by \\citet{good89} of having a direct highly absorbed light path and an indirect scattered light path which is less absorbed. \\item The X-ray spectrum shows an emission line feature around 0.91 keV (observed frame) which is most likely the NeIX triplet. \\item The non-detection of the OVII triplet suggests a super-solar Ne/O ratio. \\item To finally resolve the 0.91 keV feature, longer observations with XMM or even more powerful future X-ray missions are needed. \\end{enumerate}"
    },
    "0402/astro-ph0402574_arXiv.txt": {
        "abstract": "We present a detailed study of a binary system of tailed radio galaxies which, along with 3C75, is the only such binary known to exist. The binary is located in a region of low galaxy density at the periphery of a poor cluster Abell S345, but lies close to the massive Horologium Reticulum supercluster. The radio sources have bent tail morphologies and show considerable meandering and wiggling along the jets, which are collimated throughout their lengths. This work presents observations of the large-scale-structure environment of the binary tailed radio sources with a view to examining the influence of large-scale flows on the morphology and dynamics of the associated radio tails. We argue that the orbital motions of the host galaxies together with tidal accelerations toward the supercluster have resulted in the complex structure seen in these radio tails. ",
        "introduction": "Narrow-angle tailed (NAT) radio sources or head-tail (HT) sources consist of twin jets issuing from the active nuclei in some galaxies and then bending by an angle of almost 90\\degr. These sources are usually found in rich cluster environments, where the jets are understood to have been swept back by the deflecting pressure of the dense intracluster medium (ICM) arising from the relative motion between the host galaxy and that medium. Studies have shown that in order to bend these tails, the ICM must be hot and dense and the relative velocity between the radio galaxy and the ICM must be of order 1000\\kmsec \\citep{1994ApJ...436...67V}. These required velocities and densities make it unlikely for NATs to exist in low density environments.  Wide-angle-tail (WAT) radio sources, on the other hand, have beams which typically bend through angles less than 90\\degr. These sources are, once again, usually found in rich clusters; however, the WATs are usually associated with a dominant galaxy --- often the central galaxy --- which is not expected to be moving at high speeds with respect to the ambient medium. Therefore, the cause for their bending has been unclear: motions of the parent galaxy with respect to the ICM, buoyancy of the relatively light synchrotron plasma in the ICM thermal gas, weather/turbulence in the ICM associated with dynamical evolution and cluster-cluster mergers/interactions may all play roles in particular circumstances.  In this paper we do not attempt to sub-classify the tailed radio galaxies into WAT and NAT sources because the distinction might sometimes be confused by projection effects. Instead, we refer to them simply as bent-tail (BT) sources.  Comparisons between the radio morphologies of BT sources in cluster environments and the distribution of the intra-cluster thermal gas as seen, for example, in X-ray images, indicate that the thermal gas is almost always asymmetric and aligned towards the direction of the bending \\citep{gomez}. It is currently believed that these clusters are undergoing mergers, resulting in large scale flows of hot gas owing to the changing gravitational potential. The bending of the BTs is a consequence of these flows.  BTs in merging clusters are therefore a diagnostic of the ICM weather and, consequently, of the evolving gravitational potential resulting from the merger.  Our attention was drawn to the unusual nature of the radio source J0321--455 when it appeared in the field of the giant radio galaxy PKS~B0319--454 \\citep{1994MNRAS.269...37S} with a four-leaf-clover structure.  Follow-up observations at higher resolution revealed the source to have an extremely unusual configuration: as shown in the overview in Fig.~\\ref{image:CG}, J0321--455 is a close pair of twin-tailed radio galaxies associated with what appears to be a binary galaxy system. We label the northern radio galaxy J0321--455N and the southern source J0321--455S. The other optical objects in the field of Fig.~\\ref{image:CG} are labelled alphabetically and discussed in Section~2.3. The system  is located in a region of low galaxy density, at the periphery of the poor Abell cluster S345 \\citep{1989ApJS...70....1A}, and in the vicinity of the massive Horologium-Reticulum supercluster. The directions toward S345 and the supercluster are indicated in the figure.  The radio tails in J0321--455 are probably a result of, and a diagnostic for, (i) the motions of the host ellipticals about their centre of mass, and (ii) the relative motion of the binary with respect to the intergalactic thermal plasma as a result of the gravitational potentials of the large scale structure. Therefore, we have undertaken a study of the radio structures and the large-scale galaxy environment in this unique situation with the aim of understanding the roles of gas and galaxy dynamics in the triggering and formation of BTs. In this paper we present radio and optical observations of the binary system and its large-scale environment.  Our aim was to image the tailed radio structures and understand the large-scale mass distribution in order to infer the dynamics and kinematics of the system. The determination of dynamical and spectral ages are inputs to understanding how the nuclear activity in both members of this binary system might have been triggered and fuelled. The paper is structured as follows: Section~\\ref{observations} describes the multiwavelength observations which include high resolution radio imaging, spectroscopy of the nearby objects in Fig.~\\ref{image:CG} and the brightest members of S345, and multifibre spectroscopy of a 2\\degr field to map the 3-dimensional (3D) large-scale structure. Section~\\ref{discussion} presents a discussion wherein we have developed constraints on the source formation based on the observational data on the radio structures and the large scale environment. Throughout the paper we adopt cosmological parameters $\\Omega_{\\Lambda}=0.7$, $\\Omega_m=0.3$ and $H_0=71$~km s$^{-1}$Mpc$^{-1}$.  S345 and the binary galaxy system are at a mean redshift of about $z=0.07$ and at this distance, 1$'$ corresponds to a linear size of 71~kpc.  \\begin{figure*} \\begin{center} \\includegraphics[width=10cm]{fig1.eps} \\caption{\\small Binary radio galaxies J0321--455N \\& J0321--455S. 2.4~GHz radio contours are overlaid on a combined R and V band optical image of the field made using the ANU 2.3-m telescope. The contours correspond to $4, 16, 64, 128$ times the rms image noise of $60~\\mu$Jy~beam$^{-1}$. Spectroscopy (see Table~\\ref{2.3mdata}) reveals that the objects labelled c and d and the radio galaxies are at about the same redshift.} \\label{image:CG} \\end{center} \\end{figure*} ",
        "conclusions": "\\label{summary} This work has presented a shining example of the influence of large scale flows on the morphology of radio galaxies. We have presented multiwavelength observations of a pair of tailed radio galaxies located on the periphery of the poor Abell cluster S345, and residing approximately 3~Mpc in projection from the Horologium-Reticulum Supercluster. This system is extraordinary because it is (with the possible exception 3C75) the only known example of a binary system of powerful radio galaxies. Although the obvious triggering mechanism for the pair is their mutual interaction, the lack of other examples in the literature is hard to reconcile with such a statement. We suggest that turbulence and tidal forces due to the nearby cluster and supercluster may be another consideration. Optical spectroscopy of the host galaxies and their neighbours on the sky show that the hosts constitute a binary system. High-resolution radio imaging has shown that the tails of both galaxies are extended East-West along a galaxy filament linking S345 to the supercluster. Multi-object spectroscopy of galaxies within a 1\\degr~radius shows that S345 and other nearby clusters have a 3D spatial distribution consistent with S345 and the binary being part of a filament that is suffering tidal acceleration towards the nearby Horologium-Reticulum supercluster. The existence of X-ray substructure indicates a lack of dynamical equilibrium and is support for turbulent `cluster weather' in the gaseous environment. We propose that the BT source structure has evolved primarily as a result of the combined effects of binary orbital motion and tidal differential expansion in the environment over a period less than $10^8$~years."
    },
    "0402/astro-ph0402024_arXiv.txt": {
        "abstract": "{We report on recent progress in our theoretical understanding of X-ray binaries, which has largely been driven by new observations, and illustrate the interplay between theory and observations considering as examples intermediate-mass X-ray binaries, irradiation-driven evolution, ultraluminous X-ray sources and neutron stars with low-velocity kicks.} \\addkeyword{Stars: neutron} \\addkeyword{Binaries: close} \\addkeyword{Black holes} \\addkeyword{X-rays: stars}   \\begin{document} ",
        "introduction": "\\label{sec:intro} Our general understanding of binaries with compact components, in particular those containing neutron stars and black holes still has serious gaps, and theoretical progress is often driven by new observational discoveries. Observations not only help to guide theorists, but also provide important constraints that a successful theory has to satisfy. In this contribution, we discuss the interplay between theory and observations using several selected topics, including both neutron-star and black-hole binaries. ",
        "conclusions": ""
    },
    "0402/gr-qc0402063_arXiv.txt": {
        "abstract": "Astronomical observations have established that extremely compact, massive objects are common in the universe.  It is generally accepted that these objects are, in all likelihood, black holes.  As observational technology has improved, it has become possible to test this hypothesis in ever greater detail.  In particular, it is or will be possible to measure the properties of orbits deep in the strong field of a black hole candidate (using x-ray timing or future gravitational-wave measurements) and to test whether they have the characteristics of black hole orbits in general relativity.  Past work has shown that, in principle, such measurements can be used to map the spacetime of a massive compact object, testing in particular whether the object's multipolar structure satisfies the rather strict constraints imposed by the black hole hypothesis.  Performing such a test in practice requires that we be able to compare against objects with the ``wrong'' multipole structure.  In this paper, we present tools for constructing the spacetimes of {\\it bumpy black holes}: objects that are {\\it almost}\\ black holes, but that have some multipoles with the wrong value.  In this first analysis, we focus on objects with no angular momentum.  Generalization to bumpy Kerr black holes should be straightforward, albeit labor intensive.  Our construction has two particularly desirable properties.  First, the spacetimes which we present are good deep into the strong field of the object --- we do not use a ``large $r$'' expansion (except to make contact with weak field intuition).  Second, our spacetimes reduce to the exact black hole spacetimes of general relativity in a natural way, by dialing the ``bumpiness'' of the black hole to zero.  We propose that bumpy black holes can be used as the foundation for a null experiment: if black hole candidates are indeed the black holes of general relativity, their bumpiness should be zero.  By comparing the properties of orbits in a bumpy spacetime with those measured by an astrophysical source, observations should be able to test this hypothesis, stringently testing whether they are in fact the black holes of general relativity. ",
        "introduction": "\\subsection{Motivation} \\label{subsec:motive}  Observations have now established that the cores of nearly all nearby galaxies contain a massive, compact, dark object {\\cite{kr95,k03}}. These objects range in mass from several $10^5\\,M_\\odot$ to several $10^9\\,M_\\odot$.  Extremely compact stellar mass objects ($\\sim 10\\,M_\\odot$ or so) exist and have been studied in the galactic field (see, e.g., Ref.\\ {\\cite{bailynetal98}} for a review).  Evidence suggests the existence of objects with intermediate masses, $10^2\\,M_\\odot - 10^4\\,M_\\odot$, filling the gap between the supermassive and stellar mass objects (see, e.g., Ref.\\ {\\cite{mc03}} for a review).  The most generally accepted explanation is that these compact bodies are massive black holes.  Although this is the most generally accepted explanation for these objects, it is not the only explanation.  In some cases, the massive dark objects seen in galaxy cores can be explained quite well as dense clusters of stars or stellar remnants.  Such models are rapidly becoming disfavored in many cases as our ability to study the central regions of galaxies improves --- many of these putative clusters would have to be so compact that they would not be gravitationally stable. By using ``exotic'' matter, it becomes possible to build objects that are massive, compact, but stable.  For examine, by tuning the mass and self interaction of a massive scalar field {\\cite{csw86,r_boson}}, one can build an object that is consistent with much of the observational evidence available today.  Indeed, the fields that describe some of these black hole alternatives are similar to some dark matter candidates, leading to the suggestion that massive compact objects could be dark matter condensates rather than black holes.  Other recently proposed black hole alternatives are motivated by a desire to avoid the information paradox --- the loss of information through the black hole's event horizon.  Such models find ways of eliminating the event horizon altogether, for instance by replacing the event horizon with a hard surface surrounding a ball of negative energy density (the ``gravastar'' model) {\\cite{mm01}}, or by postulating that spacetime itself undergoes a phase transition in the presence of very strong gravitational fields {\\cite{chls}}.  If such objects exist in nature, they should have a deep, strong field structure very different from that of black holes {\\cite{birkhoff}}.  Astronomical measurements are now becoming able to probe into the very strong field of compact objects: optical and infrared observations track stellar orbits at the core of the Milky Way, probing the spacetime of the presumed black hole at Sgr A* {\\cite{ghez,genzel}}; x-ray observations of quasi-periodic oscillations from black hole candidates carry information about gas in the hole's deep strong field {\\cite{psaltis}}; and future gravitational-wave observations may be able to track the sequence of orbits followed by a compact body that slowly spirals into a massive black hole {\\cite{annals,bc04}}.  The question of whether these objects are truly the black holes of general relativity or are described by some alternative model reduces to the question of how one may use measurements of orbital properties to map the spacetime structure (i.e., the gravitational field) of these objects {\\cite{rees04}}.  One thus needs to be able to relate the properties of the measured orbits to the structure of the central gravitating objects.  A powerful way of doing this is by a multipole expansion of the compact object's spacetime.  \\subsection{Multipoles of massive compact objects} \\label{subsec:multipoles}  In Newtonian theory, the gravitational field of a body is simply described by expanding the potential in spherical harmonics.  The potential $\\Phi$ must satisfy \\begin{eqnarray} \\nabla^2\\Phi &=& 4\\pi G\\rho \\qquad \\mbox{(interior)}\\;, \\nonumber\\\\ &=& 0 \\qquad \\mbox{(exterior)}\\;. \\label{eq:poisson} \\end{eqnarray} In the exterior, the potential may be expanded as \\begin{equation} \\Phi = -G\\sum_{lm} \\frac{M_{lm}Y_{lm}}{r^{l+1}}\\;. \\label{eq:exterior} \\end{equation} By matching to an expansion of the interior solution and enforcing Eq.\\ (\\ref{eq:poisson}), we see that the coefficients $M_{lm}$ are {\\it mass multipole moments}: numbers that describe the angular distribution of matter inside the star.  For simplicity, let us focus for a moment on axially symmetric objects, so that only $m = 0$ matters.  Then, for example, $M_{00} = M$, the total mass of the object.  By an appropriate choice of the center of our coordinate system, we put the moment $M_{10} = 0$.  The first interesting moment is $M_{20}$, the quadrupole moment of the object.  This moment has the form ${\\cal Q} M L^2$, where $L$ is the object's ``size'' (e.g., its mean radius), and the dimensionless number ${\\cal Q}$ describes the quadrupolar deformation.  Higher moments can likewise be interpreted as $l$-polar moments of the mass distribution.  Because these moments directly determine the gravitational potential outside of the gravitating object, one can measure properties of its mass distribution by measuring the ``shape'' of the gravitational potential.  Doing so by studying the properties of satellite orbits is the science of {\\it geodesy}.  Somewhat remarkably, such a description describes the exterior spacetimes of bodies in general relativity as well.  For any gravitating body that is stationary, axisymmetric, and reflection symmetric across the equator --- encompassing black holes plus a wide variety of perturbations --- the exterior spacetime is fully specified by a pair of multipole moment families: the mass multipole moment $M_l$, plus the {\\it current} multipole moment $S_l$ {\\cite{g70,h74}}. (Since we have restricted ourselves to axisymmetry, we only consider $m = 0$.  We henceforth drop this subscript.)  The current moment $S_l$ does not appear in Newtonian theory; it reflects the fact that mass and energy flows gravitate in general relativity.  For example, the moment $S_1$ is the magnitude of the spin angular momentum of the body.  If the gravitating body is a Kerr black hole, then the values of the mass and current moments are strongly restricted: in units in which $G = 1$, $c = 1$ (which we use throughout this paper), we must have {\\cite{h74}} \\begin{equation} M_l + i S_l = M(i a)^l\\;. \\label{eq:kerr_multipoles} \\end{equation} This equation tells us that $M_0 = M$, the total mass of the Kerr black hole, and $S_1 = aM$, the magnitude of the black hole's spin angular momentum in these units --- precisely what we already expect. More interestingly, {\\it all} higher moments are completely determined by these two values.  The exterior spacetime of a Kerr black hole is completely determined by its two lowest multipole moments --- its mass and spin.  This is nothing more than a restatement of the ``no hair'' theorem {\\cite{nohair1,nohair2,nohair3,nohair4}}.  By analogy with geodesy, this suggests that one can test the no hair theorem by measuring orbits near black holes.  Using a spacetime that does not necessarily assume the Kerr form of the moments, one could then determine $M_l$ and $S_l$.  If that object is in fact a black hole as described by general relativity, the only free moments are those for $l = 0$ and $l = 1$.  Once they have been determined, {\\it all} higher moments are constrained via Eq.\\ (\\ref{eq:kerr_multipoles}).  Such ``geodesy for black holes'' (which has been given the names ``bothrodesy'' and ``holiodesy'' {\\cite{odyssey}}) would provide a stringent test of the black hole nature of massive compact objects in the universe.  The first detailed analysis of how one might be able to falsify the black hole nature of a massive compact object was by Fintan Ryan {\\cite{r_multi}}.  Ryan showed how to build the spacetime of an object with arbitrary multipole structure, and then analyzed the orbits of small bodies in that spacetime.  (``Small'' means that these bodies do not themselves significantly distort the spacetime, and so can be treated as following approximately geodesic trajectories).  His analysis demonstrated that the accumulated orbital phase was sensitive to these multipoles.  Orbital phase (or some surrogate of this phase) is directly observed by, for example, x-ray timing (today) and gravitational-wave detectors (future).  One could thus imagine using measurements of accumulated orbital phase to test the black hole nature of a massive compact object.  Focusing on gravitational-wave measurements with the planned space-based laser interferometer LISA {\\cite{lisa}}, Ryan showed that enough multipoles should be measurable to easily falsify the black hole hypothesis.  In many cases, enough multipoles would be measurable (up to $l \\sim 5$ or 6) to stringently constrain the object's black hole nature {\\cite{r_measure}}.  Unfortunately, the multipole expansion used by Ryan does not work very well in the deep strong fields of massive black holes, where one expects orbital phases to most stringently test the black hole hypothesis.  Multipole moments essentially label different powers of a $1/r$ expansion.  In the strong field of a black hole (small $r$), such an expansion is not going to be very useful {\\cite{moderate}}. The in-utility of this expansion is reflected by the extremely large number of terms that must be kept to describe a spacetime with arbitrary multipole moments at small radius (cf.\\ Ref.\\ {\\cite{r_multi}}, Sec.\\ III).  \\subsection{Bumpy black holes} \\label{subsec:bumpy}  We advocate a different approach.  The reason for introducing a multipolar expansion is to describe a candidate spacetime differing from that of a black hole.  If one accepts as a starting point the idea that the black hole hypothesis {\\it probably} describes the massive compact objects in question, then one just needs a spacetime to compare against that differs {\\it slightly} from that of a black hole.  Our goal is then to set up a null experiment: we find a trial spacetime that exhibits slight deviations from the spacetime of a black hole.  If the black holes of nature are the black holes of general relativity, we will measure the deviation to be zero.  Past work on candidate objects to test the black hole hypothesis has focused primarily on boson and soliton stars {\\cite{r_boson,r_measure}}.  Though of great intrinsic interest, there is no natural way for a boson star spacetime to smoothly limit to the spacetime of a black hole.  If the massive compact objects we observe in the universe are in fact black holes, then tests based on the boson star model will not provide useful constraints on orbit observations. As a ``straw man'' for the black hole hypothesis, boson stars may unfortunately contain too much straw.  We advocate instead the use of {\\it bumpy black holes}: objects that have a multipolar structure that is very nearly, but not quite, that of a black hole.  As the name suggests, these are black holes with small bumps on them.  If the universe's observed massive compact objects are in fact black holes, then we will find that the amplitude of these bumps is zero (within measurement uncertainty).  A bumpy black hole should be a trial spacetime that behaves well deep into the strong field, and that exhibits a {\\it controllable}\\ deviation from the Kerr solution.  In particular, these trial spacetimes should reduce to normal black holes when the the deviation is set to zero --- bumpy black holes become normal black hole when the bumps are removed. This reduction to normal black holes is a crucial element of using bumpy black holes as a basis for a null experiment.  A key piece of our guiding philosophy is that the notion of multipoles is most useful in the weak field of an object.  One shouldn't be too attached to multipole moments if the goal is an analysis that applies to strong gravitational fields.  Of course, by taking the weak field (large $r$) limit of the spacetimes we construct, bumpy black holes are very usefully described using multipoles.  Indeed, the goal of our detailed calculations will be to assemble a perturbation that is purely quadrupolar when examined in the weak field.  Our construction, however, works very well deep in the strong field, which is crucial for applying these notions to observations.  An important question to ask at this point is {\\it how} one can build a stationary spacetime corresponding to a bumpy black hole.  A key portion of the proof of the no hair theorem demonstrates that any deformation to a black hole will tend to radiate very quickly, removing the bump and pushing us back to the Kerr black hole solution {\\cite{nohair3,nohair4}}.  Some mechanism must exist to maintain the bump.  This is likely to require unphysical matter; the example which we describe in fact requires naked singularities.  One might object that a bumpy black hole spacetime is thus, by construction, unphysical.  Our viewpoint is that the physicality of these spacetimes is {\\it irrelevant}.  Our goal in this analysis is {\\it not} to build a spacetime which might conceivably exist in nature.  Instead, we wish to build a black hole straw man with just the right amount of straw to probe the nature of massive compact objects.  \\subsection{Overview of this paper} \\label{subsec:overview}  The goal of this paper is to present the bumpy black hole concept, to show how bumpy black hole spacetimes are generated, and to demonstrate that the magnitude of the bumps is encoded in the accumulated phase of the hole's orbits.  We focus upon axisymmetric distortions of black holes --- even in axisymmetry, an incorrect moment is enough to falsify the black hole hypothesis for a massive compact object.  We have argued that the language of multipoles is not useful for describing an object's strong field orbits.  To substantiate this argument, we review the multipole description of spacetimes and their orbits in Sec.\\ {\\ref{sec:multipole}}, summarizing the key results of Ryan {\\cite{r_multi}}.  Ryan's formulae and the detailed description of the spacetime in the multipole language show that, as we try to characterize the massive object's strong field, the description becomes extremely complicated.  Although it is possible in principle to use this description to develop tools for mapping spacetimes, it does not appear to be the best approach in practice.  We then begin our detailed presentation of bumpy black holes.  In Sec.\\ {\\ref{sec:basics}}, we show how to build a bumpy black hole spacetime.  The spacetime of a stationary, axisymmetric object is fully described by the Weyl metric {\\cite{w18}}, \\begin{equation} ds^2 = -e^{2\\psi}dt^2 + e^{2\\gamma - 2\\psi}(d\\rho^2 + dz^2) + e^{-2\\psi}\\rho^2 d\\phi^2\\;. \\label{eq:weyl} \\end{equation} Our strategy is to pick an exact solution $\\psi = \\psi_0(\\rho,z)$, $\\gamma = \\gamma_0(\\rho,z)$ for which the line element (\\ref{eq:weyl}) describes a black hole.  For this first analysis, we specialize our background to Schwarzschild black holes; generalization to Kerr black holes should be straightforward in principle (though it may be somewhat involved algebraically).  We then use this exact solution as a background against which to introduce a perturbation, putting $\\psi = \\psi_0 + \\psi_1$ and requiring $\\psi_1/\\psi_0 \\ll 1$; a similar perturbation is defined for $\\gamma$.  The perturbations are constrained by the requirement that they satisfy the vacuum Einstein equations, expanded to first order.  This formulation of the metric is particularly useful because the function $\\psi$ reduces to the Newtonian gravitational potential of the source in the weak field.  We therefore choose our perturbation $\\psi_1$ in such a manner that the weak-field perturbation can be thought of as changing the source's multipoles {\\it as measured in the weak field}.  We then solve the linearized Einstein equations in order to specify the perturbation throughout the exterior spacetime of the bumpy black hole.  Before specifying our perturbations, we first examine the properties of orbits in the bumpy hole's equatorial (reflection symmetry) plane (Sec.\\ {\\ref{sec:orbits}}).  Many useful quantities can be computed in terms of the perturbations $\\psi_1$ and $\\gamma_1$ --- the orbit's energy $E$, angular momentum $L$, and the location of the last stable orbit (a separatrix in orbital phase space dividing dynamically stable and unstable orbits).  We also write down an equation describing the differential advance of the orbit's periapsis.  The periapsis shift arises from a mismatch between the radial and azimuthal orbital frequencies; as such, it can be a sensitive probe of the spacetime. Deviations of this shift from the canonical Schwarzschild value encode the black hole's bumpiness.  We choose particular perturbations in Secs.\\ {\\ref{sec:points}} and {\\ref{sec:ring}}.  A very simple and useful one is that of a point mass near the black hole.  We build a bumpy black hole in Sec.\\ {\\ref{sec:points}} by placing a pair of point masses with mass $\\mu/2$ each near the hole's ``north'' and ``south'' poles.  The same system was used by Suen, Price, and Redmount (SPR) {\\cite{spr88}} to set up an analysis of a black hole with a deformed event horizon.  Our analysis is similar to that of SPR, though we do not focus on the region of spacetime near the horizon.  We build a second type of bumpy black hole in Sec.\\ {\\ref{sec:ring}} by placing a ring of mass $\\mu$ about the hole's equator.  As the analysis of Secs.\\ {\\ref{sec:points}} and {\\ref{sec:ring}} shows, both the polar point mass and the equatorial ring do indeed change the metric's quadrupole moment.  We demonstrate this by calculating weak-field periapsis precession in these spacetimes and showing that the shift contains a term which is exactly what we expect for weak-field quadrupolar distortions (computed in Appendix {\\ref{sec:newton}}).  Unfortunately, these perturbations also change the metric's {\\it monopole}\\ moment (i.e., its mass).  Fortunately, we can build a purely quadrupolar distortion by combining {\\it negative} mass polar points with a positive mass equatorial ring, or vice versa. (Bearing in mind that our goal is to build trial spacetimes for testing the black hole hypothesis, the unphysicality of a negative perturbing mass is not a concern.)  In Sec.\\ {\\ref{subsec:both_weak}}, we show that the weak-field periapsis precession with this combined mass distribution is identical to that of a Schwarzschild black hole plus a quadrupolar deformation.  The points + ring perturbation to a Schwarzschild black hole thus perfectly matches our requirements for a bumpy black hole.  We investigate the strong-field character of this spacetime in Sec.\\ {\\ref{subsec:both_strong}}, showing in particular that the hole's bumpiness is usefully encoded in the strong-field periapsis precession.  Concluding discussion is given in Sec.\\ {\\ref{sec:concl}}.  In particular, we outline further work that should be done to connect the bumpy black hole concept to future astrophysical observations.  Chief among the tasks needed is a generalization to Kerr black hole backgrounds; some steps in this direction are outlined in Appendix {\\ref{sec:kerr}}. ",
        "conclusions": "\\label{sec:concl}  In this paper, we have laid the foundations for a null experiment to test whether a massive compact object is a black hole.  The bumpy black hole spacetimes we construct differ only slightly from normal black hole spacetimes; and, the difference is controlled by a simple adjustable parameter --- the hole's ``bumpiness''.  It should be possible to compare the properties of black hole candidates in nature with these bumpy black hole spacetimes.  If these objects are in fact the black holes of general relativity, measurements will show that the natural spacetimes have a bumpiness of zero.  Quite a bit more work is needed in order to make the bumpy black hole concept useful in practice for astrophysical measurements:  \\begin{itemize}  \\item Foremost is the need to generalize this analysis to bumpy Kerr black holes --- zero angular momentum is a highly unrealistic idealization.  In Appendix {\\ref{sec:kerr}}, we show how, by choosing $\\psi$, $\\gamma$, and using an appropriate coordinate transformation, the Weyl metric (\\ref{eq:weyl_cyl}) encompasses Kerr black holes (in Boyer-Lindquist coordinates).  There should then be no severe difficulty perturbing this metric to build bumpy Kerr black holes, though the details are likely to be complicated.  \\item Probably next in importance is generalizing the orbits which we analyze to include inclination with respect to the hole's equatorial plane.  Besides the astrophysical motivation --- we do not often expect orbits to be confined to a special plane --- the inclusion of an extra degree of orbital freedom offers opportunity.  Motions out of the plane are characterized by oscillations with a frequency $\\Omega_\\theta$ which is generically different from the frequencies $\\Omega_\\phi$ and $\\Omega_r$ discussed in this paper.  These oscillations thus offer an additional ``handle'' by which deviations from the black hole spacetime can be characterized.  As already mentioned in Sec.\\ {\\ref{sec:orbits}}, the equations of motion for inclined orbits of bumpy black holes do not appear to separate (as they do for normal Kerr black holes).  However, the fact that, by definition, bumpy black hole spacetimes are {\\it nearly}\\ the spacetimes of black holes suggests that the equations of motion must {\\it nearly} separate.  In other words, the degree to which the $\\theta$ motion couples to the $r$ and $\\phi$ motion must be small --- no doubt controlled by a coupling factor that is of order $\\mu/M$.  It may be possible to take advantage of this smallness to usefully describe inclined and eccentric bumpy black hole orbits.  \\item We have focused on perturbing mass distributions that produce, in the weak field, a purely quadrupolar spacetime distortion.  We chose to focus on this case because an incorrect quadrupole moment is sufficient to falsify the black hole hypothesis.  There is no reason why we could not go beyond this: by using rings out of the hole's equatorial plane, one could imagine building essentially any multipolar distribution whatsoever.  Indeed, the fact that the equation for the metric function $\\psi$ in Weyl coordinates [Eq.\\ (\\ref{eq:d2psidrho2})] is simply $\\nabla^2\\psi = 0$ tells us that it is simple in principle to specify perturbations whose weak-field multipolar structure is completely arbitrary: the perturbation \\begin{equation} \\psi_1 = \\sum \\frac{{\\cal B}_l Y_{l0}}{(\\rho^2 + z^2)^{(l+1)/2}} \\end{equation} will work perfectly.  The parameter ${\\cal B}_l$ is a generalized $l$-polar ``bumpiness''.  With this ansatz to define our deviations, we can build a bumpy black holes with almost arbitrarily shaped bumps. This will make it possible to strongly constrain the properties of black holes in nature.  As this paper was being completed, an analysis appeared on the gr-qc archive by Ashtekar et al.\\ {\\cite{ih_multipoles}} of the multipole moments of isolated horizons {\\cite{ih_prl}}.  Although we have not investigated this in any depth, it may be beneficial to pursue a connection between the bumpiness of a black hole and the multipole moments expressed in the language of Ref.\\ {\\cite{ih_multipoles}}.  \\end{itemize}  With these generalizations in hand, it should be possible to begin examining detailed mechanisms by which orbital frequencies can be imprinted on astrophysical observables.  For example, one can imagine analyzing accretion disk models to see how a spacetime's bumpiness is imprinted on quasi-periodic oscillations in a source's x-ray spectrum {\\cite{psaltis,sb04,r04}}.  Another example is in gravitational-wave science.  For these ideas to be useful for testing the nature of black hole candidates, we will need to model the gravitational-wave emission and inspiral of compact bodies captured by bumpy black holes.  This problem may not be much more difficult than the corresponding problem for normal Kerr black holes --- because the wave amplitude is itself perturbatively small, the inspiral and wave generation should decouple (at least to first order) from the spacetime's bumps.  We hope to address at least some of these issues in future work."
    },
    "0402/astro-ph0402354_arXiv.txt": {
        "abstract": "The oscillation spectrum of $\\nu$ Eri is the richest known for any variable of the $\\beta$ Cephei type. We interpret the spectrum in terms of normal mode excitation and construct seismic models of the star. The frequency data combined with data on mean colours sets the upper limit on the extent of overshooting from the convective core. We use data on rotational splitting of two dipole ($\\ell=1$) modes (g$_1$ and p$_1$) to infer properties of the internal rotation rate. Adopting a plausible hypothesis of nearly uniform rotation in the envelope and increasing rotation rate in the $\\mu$-gradient zone, we find that the mean rotation rate in this zone is about three times faster than in the envelope. In our standard model only the modes in the middle part of the oscillation spectrum are unstable. To account for excitation of a possible high-order g-mode at $\\nu=0.43\\mbox{ cd}^{-1}$ as well as p-modes at $\\nu > 6\\mbox { cd}^{-1}$ we have to invoke an overabundance of Fe in the driving zone. ",
        "introduction": "}  The rich oscillation spectrum of $\\nu$ Eri obtained in a recent multisite campaign (Handler et al. 2004, Aerts et al. 2004) holds the best prospects for seismic sounding of the interior of a B-star but also presents a considerable challenge to stellar pulsation theory. The sounding is facilitated by the fact that for several excited modes the spherical harmonic indices were determined (De Ridder et al. 2004). The challenge is to explain mode excitation in an unexpectedly broad frequency range.   The oscillation spectrum of $\\nu$ Eri is shown in Fig.\\,1. Before the campaign only the $\\ell=0$ mode and the $(\\ell=1,\\mbox{ g}_1)$ triplet were known. However, for the triplet, the identification of the spherical harmonic was uncertain. Now, thanks to the analysis of accurate multicolour photometry and high-resolution spectroscopy, the uncertainty was eliminated. Reliable information about $\\ell$ is a crucial input for constructing seismic models of a star, that is, the models constrained by the frequency data.  \\begin{figure} \\begin{center} \\includegraphics[width=88mm,viewport=05 30 555 381]{figure1.eps} \\end{center} \\caption{Oscillation spectrum of $\\nu$ Eri. The peaks represented with the broken vertical lines are regarded uncertain. The four modes known before the campaign are between $\\nu=5.6\\mbox{ and 5.8 cd}^{-1}$. De Ridder et al. (2004) identified them as an $\\ell=1$ triplet and an $\\ell=0$ singlet. The newly discovered triplet at $\\nu\\approx6.25\\mbox{ cd}^{-1}$ is identified as $\\ell=1$ (De Ridder et al. 2004). In this paper we will show that the radial mode is the fundamental and the two dipole modes are ${\\rm g}_1$ and ${\\rm p}_1$, respectively. Only these mode frequencies are used in our seismic sounding.} \\end{figure}  In this work, we attempt to make the best use of the frequency data to address unsolved problems in stellar evolution theory concerning element mixing in convectively stable layers and angular momentum evolution. The problems are related because rotation induces fluid flows that may cause mixing. It is important to disentangle effects of overshooting from the convective core and effects of rotationally induced element mixing. The first effect depends only on stellar mass. The second one depends on randomly distributed stellar angular momenta. In this context, $\\nu$ Eridani, which is a very slowly rotating B star, yields a useful extreme case. However, it would be important to try to obtain similar quality data for more rapidly rotating objects.  We are very curious how the angular velocity of rotation behaves in stellar interiors. Let us remind that, while the results of helioseismic sounding of radial structure essentially confirmed the standard solar model, the results for the internal rotation rate came as a surprise. The outstanding question in the case of massive stars is whether a substantial radial gradient of the angular rotation rate may be present in their interiors. The question is related to another very interesting and important problem which is the origin of magnetic fields in B type stars.  A seismic model of a pulsating star should not only account for the measured frequencies but also for the instability of the detected modes. The latter requirement sets a different kind of constraint than the former as it concerns only properties of the outer layers, where most of contribution to mode driving and damping arises. Since the frequency spectrum of $\\nu$ Eri seems unusually broad for an object of its type, constraints in this case are particularly interesting.  The structure of this paper is as follows. In Sect. 2 we construct models of the internal structure for $\\nu$ Eri and calculate its oscillation frequencies. We find models for which the oscillation frequencies of three $m=0$ modes -- $(\\ell=0\\mbox{, p}_1)$, $(\\ell=1\\mbox{, g}_1)$ and $(\\ell=1\\mbox{, p}_1)$ -- reproduce the observations as well as have effective temperatures and luminosities within the observational error box. In Sect. 3 we use these models and data on the triplet structure to infer the internal rotation of the star.  The problem of how to make the observed modes pulsationally unstable is discussed in Sect. 4. The tasks of constructing the model of $\\nu$ Eri whose mode frequencies match the data and of the identification of the driving effect are not quite independent. The driving effect in B stars strongly depends on the iron content and its distribution in outer layers. In Sect. 5, we study how modification of the iron distribution affects the frequencies. We also propose plausible identifications of all modes detected in the $\\nu$ Eri oscillation spectrum.  \\vspace{4mm} ",
        "conclusions": "We believe that our seismic models of $\\nu$ Eridani yield a good approximation to its internal structure and rotation but there is room for improvement and a need for a full explanation of mode driving. We did not fully succeed in the interpretation of the oscillation spectrum of $\\nu$ Eridani with our standard evolutionary models and our linear nonadiabatic treatment of stellar oscillations. We fail to reproduce mode excitation in the broad frequency range, as observed, and to reproduce the frequency of the ($\\ell=1,p_2$) mode, associated with the highest frequency peak in the spectrum. We showed that both problems may be cured by allowing an enhancement of the iron abundance in the zone of the iron-opacity bump. The enhancement may result from effects of radiative levitation, as first proposed by Charpinet et al. (1996, 1997) to explain oscillations in sdB stars. Our proposal is based on models with an {\\it ad hoc} iron enhancement in the bump zone and must be checked with use of stellar evolution codes that take into account effects of levitation and diffusion of chemical elements. We regard this as the most timely theoretical work aimed at the interpretation of $\\nu$ Eridani.  The proposed modification concerns only the outer layers and has a small effect on models of the stellar interior and the frequencies of g and p$_1$ modes. Therefore we believe that our seismic models describe the deep internal structure of the star. These models reproduce frequencies of the fundamental radial modes and two $\\ell=1$ modes and have effective temperatures and luminosities within the measurement error box. Satisfactory models exist in a range of the overshooting parameter, $\\alpha_{\\rm ov}=(0 - 0.12)$, which corresponds to the mass range $(9.9 - 9.0) M_\\odot$, age range $(16 - 20)$ My, and a fractional hydrogen depletion in the core range $(0.34 - 0.38)$. Here the important finding is the upper limit for the extent of overshooting. There is a prospect for getting a more stringent constraint on $\\alpha_{\\rm ov}$ with a more accurate determination of the star's effective temperature.  We believe that our most important finding concerns internal rotation. We presented evidence for a sharp decline of the rotation rate, $\\Omega$, through the $\\mu$-gradient zone, around the shrinking convective zone. With data on the splitting for only two $\\ell=1$ modes, our estimates had to rest on the assumed form of the $\\Omega(r)$ dependence. Since a significant gradient of $\\Omega$ is possible only in the $\\mu$-gradient zone, we adopted a constant $\\Omega$ above it (rotation within the convective core has a negligible influence on the splitting). With this simplification we found that data require a decline by a factor $\\sim 5$ within the zone and an equatorial velocity of about 6 km/s. The second order effect from this slow rotation in the envelope cannot account for the asymmetry of the $(\\ell=1\\mbox{, g}_1)$ triplet as determined from the pre-campaign data. Only with new observations we will be able to tell whether the problem is real. Our star is not the first $\\beta$ Cephei star for which evidence for nonuniform rotation was put forward. However, we believe that in the $\\nu$ Eri case the evidence is significantly stronger than in the case of HD 129929 (Aerts et al., 2003), because we relied on splitting data for modes having very different probing kernels.  $\\nu$ Eridani proved to be a very important pulsating star. After the Sun, it so far is perhaps the most rewarding main sequence object for asteroseismology. Many more modes were detected in several $\\delta$ Scuti stars but in none of them so many modes have unambiguously been identified. This star deserves continued observational efforts. An extended data base is needed for precise determination of the frequencies of the modes in the $\\ell=1$ triplets. This is important for a more credible assessment of differential rotation. In this context it would be important to obtain a precise value of the projected equatorial velocity of rotation from spectroscopy. The best $v\\sin i$ value currently available (Aerts et al. 2004) is higher than that of equatorial velocity inferred in this work which may however be in part due to the combined effects of pulsational and thermal line broadening."
    },
    "0402/astro-ph0402162_arXiv.txt": {
        "abstract": "We consider the statistics of pulsar binaries with white dwarf companions (\\nswd). Using the statistical analysis method developed by Kim et al.\\ (2003) we calculate the Galactic coalescence rate of \\nswd~binaries due to gravitational-wave emission. We find that the most likely values for the total Galactic coalescence rate (\\rtot) of \\nswd~binaries lie in the range 0.2--10 Myr$^{-1}$ depending on different assumed pulsar population models. For our reference model, we obtain ${\\cal R}_{\\rm tot}=4.11_{-2.56}^{+5.25}$ Myr$^{-1}$ at a 68\\% statistical confidence level. These rate estimates are not corrected for pulsar beaming and as such they are found to be about a factor of 20 smaller than the Galactic coalescence rate estimates for double neutron star systems. Based on our rate estimates, we calculate the gravitational-wave background due to coalescing \\nswd~binaries out to extragalactic distances within the frequency band of the Laser Interferometer Space Antenna. We find the contribution from \\nswd~binaries to the gravitational-wave background to be negligible. ",
        "introduction": "\\label{sec:intro}  The observed properties of double neutron star (DNS) systems along with models of pulsar survey selection effects have been used for many years in order to estimate the coalescence rate of DNS systems due to the emission of gravitational radiation {\\citep{nps91,ph91,cl95}}. In Kim, Kalogera, \\& Lorimer (2003; hereafter paper I\\nocite{kkl03}), we presented a novel method to calculate the probability distribution of the coalescence rate estimates for pulsar binaries (see also \\citet{k04}). In paper I, we applied this method to Galactic DNS systems. Having a probability distribution at hand allowed us to calculate the most likely value for the Galactic DNS coalescence rate as well as statistical confidence limits associated with it. These were then used to calculate the expected DNS inspiral detection rate for Laser Interferometer Gravitational-Wave Observatory {\\citep{ligo92}}. In this second paper, we extend our study to \\nswd~binaries that are relevant to the future NASA/ESA mission Laser Interferometer Space Antenna (LISA; Bender et al.\\ 1998{\\nocite{lisa98}}) .  There are now more than 40 neutron star-white dwarf (\\nswd) binary systems known in the Galactic disk (see e.g.~{\\citet{d01}} for a review). Here, we consider the subset of \\nswd~binaries which will coalesce due to gravitational-wave (GW) emission within a Hubble time. There are currently three such {\\it coalescing} binaries known: PSR~J0751+1807 {\\citep{lzc95}}, PSR~J1757$-$5322 {\\citep{eb01}}, and PSR~J1141$-$6545 {\\citep{k00,b03}}. We calculate the Galactic coalescence rate of \\nswd~binaries based on their observed properties using the method introduced in paper I. The GW frequencies emitted by these systems fall within the $\\sim0.1-100$ mHz frequency band of the LISA{\\nocite{lisa98}}. Using our rate estimates we calculate the GW amplitude due to the \\nswd~binaries out to cosmological distances and compare it to the sensitivity curve of LISA {\\citep{lisa00}} as well as the Galactic confusion noise estimates from white dwarf binaries {\\citep{bh90,bh97,nyp01,s01}}.  The organization of the rest of this paper is as follows.  In \\S \\ref{sec:coal}, we consider the lifetimes of \\nswd~binaries and summarize the techniques we use to calculate their coalescence rate. The results of these calculations are presented in \\S \\ref{sec:results}. In \\S \\ref{sec:lisa} we use our rate results to calculate the expected GW backgound produced by \\nswd~binaries Finally, in \\S \\ref{sec:discussion}, we discuss the implications of our results. ",
        "conclusions": "\\label{sec:discussion}  We have used detailed Monte Carlo simulations to calculate the Galactic coalescence rate of \\nswd~binaries. From the reference model, the most probable value of ${\\cal R}_{\\rm tot}$ is estimated to be $4.11_{-2.56}^{+5.25}$ Myr$^{-1}$ at a 68\\% statistical confidence limit. We find that the coalescence rate of \\nswd~binaries is about factor of 20 smaller than those of DNS for all pulsar population models we consider. As mentioned above, we did not take into account any beaming correction for \\nswd~binaries. If we assume a beaming fraction of pulsars in \\nswd~binaries similar to that of pulsars found in DNS, $f_{\\rm b}\\sim6$, then the discrepancy beween \\rpeak~(DNS) and \\rpeak~(\\nswd) is signicantly reduced. As a simple estimate, if we assume $f_{\\rm b,1141}\\sim5$, but keeping $f_{\\rm b}=1$ for the other two binaries, the estimated Galactic coalescence rate increases to $18.06_{-12.74}^{+26.05}$ Myr$^{-1}$ at a 68\\% confidence limit. Hence the ratio between DNS and \\nswd~coalescence rate decreases to about 5. Because the contribution from PSRs J0751+1807 and J1757--5322 is an order of magnitude smaller than that of J1141--6545, moderate values of beaming fraction for those recycled pulsars do not change the result significantly.  Based on the number of sources of \\nswd~binaries in our Galaxy, we estimate the effective GW amplitude from the cosmic population of these systems. We find that the GW background from \\nswd~binaries is too weak to be detected by LISA for the nominal beaming correction. Only by adopting an unreasonably large beaming correction factor, $f_{\\rm b}>10$, could these systems be detectable by LISA in the mHz range. These results are in good agreement with an independent study by Cooray (2004) based on statistics of low mass X-ray binaries.  We finally note that combining the results from paper I and this work can give us strong constraints on the population synthesis models. The preferred models, which are consistent with both ${\\cal R_{\\rm NS-WD}}$ and ${\\cal R_{\\rm DNS}}$, can then be used for the estimation of the coalescence rate of neutron star -- black hole binaries, which have not yet been observed."
    },
    "0402/astro-ph0402481_arXiv.txt": {
        "abstract": "We present three-dimensional hydrodynamic simulations of the evolution of selfgravitating, thick accretion discs around hyperaccreting stellar-mass black holes. The black hole-torus systems are considered to be remnants of compact object mergers, in which case the disc is not fed by an external mass reservoir and the accretion is non-stationary. Our models take into account viscous dissipation, described by an $\\alpha$-law, a detailed equation of state for the disc gas, and an approximate treatment of general relativistic effects on the disc structure by using a pseudo-Newtonian potential for the black hole including its possible rotation and spin-up during accretion. Magnetic fields are ignored. The neutrino emission of the hot disc is treated by a neutrino-trapping scheme, and the $\\nu\\bar\\nu$-annihilation near the disc is evaluated in a post-processing step. Our simulations show that the neutrino emission and energy deposition by $\\nu\\bar\\nu$-annihilation increase sensitively with the disc mass, with the black hole spin in case of a disc in corotation, and in particular with the $\\alpha$-viscosity. We find that for sufficiently large $\\alpha$-viscosity $\\nu\\bar\\nu$-annihilation can be a viable energy source for gamma-ray bursts. ",
        "introduction": "In a series of papers (Ruffert, Janka \\& Sch\\\"afer 1996; Ruffert et al. 1997; Ruffert \\& Janka 1999, 2001) it was shown that the neutrino emission associated with the dynamical phase of the merging or collision of two neutron stars (NS+NS) is powerful, but too short to provide the energy for gamma-ray bursts (GRBs) by neutrino-antineutrino annihilation. Significant heating of the coalescing stars occurs only after they have plunged into each other. The neutrino luminosities can then rise to several $10^{53}\\,{\\rm erg\\,s}^{-1}$ (see also Rosswog \\& Liebend\\\"orfer 2003) and even exceed $10^{54}\\,{\\rm erg\\,s}^{-1}$ in case of the more violent (although probably very rare) collisions.  After a few milliseconds, the compact massive remnant of the merger will most likely collapse to a black hole (BH) (see, e.g., Popham, Woosley \\& Fryer 1999, Di Matteo, Perna \\& Narayan 2002, Oechslin et al. 2004, Shibata \\& Ury\\={u} 2000). When this happens some matter can remain in a toroidal accretion disc around the BH and a funnel with low baryon density develops along the system axis. Hydrodynamic simulations indicate that between some 0.01$\\,M_{\\odot}$ and a few 0.1$\\,M_{\\odot}$ of matter may have enough angular momentum to resist immediate absorption into the BH (see, e.g., Ruffert \\& Janka 1999, Shibata \\& Ury\\={u} 2000). A similar result was obtained for BH+NS mergers (Janka et al.~1999, see also Lee 2001 and references therein). This matter is swallowed by the BH on the time scale of viscous transport of angular momentum, which is much longer than the dynamical time scale. Different from the collapsar case there is no external reservoir of stellar matter which feeds the accretion torus. Therefore the relevant time scale is set by viscosity-driven accretion rather than infall. The accretion is only approximately stationary, if global instabilities do not play a role (Ruffert et al.\\ 1997, Lee \\& Ramirez-Ruiz 2002). Post-merger BH accretion might explain short GRBs, but is unlikely to account for long GRBs (e.g., Narayan, Piran \\& Kumar 2001).  Post-merger accretion tori have maximum densities between some $10^{10}\\,$g$\\,$cm$^{-3}$ and more than $10^{12}\\,$g$\\,$cm$^{-3}$ and extend to outer radii of 10--20 times the Schwarzschild radius ($R_{\\mathrm{s}} = 2GM/c^2$). The accretion proceeds with rates of fractions of a solar mass per second up to several solar masses per second. In case of such ``hyperaccreting'' black hole systems (Popham et al.\\ 1999) photons are strongly coupled to the torus plasma and therefore are inefficient in transporting away energy. Neutrinos, however, are abundantly created by weak interactions in the very dense and hot tori and a fair fraction of the gravitational binding energy of the accreted matter can be radiated away by them (``neutrino-dominated accretion flow'', NDAF, Popham et al.\\ 1999, Ruffert et al.~1997, Ruffert \\& Janka 1999, Narayan et al.\\ 2001, Kohri \\& Mineshige 2002, Di Matteo et al.\\ 2002).  In this Letter we present, for the first time, three-dimensional (3D) simulations of time-dependent hyperaccretion from thick tori on a stellar-mass BH including the physical effects of on a stellar-mass BH including the physical effects of BH rotation (Artemova, Bj\\\"ornsson \\& Novikov 1996), viscosity ($\\alpha$-prescription following Shakura \\& Sunyaev 1973), a realistic finite-temperature equation of state according to Lattimer \\& Swesty~\\cite{lat91}, and energy loss and change of lepton number by neutrino emission. The annihilation of emitted neutrinos and antineutrinos is investigated for its potential to provide the energy of relativistic gamma-ray burst fireballs (see, e.g., Eichler et al.\\ 1989; Narayan, Paczy\\'nski \\& Piran 1992; Woosley 1993). Our 3D modeling of the time-dependent (non-stationary) torus evolution abolishes approximations of previous (semi-)analytic work (e.g., Popham et al.\\ 1999, Di Matteo et al.\\ 2002) or more radically simplified numerical modeling in 2D (Lee \\& Ramirez-Ruiz 2002). ",
        "conclusions": "We have performed three-dimensional hydrodynamic simulations of hyperaccreting stellar mass BHs with self-gravitating, thick accretion tori, varying the torus mass, BH spin and $\\alpha$-viscosity of the torus gas. Our models show that tori which are heated by viscous dissipation of energy get inflated by thermal pressure and are therefore unlikely to become optically thick to neutrinos. The neutrino luminosities stay well below the Eddington limit $L_{\\nu, {\\rm Edd}} = 4\\pi G M_{\\rm BH}c/\\kappa_{\\nu}\\sim 2\\times 10^{55}\\,(M_{\\rm BH}/4\\, M_{\\odot})\\,$erg$\\,$s$^{-1}$. Here $\\kappa_{\\nu} \\sim 10^{-17}\\ave{\\epsilon_{\\nu}^2}/(20\\,{\\rm MeV})^2$ is the mean neutrino opacity (e.g., Janka 2001). The latter is mostly determined by the contribution of $\\bar\\nu_e$ which dominate the energy loss of the torus. These findings are in contrast to conclusions drawn from more approximative modeling approaches which treated the vertical structure of the torus by height-averaging (Di Matteo et al.\\ 2002).  Our results also suggest that the integral rate of the energy deposition by $\\nu\\bar\\nu$-annihilation drops like $t^{-3/2}$ during the long-time, slow decay of the accretion rate that follows an initial, transient relaxation phase of 10--20$\\,$ms with very high mass accretion rates and neutrino emission. This temporal decrease is less steep than obtained in previous simulations with azimuthal symmetry which used a simple ideal gas equation state and made the assumption that all the dissipated energy is radiated away in neutrinos (Lee \\& Ramirez-Ruiz 2002).  We have shown that $\\nu\\bar\\nu$-annihilation in the low-density funnel above the poles of the BH can deposit energy at a rate of $\\sim 10^{50}\\,$erg$\\,$s$^{-1}$, accounting for a total energy release of some $10^{49}\\,$erg in case of sufficiently large torus mass ($\\ga\\,$0.1$\\,M_{\\odot}$), high disk viscosity ($\\alpha\\sim 0.1$) and BH rotation with spin parameter $a\\sim 0.6$. These conditions are expected to be generically produced by BH+NS mergers and most likely also by NS+NS mergers (Janka et al.\\ 1999). The main effect of direct BH rotation is an increase of the lifetime of the torus. The energy release by $\\nu\\bar\\nu$-annihilation satisfies the energy requirements of short GRBs in case of a moderate amount of collimation of the ultrarelativistic bipolar outflows into a fraction $f = 2\\Delta\\Omega/4\\pi$ of a few per cent of the sky (e.g., Rosswog \\& Ramirez-Ruiz 2003). Our results for the post-merging accretion of a remnant BH-torus system are therefore more optimistic than the estimates based on NS+NS merger simulations by Rosswog \\& Ramirez-Ruiz (2002). Future simulations, however, will have to show whether the deposited thermal energy can be efficiently converted to axial outflows from BH-torus systems that are sufficiently luminous and collimated to account for short GRBs (Aloy, Janka \\& M\\\"uller, in preparation)."
    },
    "0402/astro-ph0402448_arXiv.txt": {
        "abstract": "We have performed high resolution 3D simulations with adaptive mesh refinement, following the ISM evolution in a star forming galaxy both on small (<1 pc) and large (>10 kpc) scales, enabling us to track structures in cooling shock compressed regions as well as the entire Galactic fountain flow. It is shown in an MHD run that the latter one is not inhibited by a large scale disk parallel magnetic field. The fountain plays a vital r\\^ole in limiting the volume filling factor of the hot gas. Contrary to classical models most of the gas between 100K and 8000 K is found to be thermally unstable. On scales of superbubbles we find that the internal temperature structure is rather inhomogeneous for an old object like our Local Bubble, leading to low O{\\sc vi} column densities, consistent with observations. ",
        "introduction": "In their seminal paper of a three-phase model regulated by supernova explosions in an inhomogeneous medium, \\cite{mo1977} predicted a volume filling factor of the hot intercloud medium (HIM) of {\\bf $f_{v, hot} \\simeq 0.7 - 0.8$}.  However, observations point to a value of $\\sim 0.5$ (e.g., \\cite{de1992}) or even lower when external galaxies are taken into account (e.g., \\cite{bb1986}).  A way out has been suggested by \\cite{ni1989} by the so-called chimney model, in which hot gas can escape into the halo through tunnels excavated by clustered supernova (SN) explosions. Indeed X-ray observations of several nearby edge-on galaxies have revealed extended, galaxy-sized halos (e.g., \\cite{w2001}). The transport of gas into the halo is, however, still controversial, since break-out may be inhibited by a large-scale disk parallel magnetic field (see e.g., \\cite{mss1993}). However, \\cite{to1998} has performed 3D MHD simulations of expanding superbubbles, including radiative cooling, and finds that bubble confinement only occurs when the energy injection rate is below a critical value of $\\sim 10^{37} \\, {\\rm erg} \\, {\\rm s}^{-1}$ (see also \\cite{mm1988}) and/or the field scale height is infinite, which is unrealistic.  Attempts to determine the occupation fraction of the different phases, and in particular of the hot gas, by means of modelling the effects of SNe and SBs in the ISM, have been carried out by several authors (e.g., \\cite{fe1995, fe1998, ko1999}). However, these models do not include the circulation of gas between the disk and the \\emph{full} halo, thus being unable to resolve the high-$z$ region; neither do they take into account the mixing between the different phases. Therefore, an estimate of the volume filling factors may be misleading. Using the 3D supernova-driven ISM model of \\cite{av2000} incorporating magnetic fields and the adaptive mesh refinement technique in HD \\cite{av1998} and MHD (using a modified version of \\cite{balsara2001}) algorithms coupled to a 3D parallel (multi-block structured) scheme, we explored the effects of the establishment of the disk-halo-disk circulation and its importance for the evolution of the ISM in disk galaxies both with and without magnetic fields. In this paper we review some of the results from these simulations (\\S3), and compare in the case of the Local Bubble derived O{\\sc vi} column densities with observations (\\S4), followed by a discussion of the dynamical picture that emerges from these simulations. Other important issues like the volume filling factors of the ISM ''phases'', the dynamics of the galactic fountain, the conditions for dynamical equilibrium and the importance of convergence of these results with increasing grid resolution, the variability of the magnetic field with density, the importance of ram pressure in the ISM, and the amount of gas in the unstable regimes have been treated elsewhere \\cite{av2000, ab2004a, ab2004b, ab2004c, ab2004d}. ",
        "conclusions": "The dynamical picture that emerges from these simulations is that the evolution of the ISM in disk galaxies is intimately related to the vertical structure of the thick gas disk and to the energy input per unit area by supernovae. The system evolves towards a dynamical equilibrium state on the global scale if the boundary conditions vary only in a secular fashion. Such an equilibrium is determined by the input of energy into the ISM by SNe, diffuse heating, the energy lost by radiative and adiabatic cooling and magnetic compression, and is only possible \\emph{after} the full establishment of the Galactic fountain, which for the Milky Way takes about 300 Myr (\\cite{ab2004a}, see also \\cite{kahn81}). It should be emphasized, since disk and halo are dynamically coupled not only by the escape of hot gas, but also by the fountain \\emph{return} flow striking the disk, that the disk equilibrium will also suffer secular variations (see also \\cite{rb1995}).  Furthermore, the ISM in the disk is dominated by thermal pressure gradients mostly in the neighborhood of SNe, which drive motions whose ram pressures are dominant over the mean thermal pressure (away from the energy sources) and the magnetic pressure. The magnetic field is only dynamically important at low temperatures, but can also weaken gas compression in MHD shocks and hence lower the energy dissipation rate. The thermal pressure of the freshly shock heated gas exceeds the magnetic pressure by usually more than an order of magnitude and the B-field can therefore not prevent the flow from rising perpendicular to the galactic plane.  Thus, the hot gas is fed into the galactic fountain at almost a similar rate than without field.  On the scales of superbubbles, it is found that their expansion into a highly turbulent and inhomogeneous medium leads to considerable deviations from the classical model by developing internal temperature and density structure for older bubbles. Thus the O{\\sc vi} column densities we find there are fairly low -- and in agreement with observations -- while to our knowledge other Local Bubble models so far have failed this test."
    },
    "0402/astro-ph0402176_arXiv.txt": {
        "abstract": "We review the observational properties of the class of young neutron stars known as ``anomalous X-ray pulsars,'' emphasizing the tremendous progress that has been made in recent years, and explain why these objects, like the ``soft gamma repeaters,'' are today thought to be young, isolated, ultrahigh magnetic field neutron stars, or ``magnetars.'' \\vspace{1pc} ",
        "introduction": "Prior to the commissioning of {\\it BeppoSAX} and the {\\it Rossi X-ray Timing Observatory} in 1996, the so-called ``Anomalous'' X-ray Pulsars (AXPs) were considered very mysterious sources, because the energy source for their bright X-ray emission was unknown.  At the time, there were only 3 known members of this class.  They were distinguished by having periods in the narrow range 6--9~s, showing approximately steady spin-down, and having softer spectra in general that those seen in accreting X-ray pulsars.  All were known to lie within 1$^{\\circ}$ of the Galactic Plane, and interestingly, one source, 1E~2259+586, was known to reside in the supernova remnant CTB~109.  AXPs as a class were identified as having modest X-ray luminosities, in the range $L_x \\sim 10^{34}-10^{35}$~erg~s$^{-1}$.  The leading model to explain the AXPs was that they were accreting neutron stars, though with properties very different from the bulk of established accreting X-ray pulsars, including the absence of any evidence of a companion \\citep{vtv95,ms95}.  The situation post-{\\it BeppoSAX} and especially in the latter years of {\\it RXTE} is very different and much clearer.  The basic phenomenology of the sources is now well mapped out.  Here, we systematically review the most important properties of this class of objects, which now includes 5 and possibly 8 sources (see Tables~\\ref{ta:axps1} and \\ref{ta:axps2}), and summarize why today, accretion models are strongly disfavored. Rather, the magnetar model, in which AXPs are isolated young neutron stars powered by a decaying ultrahigh magnetic field, provides the most compelling explanation for the unusual AXP source properties, as it does for an equally as exotic class, the soft gamma repeaters (SGRs).  AXPs have also been reviewed recently by \\citet{mcis02}, and magnetars in general have been reviewed recently by \\citet{kas03a,kas03b}.    \\begin{sidewaystable} \\caption{Spin parameters for AXPs.} \\label{ta:axps1} \\newcommand{\\m}{\\hphantom{$-$}} \\newcommand{\\cc}[1]{\\multicolumn{1}{c}{#1}} \\renewcommand{\\arraystretch}{1.2} % \\begin{tabular*}{\\textheight}{@{\\extracolsep{\\fill}}lcccccccc} \\hline Source & \\cc{Distance$^{\\dagger}$}  & \\cc{SNR} &  \\cc{$P$}  & \\cc{$\\dot{P}$} & \\cc{$B_{dp}$}  & \\cc{$\\dot{E}_s$}  & \\cc{$\\tau_c$}  & \\cc{Ref.}\\\\ &  (kpc)                     &          &   (s)      & ($\\times 10^{-11}$) & ($\\times 10^{14}$~G) & ($\\times 10^{32}$~erg s$^{-1}$) & (kyr) & \\\\ \\hline \\oft     &  $\\gapp 1.0$ or $\\gapp2.7$ & $-$     &  8.69 & 0.196 & 1.3 & 1.2 & 7.0   &  1 \\\\ \\tfe     & $\\gapp2.7$  & $-$     &  6.45 & $\\sim 3.81$     & $\\sim 5.0$  & $\\sim 55$ & $\\sim 2.7$ & 2 \\\\ \\soe     & $\\sim8$ & $-$     & 11.00 & 1.86 & 4.6 & 5.4 & 9.4   & 3\\\\ \\efo     &  5.7-8.5 & Kes 73 & 11.77 & 4.16 & 7.1 & 9.9 & 4.5   & 4  \\\\ \\tfn     &  3 & CTB 109 &  6.98 & 0.0483   & 0.59 & 0.55& 230 & 5 \\\\ \\axj$^*$ &  $\\sim8$ & Kes 75  &  6.97 & $-$ & $-$ & $-$ & $-$   &  6  \\\\ \\cxo$^*$ & 57 & $-$     &  8.02 & $-$ & $-$ & $-$ & $-$ &  7  \\\\ \\ett$^*$ &  $\\sim10$ & $-$     &  5.54 & 1.15 & 2.6 & 26 & 7.6   &  8   \\\\ \\hline \\end{tabular*}\\\\[2pt] ($*$) not confirmed; ($\\dagger$) see   \\\"Ozel, Psaltis \\& Kaspi 2001 for a discussion on distance estimates for the confirmed AXPs; References: (1) Gavriil \\& Kaspi 2002; (2) Kaspi et al.~2001; (3) Kaspi \\& Gavriil 2003; (4) Gotthelf et al.~2002; (5) Woods et al.~2003; (6) Torii et al.~1998; (7) Lamb et al.~2003; (8) Ibrahim et al.~2003. \\end{sidewaystable}  \\begin{sidewaystable} \\caption{Spectral parameters for AXPs.} \\label{ta:axps2} \\newcommand{\\m}{\\hphantom{$-$}} \\newcommand{\\cc}[1]{\\multicolumn{1}{c}{#1}} \\renewcommand{\\arraystretch}{1.2} % \\begin{tabular*}{\\textheight}{@{\\extracolsep{\\fill}}lcccccc} \\hline Source & \\cc{$n_H$}   & \\cc{$\\Gamma$}    & \\cc{$kT$}  & \\cc{$L_x$} & \\cc{$f_{\\mathrm{pl}}$ ($\\%$)$^{\\dagger}$} & \\cc{Ref.} \\\\ & ($\\times 10^{22}$\\ cm$^{-2}$) &   &  (keV) &  (erg s$^{-1}$) & & \\\\ \\hline \\oft      & 0.88 & 3.3  & 0.42  & $3.3\\times 10^{34}$  & $\\sim88$  & 1\\\\ \\tfe      & 1.0  & 2.9  & 0.63  & $3.4\\times 10^{34}$  & $\\sim80$  & 2\\\\ \\soe      & 1.49 & 3.11 & 0.45  & $6.8\\times 10^{35}$  & $\\sim 73$ & 3\\\\ \\efo      & 2.0  & 2.26 & $-$   & $2.3\\times 10^{35}$  & $100$     & 3\\\\ \\tfn      & 0.93 & 3.6  & 0.41  & $1\\times 10^{35}$    & $\\sim50$  & 4\\\\ \\axj$^*$  & 9.0  & 4.6  & $-$   & $7.4\\times 10^{34}$  & $100$     & 5\\\\ \\cxo$^*$  & 0.14 & $-$  & 0.41  & $1.5\\times 10^{35}$  & $0$       & 6\\\\ \\ett$^*$  & 1.05 & 3.75 & 0.668 & $1.6\\times 10^{36}$  & $\\sim70$  & 7\\\\ \\hline \\end{tabular*}\\\\[2pt] ($*$) not confirmed; ($\\dagger$) contribution of the power-law component to the total flux, see Perna et al.~2001 for further discussion; References: (1) Juett et al.~2002; (2) Tiengo et al.~2002; (3) Mereghetti et al.~2002; (4) Patel et al~2001; (5) Torii et al.~1998; (6) Lamb et al.~2003; (7) Ibrahim et al.~2003. \\end{sidewaystable} ",
        "conclusions": "Since the 1996 commissioning of {\\it BeppoSAX} and {\\it RXTE}, our overall picture of AXPs has changed dramatically.  The number of likely AXPs has nearly tripled, and our understanding of these unusual sources' properties has improved tremendously.  Perhaps the single most important discovery is that the apparent resemblance of AXPs with SGRs noted by Thompson \\& Duncan in 1995 is more than skin deep: with the discovery of bursts from AXPs, the two source classes are now united unambiguously.  Our next challenge is to learn how to extract physically interesting information from AXP and SGR observations.  In this sense, their study is still in its infancy, with observations ahead of theory.  This said, there are obvious possibilities for fruitful observational investigation of AXPs. First, it would be nice to have a direct measurement of the inferred high magnetic field.  As no X-ray spectral features have been forthcoming, another avenue is needed.  X-ray polarization observations are an excellent possibility, particularly for the brightest AXPs. The detection of such polarization, in addition to confirming the high magnetic field, would be the first demonstration of the birefringence of the vacuum, as predicted by QED.  In the shorter term, glitches in AXPs may offer a practical method of constraining the structure and physics of these objects.  The simulteneity of the 1E~2259+586 glitch with its outburst and associated radiative changes, we suspect, is telling us a lot about the stellar structure.  Continued patient timing of these objects has the potential to reveal correlations between glitch properties like amplitude and relaxation time scales with radiative properties, which will help us understand properties of the highly magnetized crust and superfluid interior. Optical/IR observations also offer hope of constraining magnetar outer magnetosphere processes, although its origin is not yet clear. Finally, there is the open question of the radio pulsar/AXP connection.  Recently, several radio pulsars having inferred magnetic fields {\\it higher} than that of 1E~2259+586 have been discovered, yet with no evidence for any AXP-like X-ray emission \\citep[e.g.][]{pkc00,msk+03}.  This is puzzling.  It may simply reflect the fact that the magnetic field measured by $P$ and $\\dot{P}$ is approximate only, in which case the discovery of more such radio pulsars should eventually result in the identification of the ``missing link.''  \\bigskip Funding for this work comes from NSERC (via a Discovery Grant and Steacie Supplement), the Canada Research Chair Program, NATEQ (via Team and Observatoire de Mont Megantique Grants), CIAR (via a Fellowship), and NASA (via the LTSA program)."
    },
    "0402/astro-ph0402340_arXiv.txt": {
        "abstract": "We built three models for the gravitational field of the Galactic bar. These models are an inhomogeneous ellipsoid, an inhomogeneous prolate spheroid, and a superposition of four inhomogeneous ellipsoids. Among the three models,  the superposition provides our best approximation to the observed boxy mass distribution of the Galactic bar. Adding the bar component to an axisymmetric Galactic model, we have calculated stellar midplane orbits and orbits of some globular clusters with known kinematical data. For all models we find a secular dispersion effect upon the orbital energy and angular momentum, as measured in the Galactic inertial  frame. This effect might be relevant to explain the orbital  prograde-retrograde distribution of globular clusters. For the stellar  kinematics, we study the connection between the sense of orbital motion  in the midplane and the onset of chaos in the presence of the bar. In the inner region of the bar, chaos is induced by an axisymmetric central component (bulge) and it arises in orbits that change its orbital sense from prograde to retrograde and vice versa as seen from an inertial reference frame. Outside the bar region,  chaos appears only in prograde orbits. Our results concerning such connection are consistent and extend those obtained for midplane orbits in the presence of only a spiral pattern in the  axisymmetric Galactic model. ",
        "introduction": "In the past few years it has been finally accepted the existence  of a bar in the center of our Galaxy. Some studies providing evidence for this Galactic component are the kinematical  data obtained in HI 21-cm emission, CO, and CS (Sanders \\& Prendergast 1974;  Liszt \\& Burton 1980; Gerhard \\& Vietri 1986; Binney \\etal 1991), the  Galactic center stellar distribution with Mira Variables from  IRAS (Harmon \\& Gilmore 1988; Nakada \\etal 1991; Weinberg 1992),  and the results of the COBE/DIRBE satellite (Weiland  \\etal 1994, and models based on these observations, such as Dwek \\etal 1995; Fux 1997; Freudenreich 1998; Beaulieu \\etal 2000;  Bissantz \\etal 2003). Based on these evidences and due to the expected importance of a non-axisymmetric  galactic component to the stellar and gas dynamics, we have  constructed three models for the gravitational  potential of the Galactic bar and studied their dynamical effects on point masses orbiting the Galaxy.  Several authors have studied and modeled in many ways the gravitational potential of bars. The simplest model is the two-dimensional one, for which  the potential has the form $\\Phi_{Bar}(R,\\varphi)=g(R)cos(2\\varphi)$  (Contopoulos \\& Papayannopoulos 1980); $R,\\varphi$ are polar coordinates  in the Galactic plane, $\\varphi$ being measured with respect to the long axis of the bar; $g$ is the amplitude. Other models consider a three-dimensional mass  distribution with similar stratification in ellipsoids or prolate  spheroids given by  \\begin{equation}  \\rho_{Bar}(x,y,z) = \\left\\{ {\\rho_c{(1- m^2)}^n\\atop 0}{,\\ \\ m\\leq 1, \\atop ,\\ \\ m\\geq 1,}\\right. \\label{rhob} \\end{equation}  \\noindent with $\\rho_c$ the central density,  $m^2 = \\frac{x^2}{a^2}+\\frac{y^2}{b^2}+\\frac{z^2}{c^2}$, and  $a>b\\geq c$ the respective semi-axes (ellipsoid: $a>b>c$;  prolate spheroid: $a>b=c$).  \\noindent With $n$ an integer, the density in equation (\\ref{rhob}) corresponds to Ferrers ellipsoids (Ferrers 1877).  The ellipsoidal case with $n=0$, i.e., an  homogeneous ellipsoid, has been considered by Sanders \\&  Tubbs (1980). Inhomogeneous cases ($n\\neq 0$) with a prolate  shape have been considered by Papayannopoulos \\& Petrou (1983);  Petrou \\& Papayannopoulos (1983); Athanassoula \\etal (1983); Teuben \\& Sanders (1985); Shlosman \\& Heller (2002).  Inhomogeneous ellipsoidal models are considered by  Pfenniger (1984), and Kaufmann \\& Contopoulos (1996).  A more elaborated model has been presented by Zhao (1996), who employed the multipolar expansion technique  given by Hernquist \\& Ostriker (1992) to obtain the gravitational potential of a ``boxy'' mass distribution that is observed  in iso-density contours of edge-on galaxies, and in our own Galaxy as well (Freudenreich 1998).  With the purpose of building a complete three-dimensional model for the Milky Way, we have modeled the Galactic bar in three different ways using the available observational parameters; the main parameter being the observed density (Freudenreich 1998; see Section \\ref{models}),  that cannot be fitted with a simple model such as that of equation  (\\ref{rhob}). Our three models are based on the density considered in Model S of Freudenreich (1998). The first model is an inhomogeneous ellipsoid;  the second is an inhomogeneous prolate spheroid, and the third one is a superposition of inhomogeneous ellipsoids. The last model approximates the observed boxy-shaped density stratification of the Galactic bar.  As an application of our models we have analyzed the structure of Poincar\\'e diagrams (or sufrace of section) corresponding to orbits in the Galactic midplane and with the required energy to reach the inner Galactic region.  This work extends a recent study in which the structure was explored in a three-dimensional model for the spiral arms (Pichardo \\etal 2003, hereafter Paper I).  In both papers, the axisymmetric background (Galactic) potential is that of Allen \\& Santill\\'an (1991). A short description of the Galactic model was given in Paper I.  The present paper shows also how the kinematics of globular clusters could be altered by the presence of a bar, via numerical integrations of the orbits  of six globular clusters in our Galaxy. A more extended and detailed study of the effect of the Galactic bar on the kinematics of the whole sample of globular clusters with known absolute proper motions will be presented in a future paper.  In Section \\ref{fit_obs} we review  the   observational parameters of the Galactic bar that are used in  our models. In Section \\ref{models} we describe the three models  for the Galactic bar. Section \\ref{results} gives our results:  the analysis of stellar midplane orbits (Section \\ref{orb_an}),  and the kinematics of six globular clusters (Section \\ref{cumulos}).  In Section \\ref{conclusions} we discuss the results and give our conclusions. Finally in the Appendix, we give a detailed analytical description of our three models and we include a force field analysis. ",
        "conclusions": "We present three models for the gravitational potential of the Galactic bar, based on the mass distribution for this component given by Model S of Freudenreich (1998).  These models are an ellipsoid, a prolate spheroid, and a superposition of four ellipsoids. The models can be easily implemented for a numerical integration of orbits in a non-axisymmetric Galactic potential. In particular, our third model, which is a superposition of four inhomogeneous ellipsoids, gives a good approximation to the boxy mass distribution of the Galactic bar. Thus, in this model the resulting gravitational potential might give relevant results in the analysis of orbits reaching the region of the bar. For orbits lying outside the bar, the detailed modeling of the shape of the bar is less important, and any of the three models can be used.  We have applied our models to orbits in the Galactic plane in the inner Galactic region, and to orbits of some globular clusters. We find that the bar produces a dispersion on the energy and angular momentum, as measured in the Galactic inertial frame. In particular, for orbits with the $z$- component of angular momentum close to zero, this dispersion effect can make an orbit oscillate between prograde and retrograde, resulting in a wider separatrix. In the case of globular clusters, the bar might be responsible for the observed orbital prograde-retrograde distribution. In general, the relative importance of  the bar with respect to a central axisymmetric component determines the  dominant stellar sense of motion, i.e., the larger is this ratio  $(M_{Bar}/M_{Cen}$, where $M_{Bar}$ is the bar mass and $M_{Cen}$ is the mass of  a central component like a bulge), the larger is the population of stars  that will change their sense of motion from prograde to retrograde  (and vice versa), as seen from an inertial reference frame.  In a preliminary analysis of chaotic regions, we have found that a central axisymmetric component induces the onset of orbital chaos in the inner region of the bar, and chaos mainly appears in the orbits that change their sense of motion in the inertial frame of reference, i.e. those that form the separatrix. Outside the bar, chaos only appears after the corotation barrier and only in the prograde orbits. As the central component mass is reduced or disappeared, chaos diminishes or is completely removed for orbits in the inner region of the bar. It is remarkable that the connection between stochastic motion and sense of motion measured from the inertial frame, found for the relatively weak spiral perturbation, is preserved under the much stronger perturbation of the bar. However, the properties of the separatrix are still a subject we consider deserves a more detailed study.  For globular clusters, chaos is found (v.g., Allen and Martos 1988) in the axisymmetrical potential with no need of a perturbation such as the presence of a spiral pattern or a bar. Chaos in those systems is seemingly related to the impulsive nature of the rapid passing of the cluster through the large mass concentration at the central regions of the Galaxy. For midplane motion, results concerning the connection between the angular momentum and chaos are apparently indicating that the physical agent has to do with a secular effect; i.e., prograde orbits with respect to the general motion of the perturbing mass, spiral or bar, tell us about longer times under their influence than that from a rapid encounter, as that expected from retrograde motion. A picture able to include a general explanation for both mechanisms triggering chaotic motion with a physical flavor seems necessary. In the case of planar orbits, Paper I invoked the overlapping of resonances as the standard explanation for the different phenomenology between prograde and retrograde motion in regard to stochasticity. In three dimensional motion, resonances involving vertical oscillations could be the lacking piece for a unified scheme."
    },
    "0402/astro-ph0402206_arXiv.txt": {
        "abstract": " ",
        "introduction": "For many years, it was thought that the strong X-ray emission observed in the cores of rich galaxy clusters results in a cooling flow in which gas settles in the gravitational potential and drops out as cold condensations \\cite{NK_fab94}. Mass inflow rates were estimated to be $\\sim10^2-10^3M_\\odot$ yr$^{-1}$ in some clusters.  However, recent X-ray observations with {\\it Chandra} and {\\it XMM-Newton} have found very little emission from gas cooler than about one-third of the virial temperature \\cite{NK_pet01,NK_pet03}, suggesting that some heating source must prevent gas from cooling below this temperature. Candidate heating mechanisms include (1) energy injection from a central active galactic nucleus (AGN) \\cite{NK_cio01,NK_chu02,NK_bru02,NK_kai03}, and (2) diffusive transport of heat from the outer regions of the cluster to the center via conduction \\cite{NK_tuc83,NK_bre88,NK_nar01,NK_voi02,NK_zak03} or turbulent mixing \\cite{NK_cho03,NK_kim03b,NK_voi04}.  Heating by a central AGN is an attractive idea since many cooling flow clusters show radio jets and lobes that are apparently interacting with the cluster gas \\cite{NK_beg01}.  The power associated with the jets is often comparable to the total X-ray luminosity of the cluster. However, there are some difficulties with this model.  Observations reveal that radio lobes are surrounded by X-ray-bright shells of relatively cool gas \\cite{NK_sch02}, which is a little surprising if this gas is being heated by the bubble.  In addition, if the heating rate (per unit volume) of the gas by the AGN varies as $\\rho^\\alpha$, thermal stability requires $\\alpha>1.5$ \\cite{NK_zak03}; such a heating law does not seem natural.  (Stability is not an issue if AGN heating is episodic \\cite{NK_kai03}).  Finally, no good correlation is seen between the AGN radio luminosity and the X-ray cooling rate \\cite{NK_voi04}.  Since the cooling cores of clusters have a lower temperature than the rest of the cluster, diffusive processes can bring heat to the center from the outside, provided the diffusion coefficient is large enough. An ordered magnetic field would strongly suppress cross-field diffusion of thermal electrons, and this argument has been traditionally invoked for ignoring thermal conduction.  However, if the field lines are chaotically tangled over a wide range of length scales, the isotropic conduction coefficient $\\kappa_{\\rm cond}$ can be as much as a few tens of per cent of the Spitzer value $\\kappa_{\\rm Sp}$ \\cite{NK_nar01,NK_cha03}, which may be sufficient to supply the necessary heat to the cluster core.  Turbulent mixing is another diffusive process that can transport energy efficiently to the center \\cite{NK_cho03}.  The turbulence might be sustained by the infall of small groups or subclusters, the motions of galaxies [K. Makishima, this conference], or energy input from AGNs \\cite{NK_dei96,NK_ric01}. The diffusion coefficient required to balance radiative cooling is typically $\\kappa_{\\rm mix}\\sim 1-6\\;\\rm kpc^2\\;Myr^{-1}$, which is similar to values inferred from observations of turbulence in clusters \\cite{NK_kim03b,NK_voi04}.  In a series of papers \\cite{NK_zak03,NK_kim03b,NK_kim03a}, we have studied equilibrium models of galaxy clusters with thermal conduction and turbulent mixing.  We summarize here the main results of this work. ",
        "conclusions": "The thermal conduction and turbulent mixing models have certain attractive properties which ultimately are due to the fact that both models involve diffusive transport.  Diffusion not only allows heat to move into the cluster center from the outside, it also irons out perturbations and thereby helps to control thermal instability.  What is interesting is that the amount of diffusion required to fit the observations is comparable to that predicted by theoretical arguments.  Two caveats need to be mentioned.  First, the presence of cold fronts in many clusters \\cite{NK_mar00,NK_vik01} indicates that large temperature and entropy jumps are able to survive in some regions of the hot gas.  Diffusion is clearly suppressed across these surfaces. It is possible that cold fronts are special regions where the magnetic field is combed out parallel to the front, thereby suppressing cross-field conduction temporarily \\cite{NK_vik01,NK_zak03}.  Second, all we have shown is that a cluster with the observed density and temperature profile would be in hydrostatic and thermal equilibrium and would be fairly stable.  However, we have not explained how the cluster reaches the observed state starting from generic initial conditions.  Time-dependent simulations show that a cluster with thermal conduction would either slowly evolve to an isothermal state if its initial density is less than a critical density, or develop a catastophic cooling flow otherwise \\cite{NK_bre88}.  Does the current observed state result from an initial rapid mass dropout (which decreases the density) and subsequent slow evolution with diffusive heating of an once overdense cluster \\cite{NK_kim03a}? Are other heating mechanisms, e.g., AGNs, necessary to explain the present state of clusters?  Answers to these questions are of fundamental importance to understanding clusters and more generally galaxy formation.  \\vspace{0.5cm} \\noindent {\\bf Acknowledgments.} The work reported here was supported in part by NASA grant NAG5-10780 and NSF grant AST 0307433."
    },
    "0402/astro-ph0402083_arXiv.txt": {
        "abstract": "s{ Kilometer-scale neutrino detectors such as IceCube are discovery instruments covering nuclear and particle physics, cosmology and astronomy. Examples of their multidisciplinary missions include the search for the particle nature of dark matter and for additional small dimensions of space. In the end, their conceptual design is very much anchored to the observational fact that Nature accelerates protons and photons to energies in excess of $10^{20}$ and $10^{13}$\\,eV, respectively. The cosmic ray connection sets the scale of cosmic neutrino fluxes. In this context, we discuss the first results of the completed AMANDA detector and the reach of its extension, IceCube. Similar experiments are under construction in the Mediterranean. Neutrino astronomy is also expanding in new directions with efforts to detect air showers, acoustic and radio signals initiated by super-EeV neutrinos.} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0402/astro-ph0402560_arXiv.txt": {
        "abstract": "SS433 is well-known for its precessing twin jets having optical bullets inferred through {\\it spectroscopic} observation of $H_\\alpha$ lines. Recently, Chakrabarti et al. (2002) described processes which may be operating in accretion disk of SS433 to produce these bullets. In a recent multi-wavelength campaign, we find sharp rise in intensity in time-scales of few minutes in X-rays, IR and radio waves through {\\it photometric } studies. We interpret them to be possible evidence of ejection of bullet-like features from accretion disks. ",
        "introduction": "SS433 is a well studied bright emission line compact system which is ejecting matter in symmetrically opposite directions at a speed of $v_{jet} \\sim 0.26$c. It has a mass-losing (${\\dot M} \\sim 10^{-4} M_\\odot/$yr) companion orbiting in $13.1$d. The jets are precessing in $162.15$d around the symmetry axis. The velocity is remarkably constant to within a percent or so (Margon, 1984; Gies et al. 2002). As a result of the remarkable constancy in the precession time scales and the velocity, the instantaneous locations of the red and blue-shifted $H_\\alpha$ lines are well predicted by the so-called `kinematic-model' (Abell \\& Margon, 1979) and the compilation of twenty years of timing properties (Eikenberry et al. 2001) suggest that kinematic model can explain the general variation of the red and blue-shifts very well.  A very exciting observation, made within years of the discovery of SS433, suggests that the passage of the jets is not continuous, but as if through successive and discrete bullet-like entities, at least in the optical region (Grandi, 1981; Brown et al. 1991). The $H_\\alpha$ lines were seen to brighten up and fade away without changing their red/blue-shifts, indicating the brightened bullets are radially ejected and do not have any rotational velocity component. Since the bullets of energy $\\sim 10^{35}$ ergs do not change their speed for a considerable time ($\\sim 1-2 $ days), Chakrabarti et al (2002) postulated that they must be ejected from accretion disk itself. They presented a mechanism to produce quasi-regular bullets. Using results of numerical simulations involving oscillation of shocks in accretion disks, they concluded that in the normal circumstances, $50-1000$s interval is expected in between the bullet ejection. These bullets would be ejected from X-ray emitting region and propagate through optical, infra-red ($\\sim 10^{13-14}$cm) and finally to radio emitting region at $\\gsim 10^{15}$cm (roughly the distance covered in a day with $v\\sim v_{jet}$) or so. Thus if the object is in a low or quiescence state, each individual bullet flaring and dying away in a few minutes time scale, should be observable not only in optical wavelengths (Grandi, 1981; Margon, 1984; Brown et al. 1991; Gies et al. 2002) but also in all the wavelengths, including X-ray, IR and radio emitting regions. So far, no such observations of individual bullets has been reported in the literature and it was necessary to make a multi-wavelength observation at relatively quieter states.  In this {\\it Letter}, we present some results of our multi-wavelength studies. From the long term analysis of radio flares (Bonsignori-Facondi, Padrielli, Montebugnoli and Barbieri, 1986; Vermeulen et al. 1993) it is known that in between big flares which occupy $\\sim 20\\%$ of the time the object may go to very quiescence state. So, it is likely that one could `catch' these bullets in action provided observations are carried with very short time resolution. Our multi-wavelength observations lasted during 25th of Sept., 2002 to 6th of October, 2002 with X-ray, infra-red, optical, radio observations made simultaneously on the 27th of September, 2002. Here, we report only X-ray, infra-red and radio observations of 27th and 29th of September, 2002. Optical studies required longer integration times and these results along with other days of observations would be reported elsewhere (Chakrabarti et al. 2003).  Our main results indicate that there are considerable variations in the timescale of minutes in all the wavelengths. These may be called micro-flares. When Fourier transform is made, some excess power is observed in $2-8$ minutes time scale (often beyond 3$\\sigma$ level). The X-ray count rate was found to increase by $15-20$\\% within a minute. Since the emitting regions of X-ray, IR and radio are not well known with absolute certainly, while duration of the flares last less than a few minutes, we could not prove beyond doubt that there are indeed correlations among the micro-flares observed in these wavelengths. However variabilities we find are not flicker type or shot noise type in the sense that the power density spectrum (PDS) is not of $1/f$ type and the duration is not very short (i.e., $<1$s). We therefore believe that we may have found evidences of bullet ejection through these observations in other wavelengths. ",
        "conclusions": "In this {\\it Letter}, we presented results of our multi-wavelength observations which were save at short time intervals. From the analysis of the observations of radio, IR and X-ray in the quiescence state we conclude that we may be observing ejection events of bullet-like features from the accretion disk in time scales of $\\sim 2-8$ minutes. Identification of small micro-flare events with those those of bullet ejection is derived from the time scales of variabilities, which are roughly the same in all these wavelengths. We find their presence in X-ray ($\\lsim 10^{11-12}$cm), IR ($\\lsim 10^{13-14}$cm) and radio ($\\lsim 10^{15}$cm) emission regions. Vermeulen et al. (1993) found evidence for optical bullets with life-time of $1-2$d. This is perhaps due to the propagation of a burst of indistinguishable bullets and not due to a single one. We identified micro-flare like features in all these observations which may be signatures of the bullets. Count rate of X-ray was seen to increase $15-20$\\% in a matter of a minute. One way to actually identify each bullet could have been to follow them from X-ray region outwards. This will require very careful time delay measurements since the distances of emission regions are not very accurately known to follow a feature of duration of a minute. We exclude the possibility that what we see were flicker noise since neither the duration nor the PDS properties match with those of flicker noise. One could have perhaps distinguished the energetic bullets by observing polarization properties of the radio-emissions during the short-lived flares, but unfortunately due to technical reason this observation could not be carried out. We plan to do such an observation in near future.  SKC thanks Dr. J. Swank of NASA/GSFC and Prof. A. P. Rao of GMRT/NCRA for giving time in RXTE and GMRT for observations. He also thanks Mr. J. Kodilkar of GMRT/NCRA for his assistance in analysing the radio data. This work is supported in part by CSIR fellowship (SP) and a DST project (SKC and AN). Authors thank the unknown referee for suggestions leading to considerable improvement of the observational aspects."
    },
    "0402/astro-ph0402426_arXiv.txt": {
        "abstract": "The electric field induced by extensive air showers generated by high energy cosmic rays is considered and, more specifically, its dependence on the shower incident angle. It is shown that for distances between the shower axis and the observation point larger than a few hundred meters, non-vertical showers produce larger fields than vertical ones. This may open up new prospects since, to some extent, the consideration of non-vertical showers modifies the scope of the radio-detection domain. ",
        "introduction": "\\label{Introduction} The observation of ultra high-energy cosmic rays ($E\\geq 10^{18}$~eV) is nowadays a question of paramount importance. They can be detected indirectly via the observation of the Extensive Air Showers (EAS) they create while interacting with the atmosphere. The induced secondary particles generate a radio-electric field whose measurements may provide an alternative to the experimental techniques consisting of counting the shower particles collected in ground detectors. It turns out that the radio signal becomes significant provided that an efficient charge separation mechanism occurs~\\cite{allan}.  The radio frequency component associated with EAS was studied in the 1960's. At that time, the purpose was mainly to demonstrate the observability of such a radio-electric field~\\cite{askaryan,jelley}. The strategy was then to detect a radio pulse in coincidence with particles in an experimental set-up consisting of one antenna located as close as possible to some particle detectors. Such a device selects mainly \\emph{vertical} showers.  The purpose of this letter is to stress that, in the case of not-too-close impact parameters (i.e., large distances between the shower axis and the observation point), vertical showers produce a radio-electric signal weaker in amplitude than oblique ones. The reason is the ultra-relativistic character of most of the shower particles that leads to a strong forward enhancement of the electric field. For a given impact parameter, this means that the point at which a charge produces the largest signal along its trajectory is located far above the point of closest approach. This forward-peaked structure is counterbalanced by the rise in the number of charges as the shower develops in the dense layers of atmosphere. Quantitatively, as explained in Sect.~\\ref{sec:calc}, the combination of both effects will favor the observation of \\emph{non-vertical} showers.  Sect.~\\ref{sec:outlook} unfolds the consequence of this observation. The point is that, using a large array of scattered antennas in order to compensate for the scarcity of high-energy cosmic rays, there is some interest in examining the possibility of radio detection of EAS at a somewhat larger impact parameter. According to the above statement, such an apparatus, therefore, puts emphasis on the detection of non-vertical hadron-induced showers. ",
        "conclusions": ""
    },
    "0402/astro-ph0402092.txt": {
        "abstract": "We study the influence of shower fluctuations, and the possible presence of different nuclear species in the primary cosmic ray spectrum, on the experimental determination of both shower energy and the proton air inelastic cross section from studies of the longitudinal development of atmospheric showers in fluorescence experiments. We investigate the potential of track length integral and shower size at maximum as estimators of shower energy.  We find that at very high energy ($\\sim 10^{19}-10^{20}$ eV) the error of the total energy assignment is dominated by the dependence on the hadronic interaction model, and is of the order of 5\\%. At lower energy  ($\\sim 10^{17}-10^{18}$ eV), the uncertainty of the energy determination due to the limited knowledge of the primary cosmic ray composition is more important. The distribution of shower maximum, $X_{\\rm max}$, is discussed as a measure of the proton-air cross section. Uncertainties in a possible experimental measurement of this cross section introduced by intrinsic shower fluctuations, the model of hadronic interactions, and the unknown mixture of primary nuclei in the cosmic radiation are numerically evaluated. ",
        "introduction": "Introduction}  The fluorescence technique of ultra high energy cosmic ray (UHECR) detection was first explored in the pioneering Fly's Eye detector \\cite{Baltrusaitis88}, and is currently being used in its successor, the high resolution Fly's Eye (HiRes) \\cite{Abu-Zayyad00}, as well as in the Pierre Auger Observatory \\cite{Cronin95} that is currently under construction. The underlying idea is the detection of atmospheric nitrogen fluorescence light induced by the passage of charged particles through the atmosphere. The number of charged particles at depth $X$ in the atmosphere, $N(X)$, i.e. the longitudinal shower profile, can be extracted from data because $N(X)$ is to a good approximation proportional to the amount of emitted fluorescence light. In this approximation, the total energy that goes into electrons and positrons (the electromagnetic energy $E_{\\rm em}$ from now on) is obtained by integration of the shower longitudinal profile \\cite{Song00} % \\begin{equation} E_{\\rm em}=\\alpha_{\\rm eff} \\int_0^{\\infty} N(X)dX \\label{eq:Eem} \\end{equation} where $\\alpha_{\\rm eff}$ is the average (effective) ionization loss rate which is usually taken as a constant over the entire shower and is given by $\\sim 2.19~{\\rm MeV/g~cm^{-2}}$ \\cite{Baltrusaitis85,Song00}.  The integral on the right hand side of Eq.~(\\ref{eq:Eem}) represents the total track length of all charged particles in the shower projected onto the shower axis. Electrons and positrons constitute the bulk of the charged particles in a shower and contribute most to the production of fluorescence light.  In the following we neglect the contribution of muons and other charged particles to the production of fluorescence light, which is of the order of 2\\% (see discussion in \\cite{Risse03b}).  %%%%%% I think this is not what we do -> commented out (RE) %%%%%%%%% % The contribution of ionization loss of muons can % be estimated from Fig. 14 of Ref.~\\cite{Alvarez02} as approximately % 1\\% of that of electrons and positrons in the energy range % we consider in this paper. % Below we evaluate $\\alpha_{\\rm eff}$ % from photon-induced showers, but we use it to convert the total % track length integral (including the small contribution from muons) to % shower energy.  Thus the contribution of ionization loss by muons % in the atmosphere is included in an approximate way, but not neglected. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  It is generally assumed that the fluorescence rate is proportional to the ionization energy loss rate $dE/dX$, although this has been experimentally proved only to some extent \\cite{Kakimoto96,Nagano:2003zn}. Consequently in order to estimate shower energy, there is in principle no need to convert the measured fluorescence intensity first to a particle number, and then relate the total track length to the energy of the shower through Eq.~(\\ref{eq:Eem}). The total ionization energy deposit can instead be obtained from the fluorescence intensity and can be used directly as an energy estimate \\cite{Dawson02}. However, as long as the lateral spread of shower particles is correctly accounted for \\cite{Alvarez-Muniz03}, the conversion of fluorescence light intensity to number of particles and then to energy through Eq.~(\\ref{eq:Eem}) does not lead in principle to observable errors mainly for two reasons: Firstly the ionization energy deposit depends only weakly on the particle energy, and secondly the shape of the energy spectrum of particles in an air shower changes only slowly with the traversed depth. Only in the very early evolution stage of a shower is the particle energy spectrum significantly harder than that at the shower maximum. The corresponding energy deposit is higher by up to a factor of 1.5, but due to the small number of particles, the resulting error in the energy estimation is negligible \\cite{Risse02a}.  There are several additional factors, such as air pressure, density and humidity, that influence the relation of the fluorescence intensity to energy deposit and particle numbers. The discussion of these aspects, including the conversion of the observed light curve to a longitudinal shower profile are beyond the scope of this work.  In this article we investigate the longitudinal shower profile as an experiment-independent quantity and study its relation to the energy and mass of the primary particle. In Sec.~\\ref{sec:energy} the track length integral and the particle number at shower maximum are compared as energy estimators under the assumption of an unknown cosmic ray composition above $10^{17}$ eV. The model and mass dependence of the invisible energy carried by neutrinos and energetic muons is calculated for the QGSjet \\cite{qgsjet} and SIBYLL \\cite{Engel99,Fletcher94} models of hadronic interactions. The mean position of the shower maximum, $X_{\\rm max}$, and its distribution is discussed as a measure of the primary cosmic ray composition and proton-air cross section in Sec.~\\ref{sec:xmax}. A summary and conclusions are given in Sec.~\\ref{sec:dis}.   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "}  We have used a hybrid simulation scheme~\\cite{Alvarez02} for a quantitative evaluation of certain systematic uncertainties in the interpretation of fluorescence measurements of giant air showers. The sources of uncertainty we investigate include the model of hadronic interactions used for shower simulation, the unknown mixture of primary nuclei in the cosmic radiation and intrinsic fluctuations in shower development.  We do not investigate uncertainties in shower reconstruction, detector acceptance or other technical aspects related to properties of the environment or performance of the detectors.  As an illustration of uncertainties arising from the need to extrapolate hadronic interaction models outside the kinematical region and energies explored by accelerator experiments, we compared two specific interaction models,  QGSjet~\\cite{qgsjet} and SIBYLL~\\cite{Engel99}, both of which agree with each other and with a range of accelerator data for $\\sqrt{s}\\sim$~TeV and below.  We find that the correction for unseen energy (i.e. energy lost to neutrinos and muons that reach the ground) is consistently larger for QGSjet than for SIBYLL.  The difference is such that, for a given track length integral, the assigned energy will be about 5\\% higher when the same data are interpreted with QGSjet rather than with SIBYLL.  For primary energies below about $10^{19}$~eV/nucleus the fraction of unseen energy depends significantly on the mass of the primary nucleus. For example, at $10^{17}$~eV with SIBYLL about 7\\% of the primary energy is lost to neutrinos and muons as compared to 13\\% for iron.  In view of the steep cosmic-ray energy spectrum, knowledge of the energy resolution is of great importance.  The track length integral can be used to assign energy when a sufficient portion of the profile is measured to fix the parameters needed to complete the integral.  It has an intrinsic resolution of 2-4\\%, depending somewhat on interaction model and energy (narrower at higher energy). Size of shower at maximum gives only a marginally broader energy resolution (3-5\\%) and can be used when much of the profile after maximum is not measured, provided care is taken to correct for a slightly asymmetric distribution.   There is in both methods some dependence on primary mass of the relation between the measured quantity and the primary energy which leads to a $\\sim 5$\\% systematic uncertainty if the primary mass is not separately determined.  Intrinsic fluctuations in shower development (after the first interaction) affect the relation between the interaction length ($\\lambda_{\\rm int}$) and the slope $\\Lambda$ that describes the exponential tail of the $X_{\\rm max}$ distribution. The relation is often expressed with a `$k$ factor\u00b4 as $\\Lambda\\,=\\,k\\times\\lambda_{\\rm int}$.  Differences in $k$ factors for the range of models studied here are at the level of 5-7\\%, implying a similar uncertainty in the p-air cross section that may be inferred from measurements of shower profiles. Further uncertainties arise to the extent that an unknown fraction of helium and other nuclei contaminate the tail of the measured $X_{\\rm max}$ distribution.  \\noindent {\\bf Acknowledgments} J.A. Ortiz is supported by CNPq/Brasil and acknowledges the Bartol Research Institute for its hospitality. This research is supported in part by NASA Grant NAG5-10919. RE, TKG \\& TS are also supported by the US Department of Energy contract DE-FG02 91ER 40626. JA-M is also supported by MCyT (FPA 2001-3837 and FPA 2002-01161). The simulations presented here were performed on Beowulf clusters funded by NSF grant ATM-9977692."
    },
    "0402/astro-ph0402610_arXiv.txt": {
        "abstract": "{ We present detailed gas-phase chemical models for the envelope of the low-mass star-forming region IRAS 16293-2422.  By considering both time- and space-dependent chemistry, these models are used to study both the physical structure proposed by Sch\\\"oier et al.  (\\cite{schoieretal2002}), as well as the chemical evolution of this region.  A new feature of our study is the use of a detailed, self-consistent radiative transfer model to translate the model abundances into line strengths and compare them directly with observations of a total of 76 transitions for 18 chemical species, and their isotopes.  The model can reproduce many of the line strengths observed within 50\\%. The best fit is for times in the range of $3\\times 10^3 - 3\\times 10^4$ yrs and requires only minor modifications to our model for the high-mass star-forming region AFGL 2591.  The ionization rate for the source may be higher than previously expected -- either due to an enhanced cosmic-ray ionization rate, or, more probably, to the presence of X-ray induced ionization from the center.  A significant fraction of the CO is found to desorb in the temperature range of 15--40~K; below this temperature $\\sim$90\\% or more of the CO is frozen out. The inability of the model to explain the HCS$^{+}$, C$_{2}$H, and OCS abundances suggests the importance of further laboratory studies of basic reaction rates.  Finally, predictions of the abundances and spatial distributions of other species which could be observed by future facilities (e.g. Herschel-HIFI, SOFIA, millimeter arrays) are provided. ",
        "introduction": "The distribution and composition of dust and gas around isolated low-mass young stellar objects (YSOs) is receiving increased attention both observationally and theoretically.  While the general process of low mass star formation is relatively well understood (see e.g. Shu et al. \\cite{shu1993}; Evans \\cite{evans1999}, and others), many details on the chemical and physical structure at different stages of evolution remain uncertain.  In particular, the warm and dense gas in the very interior of these star-forming regions provides a rich chemical environment with which to probe their structure, properties, and evolution.  Beyond their own intrinsic interest, these regions may provide a link to the so-called hot cores observed toward many high-mass star-forming regions (e.g., Walmsley \\& Schilke \\cite{walmsleyetal1993}).  Rapid advances in both observational and modeling capabilities allow much more quantitative studies of the chemistry in YSO envelopes than was possible even a few years ago.  Several different steps can be distinguished (see Fig. \\ref{modelingflowchart}). Thanks to the advent of large-format bolometer arrays, most studies nowadays start with an analysis of the spatial distribution of the submillimeter continuum emission from dust and its spectral energy distribution (SED) (e.g., Shirley et al.\\ \\cite{shirleyetal2000}; J{\\o}rgensen et al.\\ \\cite{joergensenetal2002}; Sch\\\"oier et al.\\ \\cite{schoieretal2002}). Through continuum radiative transfer calculations, both the density profile $n\\propto r^{-p}$ as a function of radius $r$ and the dust temperature structure $T_{\\rm dust}(r)$ can be determined self-consistently. For the gas, two approaches can subsequently be taken. In atmospheric chemistry, these two cases are commonly known as the `forward' and `backward' or `retrieval' methods. In the `empirical model' (`retrieval'), $T_{\\rm gas}$ is taken to be equal to $T_{\\rm dust}$ and the excitation, radiative transfer and fluxes of the various molecular lines are calculated for an assumed abundance profile $x(r)=n({\\rm X})(r)/n$(H$_2$)($r$). This trial abundance profile is then varied until the best agreement with observations is obtained. In practice, only two types of abundance profiles are considered: a constant abundance throughout the envelope or a `jump' profile in which the abundance is increased by a large factor in the inner warm region due to ice evaporation.  Such models have successfully been applied to both high- (e.g., van der Tak et al.\\ \\cite{vdt2000}) and low-mass YSOs (e.g., Ceccarelli et al.\\ \\cite{ceccarellietal2000b}, Sch\\\"oier et al.\\ \\cite{schoieretal2002}), and work best for molecules for which a large set of lines originating from levels with a range in energy has been observed.  This method provides abundances for comparison with chemical models, but does not test the chemical networks directly. \\begin{figure*} \\resizebox{\\hsize}{!} {\\includegraphics[angle=-90]{0476f1.eps}} \\caption[]{An overview of methods used for constraining the physical and chemical structure of YSO envelopes. In general, the steps include the determination of the physical structure through dust modeling, calculation of gas temperature, and adoption of a chemical model.  This combination produces observables (column densities, line fluxes, etc.) to compare with observations.  The source parameters are determined by adjusting them until a best fit is obtained.  Two important points of divergence include the use of self-consistent vs. approximate radiative transfer, and the use of a time-dependent chemical network (`Full chemical model') vs. simple trial abundances (`Empirical model') (figure adapted from van Dishoeck \\& van der Tak \\cite{vDandvdTreview2000}, and van Dishoeck \\cite{vdreview2003}). } \\label{modelingflowchart} \\end{figure*}   The second, {\\it ab initio} or `forward' approach is the `full chemical model', in which only the density structure derived from the dust is adopted as a starting point. Given initial abundances and a detailed chemical network, the abundances of various molecules can be solved as functions of position and time and the gas temperature can be calculated explicitly by solving the full thermal balance of the gas.  The physical structure $n(r)$ and $T(r)$ can either be taken to be constant with time or to vary according to some (dynamical) prescription. Such time- and space-dependent chemical models have been applied to low-mass YSOs by Ceccarelli, Hollenbach, \\& Tielens (\\cite{cht96}, hereafter CHT96) and Rodgers \\& Charnley (\\cite{rodgerscharnley2003}), and to specific high-mass sources by Millar, MacDonald \\& Gibb (\\cite{millaretal1997}; G34.3+0.15) and Doty et al.\\ (\\cite{dotyetal2002}; AFGL 2591). The main quantities to be determined are the best-fit time (or, in a dynamical model, mass-accretion rate) and other parameters which enter the chemical models, such as the cosmic ray ionization rate. The chemical models themselves can be tested by comparing the abundance profiles $x(r,t_{\\rm fit})$ at the best-fit time with those derived through the empirical method. In this way, they also provide a guideline for more complicated abundance profiles to adopt in the empirical method.  In this paper, we describe a detailed `full chemical model' of the best studied low-mass YSO, IRAS 16293-2422. A novel feature is the addition of a full Monte Carlo radiative transfer calculation of the resulting line fluxes for direct comparison with observations (see bottom-right part of Fig. \\ref{modelingflowchart}).  Such models provide the most complete test of our understanding of the physical and chemical structure of YSO envelopes.  IRAS 16293-2422 is a nearby ($\\sim 160$ pc, Whittet \\cite{whittet1974}) low mass, low luminosity (27 L$_{\\odot}$) protostellar object located within the $\\rho$ Ophiuchus molcular cloud.  It has an exceptionally rich and well-studied spectrum (e.g., Blake et al. \\cite{blakeetal1994}; van Dishoeck et al. \\cite{vandishoecketal1995}; Cecarelli et al. \\cite{ceccarellietal2000a}, \\cite{ceccarellietal2000b}; Sch\\\"oier et al. \\cite{schoieretal2002}), and is therefore considered the prototypical low-mass source for chemical studies, much like Orion is for high-mass objects.  Ceccarelli et al.\\ (2000a,b) used the physical structure based on CHT96 combined with a restricted chemical network to analyze data of H$_2$CO, H$_2$O, and SiO in a `full chemical model', and found strong evidence for large abundance enhancements of these species in the innermost part ($\\leq$150 AU) of the envelope. The evaporated species may subsequently drive a complex `hot core' chemistry leading to the even more complex organic molecules which have recently been detected in IRAS16293-2422 (Cazaux et al. \\cite{cazauxetal2003}).   In a later analysis, Sch\\\"oier et al. (\\cite{schoieretal2002}) combined dust/SED modeling of the physical structure, multiple line observations covering a range of excitation conditions, and a detailed radiative transfer analysis in an `empirical model' to infer the structure of IRAS 16293-2422.  This work supported the general conclusion of a `hot core' where the abundances of key molecules are enhanced by several orders of magnitude due to evaporation of ices.  The model employed only uniform and `jump' abundances, however, which may not be representative of the detailed time- and space-dependent chemistry.  Recently, Doty et al. (\\cite{dotyetal2002}) described such a time- and space-dependent physical/chemical model for static YSO envelopes including the hot core chemistry.  By combining the model results with observations of many species of one particular high-mass YSO, AFGL~2591, it has been shown that it may be possible to not only confirm the gross source structure, but also constrain source properties such as age, ionization rate, and role of grains in determining the chemical structure (see also Boonman et al. \\cite{boonmanetal2003}).  Here the `full chemical model' of Fig. \\ref{modelingflowchart} was adopted, but the self-consistent line radiative transfer was performed for only a subset of the species.   In this paper, we report on the application of the physical/chemical model of Doty et al. (\\cite{dotyetal2002}) to the low-mass YSO IRAS 16293-2422.  These results are combined with a self-consistent radiative transfer model, and applied to the full multi-species, multi-transition dataset of Sch\\\"oier et al. (\\cite{schoieretal2002}). By comparison with the case of AFGL~2591, we can also directly determine the differences in derived model parameters for a low- and a high-mass YSO (van Dishoeck 2003). The models and observations are briefly described in Section 2.  The observations are then used with the models to constrain the source properties in Section 3.  Finally, we summarize the results and conclude in Section 4. ",
        "conclusions": "We have constructed detailed thermal and gas-phase chemical models for IRAS 16293-2422 based upon the physical model of Sch\\\"oier et al. (\\cite{schoieretal2002}). These models were used to probe the validity of the proposed physical structure, as well as study the chemical evolution of the source, and to test the application of our combined `hot-core'/envelope chemistry model of AFGL 2591 to a low-mass `hot-core'-like source. In particular, we find that: \\begin{enumerate}  \\item The combined application of a physical, thermal, and chemical model with detailed radiative transfer is a powerful tool in constraining the structure and evolution of depth-dependent sources.  \\item The time and position dependent model of Doty et al. (\\cite{dotyetal2002}) can be meaningfully applied to IRAS 16293-2422, yielding results qualitatively similar to the massive YSO AFGL 2591.  \\item The best fit for IRAS 16293-2422 occurs for times in the range $3\\times10^{3} < t(\\mathrm{yrs}) < 3 \\times 10^{4}$, consistent with existing infall models (Sect. 3.1).  \\item The best-fit ionization rate in IRAS 16293-2422 is high, $\\sim 5 \\times 10^{-16} - 10^{-15}$ s$^{-1}$, compared with previous results for dense clouds. We propose that this may be due to X-ray emission from the central sources (Sect. 3.2).  \\item Our best fit suggests that important CO desorption occurs at low temperatures, $\\sim 20$ K, and constrained to the range $15 < T_{\\mathrm{CO}} \\mathrm{(K)} < 40$.   Some solid CO may remain at higher temperatures for these timescales. These results are in agreement with recent laboratory data of CO on -- but not mixed with -- a water ice  (Sect. 3.3).  \\item We can also constrain the warm ($T>T_{\\mathrm{CO}}$) and cold ($T<T_{\\mathrm{CO}}$) CO abundances to $x(\\mathrm{CO,warm}) \\sim 10^{-4}$ and $x(\\mathrm{CO,cold}) < 10^{-5}$ and most probably $< 10^{-6}$.  These results are reflected in both the CO lines and in the results from the greater chemical network, and suggest significant ($>$ 90\\%) depletions at low temperatures (Sect. 3.4).  \\item CH$_{3}$OH appears to desorb at temperatures $60 < T(\\mathrm{K}) < 100$, consistent with laboratory expectations.  On the other hand, the comparison between the H$_{2}$CO predictions and observations are insensitive to the amount of cold H$_{2}$CO present (Sect. 3.5).  \\item The chemistry of HNC, C$_{2}$H, and HCS$^{+}$ may not be fully understood.  In particular, it would be useful to measure the branching fraction of dissociative recombination of H$_{2}$NC$^{+}$, the temperature dependence of the reaction O$+$C$_{2}$H, and the ion-molecule reaction rates (HCO$^{+}$, H$_{3}^{+}$, H$_{3}$O$^{+}$) + CS (Sect. 3.6).  \\end{enumerate}"
    },
    "0402/astro-ph0402361_arXiv.txt": {
        "abstract": "{ Using hydrodynamic simulations we investigate the rotational properties and angular momentum evolution of prestellar and protostellar cores formed from gravoturbulent fragmentation of interstellar gas clouds. We find the specific angular momentum~$j$ of molecular cloud cores in the prestellar phase to be on average $\\langle j \\rangle = 7\\times10^{20}\\,\\mathrm{cm^2\\,s^{-1}}$ in our models.  This is comparable to the observed values. A fraction of those cores is gravitationally unstable and goes into collapse to build up protostars and protostellar systems, which then have $\\langle j \\rangle = 8\\times10^{19}\\,\\mathrm{cm^2\\,s^{-1}}$. This is one order of magnitude lower than their parental cores and in agreement with observations of main-sequence binaries. The loss of specific angular momentum during collapse is mostly due to gravitational torques exerted by the ambient turbulent flow as well as by mutual protostellar encounters in a dense cluster environment.  Magnetic torques are not included in our models, these would lead to even larger angular momentum transport.  The ratio of rotational to gravitational energy~$\\beta$ in cloud  cores that go into gravitational collapse turns out to be similar to the observed values. We find, $\\beta$ is roughly conserved during the main collapse phase. This leads to the correlation $j \\propto M^{2/3}$, between specific angular momentum $j$ and core mass $M$.  Although the temporal evolution of the angular momentum of individual protostars or protostellar systems is complex and highly time variable, this correlation holds well in a statistical sense for a wide range of turbulent environmental parameters. In addition, high turbulent Mach numbers result in the formation of more numerous protostellar cores with, on average, lower mass. Therefore, models with larger Mach numbers result in cores with lower specific angular momentum.  We find, however, no dependence on the spatial scale of the turbulence.  Our models predict a close correlation between the angular momentum vectors of neighboring protostars during their initial accretion phase. Possible observational signatures are aligned disks and parallel outflows. The latter are indeed observed in some low-mass isolated Bok globules. ",
        "introduction": "Angular momentum plays a pivotal role in star formation.  The amount of specific angular momentum determines whether a collapsing protostellar core will form a single star or a binary or higher-order multiple system.  Stars are thought to form by gravoturbulent fragmentation in interstellar clouds. The supersonic turbulence ubiquitously observed in molecular gas generates strong density fluctuations with gravity taking over in the densest and most massive regions.  Once gas clumps become gravitationally unstable, collapse sets in and the central density increases until a protostellar object forms and grows in mass via accretion from the infalling envelope. Various aspects of the relation between supersonic turbulence and star formation have been discussed, e.g., by \\citet{HUN82}, \\citet{ELM93}, \\citet{LAR95}, \\citet{PAD95}, \\citet{BAL99a,BAL99b,BAL03}, \\citet{PAD99,PAD02}, \\citet{VAZ00}, \\citet*{KLE00b}, \\citet*{HEI01}, \\citet{KLE01c,KLE01a}, \\citet{GAM03}, or \\citet*{VAZ03}. In particular see the reviews by \\citet{LAR03} and \\citet{MAC04}.  This dynamic picture of gravoturbulent star formation challenges the so called ``standard theory'' where stars build up from the ``inside-out'' collapse of singular isothermal spheres, which are generally assumed to result from the quasistatic contraction of magnetically supported cloud cores due to ambipolar diffusion \\citep*{SHU77,SHU87}. This picture, however, has always received strong criticism \\citep[e.g.,][ for a critical discussion]{WHI85,WHI96,NAK98}.  In particular, it seems strongly biased toward the formation of single stars \\citep{WHI96} which is in contradiction to the observational fact that most (if not all) stars form as members of binary or higher-order multiple systems \\citep[see, e.g., the reviews by][ and references therein]{BOD00,MAT00}.  Gravitational collapse in the astrophysical context always involves solving the angular momentum problem.  It results from the blatant discrepancy between the specific angular momentum observed in low-density gas on large scales and the amount of rotation present after collapse \\citep{SPI68,BOD95}. The source of angular momentum on large scales lies in the differential rotation of the galactic disk and, closely related to that, on intermediate to small scales it results from the high degree of vorticity inextricably adherent to turbulent flows. The typical specific angular momentum $j$ of molecular cloud material, e.g., on scales of about 1$\\,$pc is $j\\approx 10^{23}\\,$cm$^2$\\,s$^{-1}$, while on scales of cloud cores, say below $0.1\\,$pc, it is of order of $10^{21}\\,$cm$^2$\\,s$^{-1}$. A binary G star with a  orbital period of 3 days has $j \\approx 10^{19}\\,$cm$^2$\\,s$^{-1}$, while the spin of a typical T Tauri star is a few $\\times 10^{17}\\,$cm$^2$\\,s$^{-1}$. Our own Sun rotates only with $j\\approx 10^{15}\\,$cm$^2$\\,s$^{-1}$. That means, during the process of star formation most of the initial angular momentum is removed from the collapsed object.   The presence of magnetic fields, in principle, provides a viable mechanism for locally reducing the angular momentum through magnetic braking. This was treated approximately by \\citet*{EBE60}, and later calculated accurately by \\citet{MOU79,MOU80}.  The criterion for effective braking is essentially that the outgoing helical Alfv\\'en waves from the rotating cloud have to couple to the ambient medium over a volume that contains roughly the same mass as the cloud itself. For the strong magnetic fields required by the standard theory of star formation, the deceleration time can be less than the free-fall time, leading to efficient transfer of angular momentum away from collapsing cores, and thus, to the formation of single stars. Field strengths small enough to allow for binary formation cannot provide support against collapse, thus pointing toward a more dynamic picture of star formation as offered by gravoturbulent fragmentation.  It is therefore a crucial test for any theory of star formation whether it can produce the required angular momentum loss during collapse while at the same time explain the high numbers of binaries and multiple stellar systems observed \\citep[e.g.,][]{DUQ91b,HAL03}.  In a semiempirical analysis of isolated binary star formation \\citet{FIS04} presented the effects of turbulence in the initial state of the gas on binary orbital parameters. These properties were in agreement with observations if a significant loss of angular momentum was assumed. In the current investigation we focus on models of non-magnetic, supersonically turbulent, self-gravitating clouds and analyze the time evolution of angular momentum during formation and subsequent collapse of protostellar cores. Our main question is whether gravoturbulent fragmentation can solve or at least ease the so called ``angular momentum'' problem without invoking the presence of magnetic fields.  The structure of the paper is as follows. In Section~\\ref{sec:models} we introduce and discuss the numerical method to calculate the dynamical cloud evolution and the suite of models included in the current analysis. In Section~\\ref{sec:prestellar} we present results on the angular momentum distribution of starless molecular cloud cores. We call them prestellar cores. Some of them collapse to become protostellar cores. In Section~\\ref{sec:protostellar} we investigate the angular momentum evolution of their central protostellar objects. We report a statistical correlation between specific angular momentum and mass, and analyze its dependence on the turbulent environment. The angular momentum vectors of neighboring protostars tend to be aligned, at least in the early accretion phase. This is discussed in Section~\\ref{sec:corr}. Finally, in Section~\\ref{sec:summary}, we summarize and conclude. ",
        "conclusions": "\\label{sec:summary}  We studied the rotational properties and time evolution of the specific angular momentum of prestellar and protostellar cores formed from gravoturbulent fragmentation in numerical models of supersonically turbulent, self-gravitating molecular clouds. We considered rms Mach numbers ranging from 2 to 10, and turbulence that is driven on small, intermediate, and large scales, as well as one model of collapse from Gaussian density fluctuations without any turbulence. Our sample thus covers a wide range of properties observed in Galactic star-forming regions, however, our main focus lies in typical low- to intermediate-mass star-forming regions like $\\rho$-Ophiuchus or Taurus.  With the appropriate physical scaling, we find the specific angular momentum~$j$ of prestellar cores in our models, i.e.\\ cloud cores as yet without central protostar, to be on average $\\langle j \\rangle = 7\\times10^{20}\\,\\mathrm{cm^2\\,s^{-1}}$. This agrees remarkably well with observations of cloud cores by \\citet{CAS02} or \\citet{GOO93}. Some prestellar cores go into collapse to build up stars and stellar systems. The resulting protostellar objects have on average $\\langle j \\rangle = 8\\times10^{19}\\,\\mathrm{cm^2\\,s^{-1}}$. This is one order of magnitude less, and falls into the range observed in G-dwarf binaries \\citep{DUQ91b}. Collapse induced by gravoturbulent fragmentation is accompanied by a substantial loss of specific angular momentum. This is mostly due to gravitational torques exerted by the ambient turbulent flow and to close encounters occurring when the protostars are embedded in dense clusters. This eases the ``angular momentum problem'' in star formation without invoking the presence of strong magnetic fields.  The time evolution of $j$ is intimately connected to the mass accretion history of a protostellar core. As interstellar turbulence and mutual interaction in dense clusters are highly stochastic processes, the mass growth of individual protostars is unpredictable and can be very complex.  In addition, a collapsing cloud core can fragment further into a binary or higher-order multiple or evolve into a protostar with a stable accretion disk. It is the ratio of rotational to gravitational energy~$\\beta$ that determines which route the object will take.  This is seen in the turbulent cloud cores studied here as well as in simulations of isolated cores where magnetic fields are important \\citep[e.g.,][]{BOS99}. The $\\beta$-distribution resulting from gravoturbulent cloud fragmentation reported here agrees well with $\\beta$-values derived from observations \\citep{GOO93}. The average value is $\\beta\\approx0.05$. It is important to note, that we find that the distribution of $\\beta$ stays essentially the same during collapse and accretion \\citep[see also][]{BUR00,GOO93}.  Although the accretion history and thus the evolution of the specific angular momentum of a single protostellar object is complex, we find a clear correlation between $j$ and mass $M$. This can be interpreted conveniently assuming collapse of an initially uniform density sphere in solid body rotation. Our models of gravoturbulent cloud fragmentation are best represented by the relation $j\\propto M^{2/3}$.  When prestellar cores form by compression as part of supersonically turbulent flows and then go into collapse and possibly break apart into several fragments due to the continuing perturbation by their turbulent environment, then we expect neighboring protostars to have similarly oriented angular momentum, at least during their early phases of accretion.  Star clusters form hierarchically structured, with  several young stellar objects being embedded in the same clump of molecular cloud material.  These protostars accrete from one common reservoir of gas and consequently gain similar specific angular momentum. Their disks and protostellar outflows therefore will closely align. Indeed, there are several examples of parallel disks and outflows seen in low-mass, isolated Bok globules \\citep{FRO03,KAM03,NIS01,SAI95}. During later phases of cluster formation, the initial substructure becomes erased by dynamical effects and the correlation between the angular momenta of neighboring protostars vanishes. This is in agreement with our numerical calculations of gravoturbulent cloud fragmentation. They show small groups of close protostellar objects that have almost aligned specific angular momenta. As expected, the alignment occurs during the early phase of accretion as neighboring protostars accrete material from the same region with similar angular momentum. During the subsequent evolution the correlation length decreases. This is either because protostellar aggregates disperse, or because infalling new material with different angular momentum becomes distributed unevenly among the protostars.  Altogether, the process of gravoturbulent fragmentation, i.e.\\ the interplay between supersonic turbulence and self-gravity of the interstellar gas, constitute an attractive base for a unified theory of star formation that is able to explain and reproduce many of the observed features in Galactic star forming regions \\citep{MAC04}. Our current study contributes with a detailed analysis of the angular momentum evolution during collapse."
    },
    "0402/astro-ph0402157_arXiv.txt": {
        "abstract": "We show how an adequate post--Newtonian generalization can be obtained for Newtonian dark matter halos associated with an empiric density profile. Applying this approach to halos that follow from the well known numerical simulations of Navarro, Frenk and White (NFW),  we derive all dynamical variables and show that NFW halos approximatelly follow an ideal gas type of equation of state which fits very well to a polytropic relation in the region outside the core. This fact suggests that ``outer'' regions of NFW halos might be related to equilibrium states in the non--extensive Statistical Mechanics formalism proposed by Tsallis. ",
        "introduction": "The issue of dark matter clustering in halos of virialized galactic structures is one of the most interesting open problems in astrophysics and cosmology \\cite{KoTu,Padma1,Peac,Ellis,Fornengo}. The physical properties of this dark matter are uncertain, leading to various proposed physical matter models: thermal sources, meaning a colissionless gas of weakly interacting massive particles (WIMP's), which can be very massive ($m\\sim 100-200$ GeV) supersymmetric~\\cite{Ellis} (``cold dark matter'' CDM) or self--interacting less massive ($m\\sim$ KeV) particles~\\cite{scdm,wdm} (``warm'' DM). Other proposals include scalar fields (real and complex)~\\cite{sfe_1,sfe_2}, global momopoles~\\cite{sfe_3}, axions, etc. However, all these models must comply from inferred direct and indirect observations that reveal the presence of DM: velocity profiles of rotating stars, microlensing and tidal effects affecting satelite galaxies and galaxies within galactic clusters. Galactic DM is mixed with visible baryonic matter (stars and gas) clustering in galactic disks, making up about 5--10 \\% of the total galactic mass. Hence, it is a good approximation to identify the gravitational field of a galaxy with that of its DM halo, considering visible matter as ``test particles'' in this field~\\cite{HGDM, Lake}.  Realistic galactic halos are obviously not spherically symmetric, but they are approximatelly so, since their global rotation is not dynamically significant~\\cite{urc}. Hence, we will consider throughout this paper that halos are spherically symmetric equilibrium configurations. Assuming the CDM paradigm and spherical symmetry, we can distinguish two types of halo models: those obtained from a Kinetic Theory approach, whether based on specific theoretical considerations or on convenient ansatzes that fix a distribution function satisfying Vlassov's equation~\\cite{BT}, or those emerging from ``universal'' density profiles obtained empirically by N--body numerical simmulations~\\cite{nbody_1,nbody_2,nbody_3}. In this paper we will consider the latter apprach, based on the well known numerical simulations of Navarro, Frenk and White (NFW)~\\cite{nbody_1,LoMa}. Although these simulations yield virialized equilibrium structures that reasonably fit CDM structures at a cosmological scale ($\\agt 100$ Mpc), some of their predictions in smaller scales (``cuspy'' density profiles and excess substructure) seem to be at odds with observations~\\cite{cdm_problems_1,cdm_problems_2}, especially those based on galaxies with low surface brightness (LSB), which are supposed to be overwhelmingly dominated by DM and so well suited to examine the predictions of various DM models.  Galactic halos are newtonian systems characterized by typical velocities, ranging from $5-10$ km/sec for dwarf galaxies up to about $1000-3000$ km/sec for rich clusters. However, we believe that a study of these systems under General Relativity, as a post--Newtonian approximation, might provide new information that can be useful and interesting for gravitational lensing and for the study of the interplay between cosmological scale evolution and galactic DM. In any case, since General Relativity is the best available theory of classical gravity, it is relevant from a theoretical point of view to be able to construct  spacetimes that are suitable for important self--gravitating structures like galactic halos.  All DM halo models derive a full set of dynamical variables from a given ``mass--density'' profile. In a post--newtonian relativistic generalization we will assume that this density is the dominant rest--mass contribution to the matter--energy density, made up by rest--mass and an internal energy term proportional to suitably defined temperature and pressure (of a kinetic nature). Thus, we will assume an ``ideal gas'' type of equation of state~\\cite{RKT,Padma2,Padma3} in which this internal energy density becomes determined only by the hydrostatic equations themselves in the case of isotropic velocity distributions. In the anistropic case, which we leave for a future paper~\\cite{enproceso}, various empirical ansatzes can be assumed in order to relate radial and angular components of the stress tensor. ",
        "conclusions": "The fact that NFW halos asymptoticaly comply with a polytropic relation with $n\\approx 5.5$ is quite significant, since stellar polytropes characterized by (\\ref{poly}) are the equilibrium state associated with the entropy functional in the non--extensive entropy formalism derived by Tsallis and coworkers~\\cite{PL,Tsallis,TS1, TS2}. In its application to self--gravitating collisionless systems this formalism is characterized by the free parameter $q=(2n-1)/(2n-3)$, so that the isothermal sphere (equilibrium state for the usual Boltzmann--Gibbs entropy functional) follows in the ``extensive entropy'' limit $n\\to\\infty$ (or $q\\to 1$). Assuming Tsallis theory to be correct, the empiric verification (see Figure \\ref{fig:1}) that NFW halos outside the ``inner'' core satisfy a polytropic relation might indicate that in this ``outer'' region the NFW numerical simulations yield self--gravitating configurations that approach an equilibrium state characterized by the Tsallis parameter $q\\approx 1.25$. However, while the central cusps in the density profile that are predicted by NFW simulations seem to be at odds with observations~\\cite{cdm_problems_1,cdm_problems_2}, there is no conflict between these observations and the $1/x^3$ scaling of the NFW density profile outside the core region (as well as the rotation velocity profile from (\\ref{VrotPN})). Although the issue of the cuspy cores is still controversial, if galactic halos seem to exhibit flat density cores, their profiles could be adjusted to stellar polytropes and this might be helpful in providing a better empirical verification of Tsallis' formalism. However, this idea must be handled with due case, since stellar polytropes follow from an isotropic velocity distribution, while galactic halos with such distributions could be unrealistic.  As pointed out before, the density profile of NFW halos diverges at the center. Apparently this issue has not bothered astrophysicists, since (as mentioned before) the cuspy cores of NFW numerical simulations are meant to show a density scaling of $1/x$ near the center and these simulations cannot resolve distances to the halo center smaller than 1\\,\\% of the virial radius~\\cite{NSres}. One way to deal with this unphysical feature, leading to a better description of these halos, would be to perform a smooth matching between NFW spacetimes and a small central region with a regular density profile. An adequate radius for this ``inner'' region could be the minimal resolution scale in numerical simulations ($x=x_0\\sim 0.01\\, c_0$). Another improvement could be a smooth matching of the NFW spacetime to a Schwarzschild vacuum exterior at the virial radius $x=c_0$, which is the physical radius of the halo. One of the matching conditions in this latter case would be $\\CP(c_0)=0$, implying a different choice of the integration constant $C$ in (\\ref{eq:CP1}). Another necessary improvement is the study of the anisotropic cases for which $\\alpha\\ne 0$.  We have constructed the spacetime corresponding to post--Newtonian generalizations suitable to NFW halos. Although we have presented only the idealized case with isotropic pressure, the methodology that we followed here can be applied, in principle, to any Newtonian model of galactic halos. We believe that it is necessary to study galactic halo models (NFW, as well as other empiric or theoretical models) within a wider framework including the usual thermodynamics of self--gravitation systems~\\cite{Padma2,Padma3}, as well as alternative approaches such as Tsallis' formalism~\\cite{Tsallis,PL,TS1,TS2}. Such an improvement and extension of the present study of NFW halos are being pursued elsewhere~\\cite{enproceso}."
    },
    "0402/astro-ph0402011_arXiv.txt": {
        "abstract": "The giant molecular cloud complex, Sagittarius B2 (Sgr B2), was observed with the ISO Long Wavelength Spectrometer over the range 47--197~$\\umu$m. The ground state rotational transitions of OH lie in the survey range and we have analysed results from the $^{16}$OH, $^{17}$OH and $^{18}$OH isotopomers. Absorption in these lines due to clouds along the line of sight towards Sgr B2 has allowed us to compare the isotopomer abundance at different galactocentric radii. This is an excellent test of previous results for the isotope ratios which were measured at radio wavelengths and generally required a double ratio with $^{12}$C/$^{13}$C. ",
        "introduction": "Sgr B2 was observed over the wavelength range 47--197~$\\umu$m in a full and unbiased spectral survey using the Infrared Space Observatory Long Wavelength Spectrometer (LWS). A spectral resolution of 30--40~km~s$^{-1}$ was achieved across this range using the instrument in its Fabry-P\\'{e}rot mode, L03. Lines observed in the survey include rotational transitions of molecules (e.g. OH, H$_{2}$O, H$_{3}$O$^{+}$, NH, NH$_{2}$, NH$_{3}$, CH) as well as atomic fine structure lines (e.g. [OI], [OIII], [CII], [NII], [NIII]). All of the low lying rotational transitions of OH were covered by the survey.  Sgr~B2 is a giant molecular cloud complex, located close to the Galactic Centre. It emits a strong far-infrared (FIR) continuum spectrum with a peak near 80~$\\umu$m \\cite{goicoechea_c}. This is a perfect background against which to observe line of sight absorption features.  We detected several $^{16}$OH transitions as well as one $^{18}$OH transition. Each line shows two resolved $\\varLambda$-doublet components. Figure~\\ref{obs_lev} shows the detected rotational transitions. Transitions from the lowest level show absorption due to the whole line of sight in the range $-107$~km~s$^{-1}$ to $+30$~km~s$^{-1}$. This is due to galactic spiral arm clouds located between the Sun and Galactic Centre \\cite{greaves94}. Higher transitions are only seen at the velocity of Sgr B2 itself ($\\sim $+65~km~s$^{-1}$; \\cite{martin-pintado}) and these have been modelled by \\cite{goicoechea_b}.  \\begin{figure}[!t] \\centering \\includegraphics[height=6cm]{18oh_levels.eps} \\caption{LWS Fabry-P\\'{e}rot observations of $^{16}$OH (left) and $^{18}$OH (right) shown with their rotational levels. The $J$ value of each level is indicated. The two resolved $\\varLambda$-components have been averaged together for each transition except at 79~$\\umu$m.} \\label{obs_lev} \\end{figure} ",
        "conclusions": ""
    },
    "0402/astro-ph0402294_arXiv.txt": {
        "abstract": "{We report the discovery of narrow  \\fetfive\\ and \\fetsix\\ \\ka\\ X-ray absorption lines at 6.65$\\,^{+0.05}_{-0.02}$ and 6.95$\\,^{+0.05}_{-0.04}$~keV in the persistent emission of the dipping low-mass X-ray binary (LMXB) \\nineteen\\ during an XMM-Newton observation performed in September 2002.  In addition, there is marginal evidence for absorption features at 1.48~keV, 2.67~kev, 7.82~keV and 8.29~keV consistent with \\ion{Mg}{xii}, \\ssixteen, \\nitseven\\ \\ka\\ and \\fetsix\\ \\kb\\ transitions, respectively.  Such absorption lines from highly ionized ions are now observed in a number of high inclination (ie. close to edge-on) LMXBs, such as \\nineteen, where the inclination is estimated to be between 60--80\\degree.  This, together with the lack of any orbital phase dependence of the features (except during dips), suggests that the highly ionized plasma responsible for the absorption lines is located in a cylindrical geometry around the compact object. Using the ratio of \\fetfive\\ and \\fetsix\\ column densities, we estimate the photo-ionization parameter of the absorbing material, $\\xi$, to be $10^{3.92}$ \\xiunit. Only the \\fetfive\\ line is observed during dipping intervals and the upper-limits to the \\fetsix\\ column density are consistent with a decrease in the amount of ionization during dipping intervals. This implies the presence of cooler material in the line of sight during dipping.  We also report the discovery of a 0.98~keV absorption edge in the persistent emission spectrum.  The edge energy decreases to 0.87~keV during deep dipping intervals. The detected feature may result from edges of moderately ionized Ne and/or Fe with the average ionization level decreasing from persistent emission to deep dipping. This is again consistent with the presence of cooler material in the line of sight during dipping. ",
        "introduction": "\\label{sec:intro}  \\nineteen\\ is a LMXB showing periodic intensity dips in its light curve \\citep{1916:walter82apjl,1916:white82apjl}. Dips are believed to be due to obscuration of the central X-ray source by a vertical structure located at the outer edge of the accretion disk and resulting from the impact of the accretion flow from the companion star into the disk \\citep{1916:white82apjl}.  The presence of dips in \\nineteen\\ and the lack of X-ray eclipses from the companion star indicate that the system is viewed relatively close to edge-on, at an inclination angle in the range $\\sim$60--80\\degree\\ \\citep{frank87aa,1916:smale88mnras}.  The X-ray dip period is 50 minutes \\citep[e.g., ][]{1916:white82apjl}, the shortest amongst the dipping sources.  The optical counterpart of \\nineteen\\ shows a modulation with a period $\\sim$1~\\% longer than the X-ray dip period \\citep[e.g., ][]{1916:callanan95pasj}. This discrepancy has led to several interpretations including superhumps \\citep{1916:schmidtke88aj} and a hierarchical triple system model \\citep{1916:grindlay88apjl}. \\citet{1916:retter02mnras} recently favored the superhump model, which invokes a precessing accretion disk, and which identifies the X-ray period as orbital. \\nineteen\\ is a  type~I X-ray burster \\citep{1916:becker77apjl}, indicating that the compact object is a neutron star. \\nineteen\\ has shown quasi-periodic oscillations  at various frequencies ranging from $\\sim $0.2 to 1300 Hz \\citep{1916:boirin00aa}, and a 270~Hz highly coherent oscillation during an X-ray burst \\citep{1916:galloway01apjl}.  \\nineteen, as most dippers, shows spectral hardening during dipping. However, the spectral evolution is not consistent with a simple increase of photo-electric absorption by cool material, as an excess of photons is present at low energy. Two approaches have been proposed for modelling the spectral evolution during dips. In the ``absorbed plus unabsorbed'' approach \\citep[e.g., ][]{0748:parmar86apj}, the persistent (non-dipping) model is used to fit the intensity-selected dip spectra, but is divided into two parts. One part is allowed to be absorbed, whereas the other one is not. The spectral evolution during dipping is well accounted for by a large increase in the column density of the absorbed component, and a decrease of the normalization of the unabsorbed component. The latter decrease has been attributed to effects of electron scattering in the absorber.  In the ``progressive covering'', or ``two-component'' or ``complex continuum'' approach \\citep[e.g., ][]{1916:church97apj}, the X-ray emission originates from two components. The first one, generally modelled as a blackbody, is from a point-like region, such as a boundary layer around the compact object.  The second component, generally modelled as a power-law, comes from an extended region, such as an accretion disk corona. The complex continuum approach allows both the persistent and the various intensity-selected dipping spectra to be modelled with the same two components, and explains the spectral changes by allowing partial and progressive covering of the extended source by an opaque absorber that occults various fractions of the source. The absorption applied to the point-source component is allowed to vary independently of the absorption applied to the extended component, but no partial covering is needed as the source is supposed to be point-like and thus fully covered by the absorber in the line of sight during the dips. Both approaches have been applied to \\nineteen\\ \\citep[e.g., ][ respectively]{1916:yoshida95pasj,1916:church97apj}.   In addition to the continuum emission, a broad emission feature interpreted as fluorescent line emission from neutral Fe was detected from \\nineteen, at 5.60$^{+0.53}_{-0.43}$~keV \\citep{1916:smale92apj}, 6.14$^{+0.18}_{-1.07}$~keV \\citep{1916:bloser00apj} and 5.9$^{+0.2}_{-0.1}$~keV \\citep{asai00apjs}.  The low energy of the emission lines compared to the 6.4~keV expected for neutral Fe led \\citet{1624:parmar02aa} to propose that the broad emission feature was modulated with $\\sim$7~keV absorption features, which were not included in the spectral model, reducing the measured energy of the emission feature. Narrow absorption features from highly ionized Fe and other metals are now observed in a growing number of LMXBs (see Table~\\ref{tab:sources}), thanks to the improved sensitivity and spectral resolution offered by the new generation of instruments on-board of {\\it Chandra} and XMM-Newton.  These discoveries indicate that a highly ionized plasma is present in these systems, that was so far not taken into account in models. Therefore, the study of these lines appears to be extremely important to characterize the geometry and physical properties of this plasma, that could be a common property of accreting systems. The presence of absorption lines could be related to the viewing angle of the system.  In this paper, we report the discovery of narrow \\fetfive\\ and \\fetsix\\ \\ka\\ X-ray absorption lines near 7~keV in the dipping LMXB \\nineteen. Absorption features due to highly ionized Mg, S and Ni may also be present. The \\fetfive\\ line is also observed during dipping intervals.  We also report the discovery of an absorption edge at an energy of 0.98~keV and 0.87~keV during the persistent and dipping intervals, respectively.  Detailed modelling of the continuum emission, comparison of the ``absorbed plus unabsorbed'' and ``complex continuum'' approaches, and discussion of their physical interpretations are the subject of a separate paper (Webb et al., in preparation), and are therefore not presented here. In this paper, we focus on the X-ray absorption features, and, in order to extract their properties, we fit the continuum with a simple model and follow the complex continuum approach. ",
        "conclusions": "We have reported the discovery of narrow absorption lines at 6.65 and 6.95~keV, consistent with resonant \\fetfive\\ and \\fetsix\\ \\ka\\ transitions, in the persistent emission of the LMXB \\nineteen. There is also marginal evidence for absorption features consistent with \\ion{Mg}{xii}, \\ssixteen, \\nitseven\\ \\ka\\ and \\fetsix\\ \\kb\\ transitions. Such absorption lines from highly ionized ions are now observed in a number of LMXBs seen close to edge-on ($i\\sim70$\\degree), such as \\nineteen.  This suggests that the highly ionized plasma responsible for the absorption lines is related to the accretion disk.  We have reported the dectection of the \\fetfive\\ line in the dipping emission  of \\nineteen, and the upper-limits to the \\fetsix\\ column densities are consistent with a decrease in the amount of ionization during dipping intervals.  This implies the presence of cooler material in the line of sight during dipping.  We have also reported the discovery of a 0.98~keV absorption edge in the persistent emission spectrum.  The edge energy decreases to 0.87~keV during deep dipping intervals. The detected feature may result from edges of moderately ionized Ne and/or Fe with the average ionization level decreasing from persistent emission to deep dipping. This is again consistent with the presence of cooler material in the line of sight during dipping."
    },
    "0402/astro-ph0402541_arXiv.txt": {
        "abstract": "We investigate the effects of thermal interactions on tracking models of quintessence. We show that even Planck-suppressed interactions between matter and the quintessence field can alter its evolution qualitatively. The dark energy equation of state is in many cases strongly affected by matter couplings. We obtain a bound on the coupling between quintessence and relativistic relic particles such as the photon or neutrino. ",
        "introduction": "Recent evidence \\cite{observation} suggests that a large fraction of the energy density of the universe has negative pressure, or equation of state with $ w \\equiv p /\\rho < 0$. One candidate source of this dark energy is a slowly varying and spatially homogeneous scalar field called quintessence \\cite{darkenergy}. Because the dark energy redshifts more slowly than ordinary matter or radiation, it appears that the ratio of energy density of quintessence to that in ordinary particles must be fine tuned to a specific infinitesimal value in the early universe in order to explain its current observed value. One class of models that ameliorate this problem describe {\\it tracker fields} \\cite{tracker} whose evolution is largely insensitive to initial conditions and at late times begin to dominate the energy density of the universe with a negative equation of state. Tracker models have difficulty producing $w_\\phi$ consistent with observational data: they generally imply $w^{\\rm eff}_\\phi \\gsim -0.7$ , whereas WMAP implies $w_{\\phi}^{\\rm eff} < -0.78$ (95\\% CL) \\cite{observation}. (Here, effective means as measured observationally, so integrating over redshifts less than of order $10^3$.) Nevertheless, they provide an interesting class of models describing dark energy as a slowly evolving scalar field. An alternative class of models, which avoids the extremely flat potentials required at late times of tracker models, utilizes nonlinear field oscillations that exhibit $w < 0$ \\cite{nonlinear}.  Tracker models generally require only a single adjustable parameter. Once this parameter is appropriately chosen, a wide range of initial values of the tracker field, $\\phi$, and its derivative, $\\dot{\\phi}$, result in similar values of its energy density today. This is due to an attractor-like property of the tracker equations of motion.  In this paper we investigate the effects of interactions between the quintessence field, $\\phi$, and ordinary matter particles in the early universe. It is important to note that while the zero temperature potential may be fine tuned in order for the evolution of $\\phi$ to have the attractive properties mentioned above, the same may not be done with the finite temperature effective potential. That is to say, once the form of the renormalized zero temperature potential is determined, no additional freedom remains to fine-tune away unwanted thermal effects. Therefore, such effects must be considered. We expect thermal interactions to be at least of gravitational strength (even in the case where $\\phi$ is a ``hidden sector'' field). At minimum, quantum gravity is likely to produce interactions of the type \\cite{Planck} \\beq \\label{int} {\\beta_i \\over M_P} \\, \\phi \\, {\\cal L}_i ~~,\\eeq where $M_P$ is the Planck scale, and ${\\cal L}_i$ are terms in the standard model lagrangian, including for example $$ F_{\\mu \\nu}^2~,~F_{\\mu \\nu} \\tilde{F}^{\\mu \\nu} ~,~ \\bar{\\psi} \\DS \\, \\psi ~,~ \\cdots $$ where $F$ is the field strength of any gauge field (including the photon, but not excluding gluons or the W or Z) and $\\psi$ is any fermion field from neutrinos to the top quark. Even if $\\phi$ were a pseudo-Goldstone boson \\cite{constraints}, it would be surprising not to find at least Planck-suppressed violations of the resulting $~\\phi \\rightarrow \\phi + {\\rm constant}~$ symmetry. String theory, for example, is believed to not exhibit any exact global symmetries \\cite{strings}. Previous constraints on certain $\\beta_i$ are quite strong, where the coupling is to the photon or gluon \\cite{constraints}. However, some $\\beta_i$ could be much larger, such as when the interaction is with the W, Z or even a neutrino\\footnote{Direct coupling to relic neutrinos has been considered previously \\cite{neutrino}.}.   In the early universe matter particles are in thermal equilibrium, and the interactions in (\\ref{int}) produce a thermal mass for $\\phi$ of the form \\beq  \\label{mass} \\left( {\\beta_i \\over M_P} \\right)^2 \\, \\phi^2 \\, T^4~~, \\eeq where $T$ is the temperature. If the thermal degree of freedom is massive, the thermal effect goes to zero exponentially as $e^{- m /T}$ when the temperature drops below the mass $m$. We note that the thermal effects of interest here are {\\it in addition} to any quantum corrections to the effective potential resulting from the interactions between matter and quintessence (see, for example, \\cite{qloop}). In general, the bare parameters of quintessence models must be fine-tuned in order to obtain potentials of the necessary type. We assume here that this fine-tuning is achieved (whatever its consequences for the plausibility of the model) and focus on thermal effects which must also arise.     Although $\\phi$ may not itself be in equilibrium, nevertheless its dynamical evolution will be affected by these thermal interactions, just as for the axion field near the QCD phase transition \\cite{axion}. We can derive the correction (\\ref{mass}) to the effective potential for $\\phi$ as follows. Let the cold $\\phi$ field be a static, external source for a Euclidean path integral describing the thermal degrees of freedom. The timelike boundary conditions for the path integral have period given by the inverse temperature. Performing the integration over the thermal fields yields a contribution to the effective potential for $\\phi$, and the usual perturbative analysis identifies the leading effect to be the thermal mass term in (\\ref{mass}). In this calculation, we need never assume that $\\phi$ itself is in thermal equilibrium, yet its effective potential receives temperature-dependent contributions.   In what follows we will examine how the contribution of (\\ref{mass}) to the tracker potential modifies its evolution. We can make a simple argument for why (\\ref{mass}) is non-negligible at late times. At late times, the quintessence field must have a very small mass: $V''(\\phi)^{1/2} \\lsim H_0 \\sim 10^{-33}$ eV, and contribute of order closure density to $\\Omega$: $V(\\phi) \\sim (10^{-3} {\\rm eV})^4$, which implies that $\\phi \\sim M_P$. This means that the mass term in (\\ref{mass}) can be roughly the same size as $V(\\phi)$, up to powers of $\\beta_i$.   At early times, (\\ref{mass}) also affects the evolution in many cases. Suppose the tracker potential is given by $V(\\phi) = M^{l+4} \\phi^{-l}$, where $l > 4$. Then $V(\\phi_*)$ and (\\ref{mass}) are comparable at the minimum of the combined potential: \\beq \\phi_* \\sim M \\left( {M^2 M_P^2 \\over \\beta^2 T^4} \\right)^{1 \\over l+2} ~~,\\eeq where the potential energy density is roughly \\beq V (\\phi_*) \\equiv V_* \\sim M^4 \\left( { \\beta^2 T^4 \\over M^2 M_P^2} \\right)^{l \\over l+2} ~~.\\eeq In many cases, $\\phi$ oscillates about the temperature-dependent minimum $\\phi_*$. The oscillation energy redshifts faster than the potential energy at the minimum, $V_* \\sim T^{4l / (l+2)}$, so $\\phi$ simply tracks $\\phi_*$ with oscillations that decrease in amplitude over time. Interestingly, $V_*$ redshifts exactly as the energy density of the tracker solution \\cite{pr1988} (assuming radiation domination; during a matter dominated epoch $V_*$ redshifts somewhat faster than the usual tracker energy density). This means that thermal effects will keep $\\phi$ and its energy density near their desired values, even though the physics responsible is very different. When the thermal term eventually either disappears due to the crossing of a particle mass threshold, or becomes negligible due to redshift, $\\phi$ will merge back to a tracker solution. ",
        "conclusions": "Our analysis shows that the evolution of the tracker field depends quite sensitively on its interaction with thermal matter, even when the strength of the couplings is as small as one would imagine they may possibly be; i.e., Planck-suppressed. When $\\beta$ is larger than of order unity, there is a tendency for the evolution to be controlled by that of $\\phi_*$, once initial oscillations have damped away. By a lucky coincidence, the redshift of $V_* \\sim T^{4l /(l+2)}$ is the same as that of the tracker solution (during radiation domination), so that $\\phi$ can rejoin a tracker solution at late times. The most dangerous possibility (which is realized when $\\beta \\gsim \\beta_c$) is that $\\phi$ is still following $\\phi_*$ at late times, in which case its equation of state will be far from the observationally favored $w_\\phi = -1$. For general $l$, the equation of state obeyed by $V_*$ is $w_{\\phi *} = (l-6)/ 3(l+2)$, which is never consistent with observational bounds for $l > 4$. We have checked that the behavior described above is qualitatively similar when higher order terms such as \\beq \\label{cubic} \\left( \\beta \\over M_P \\right)^3 \\phi^3 T^4 \\eeq are included in the potential.  Hence, we conclude that there are stringent limits on the coupling between the tracker field and any particles which are still relativistic today, such as the photon or neutrinos. Such limits cannot be avoided through fine tuning of the finite temperature effective potential; once the zero temperature potential has been computed the finite temperature effects are determined. Interactions which are more than roughly two orders of magnitude stronger than Planck-suppressed lead to a problematic equation of state. Couplings of the tracker to heavy particles, which freeze out at $T \\sim m$, may alter the tracker evolution at early times, but do not affect the observed dark energy equation of state and are hence poorly constrained. The best hope of directly detecting the quintessence field may be through its interaction with massive particles.  Finally, although our analysis has focused on tracker models, similar results apply for any quintessence model in which the field is today slowly evolving in a very flat potential. As we argued in the introduction, in {\\it any} such model the value of $\\phi$ must be of order $M_P$ today, which means that thermal terms such as (\\ref{mass}) or (\\ref{cubic}) will be important for sufficiently large $\\beta$. Large couplings to relic particles such as neutrinos can be ruled out as they lead to a problematic equation of state."
    },
    "0402/astro-ph0402588_arXiv.txt": {
        "abstract": "The annihilation of neutralino dark matter in the Galactic Center (GC) may result in radio signals that can be used to detect or constrain the dark matter halo density profile or dark matter particle properties. At the Galactic Center, the accretion flow onto the central Black Hole (BH) sustains strong magnetic fields that can induce synchrotron emission by electrons and positrons generated in neutralino annihilations during advection onto the BH. Here we reanalyze the radiative processes relevant for the neutralino annihilation signal at the GC, with realistic assumptions about the accretion flow and its magnetic properties. We find that neglecting these effects, as done in previous papers, leads to the incorrent electron and photon spectra. We find that the magnetic fields associated with the flow are significantly stronger than previously estimated. We derive the appropriate equilibrium distribution of electrons and positron and the resulting radiation, considering adiabatic compression in the accretion flow, inverse Compton scattering off synchrotron photons (synchrotron self-Compton scattering), and synchrotron self-absorption of the emitted radiation. We derive the signal for a Navarro-Frenk-White (NFW) dark matter halo profile and a NFW profile with a dark matter spike due to the central BH. We find that the observed radio emission from the GC is inconsistent with the scenario in which a spiky distribution of neutralinos is present. We discuss several important differences between our calculations and those previously presented in the literature. ",
        "introduction": "The combination of cosmic microwave background anisotropy studies with observations of distant type Ia supernovae and measurements of the large scale structure of the universe reveal that the energy density in our universe is dominated  by dark energy ($\\sim 70\\%$) followed by a significant contribution from matter ($\\sim 27\\%$). In addition, these studies show that the baryonic content of the universe is limited to less than $\\sim 4 \\%$ in agreement with big bang nucleosynthesis predictions. More specifically, recent WMAP results give the mean baryonic density in the universe to be within 0.044 $\\pm 0.004$ of the critical density while the matter density is 0.27 $\\pm 0.04$ of critical \\cite{WMAP}. The remaining $\\sim$ 0.23 of the matter density, i.e., the bulk of the matter in the universe, is believed to be a yet undiscovered form of dark matter.  Weakly interacting massive particles (WIMPs) are natural candidates for the dark matter \\cite{kamion}. Particles with masses around $\\sim 100$ GeV that interact only weakly have freeze-out densities in the required range. In addition, particle physics models that invoke supersymmetry generate a number of plausible WIMP candidates. In supersymmetric extensions of the standard model, the lightest supersymmetric particle may be stable due to conservation of R-parity enabling their survival to the present. In addition to massive and weakly interacting, dark matter particles are expected to be neutral. A class of neutral lightest supersymmetric particles is a combination of gauginos and higgsinos, named the neutralino often represented by $\\chi$.  A number of experiments are presently searching for the neutralino (see, e.g., \\cite{bergst}). Accelerator  experiments have placed lower limits on the neutralino mass of $m_{\\chi} \\ga 37$ GeV  \\cite{accel}.  The neutralino mass in the minimal supersymmetric extension of  the standard model that can explain the dark matter density is also bounded from above, $m_{\\chi} \\la 7$ TeV. In some constrained supersymmetric models $m_{\\chi} \\la 500$ GeV due to the  cosmological constraints from WMAP \\cite{ellis}. Direct searches for neutralinos that cross the Earth are presently reaching the required sensitivity to probe the relevant range in  parameter space of supersymmetric models.  A good complement to accelerator and direct detection methods are indirect searches based on the emission from neutralino self-annihilation in astrophysical environments \\cite{bere1,bere2,berg1,gon1,calc,berg2,bert,ullio,blasi1,tasitol,cesar}. Neutralinos self-annihilate at a rate proportional to the square of the neutralino density. Thus, the highest density dark matter regions are the best candidates for indirect searches. In fact, the GC region may potentially be so dense that all neutralino models would be ruled out \\cite{gon1}. This strong constraint arises in models where the super-massive black hole (BH) at the GC induces a strong dark matter density peak called {\\it the spike}. The existence of such a spike is strongly dependent on the formation history of the Galactic Center BH \\cite{merrit}. If the BH formed adiabatically a spike would be present, while a history of major mergers would most likely not allow the survival of a spike.  In contrast to the uncertain presence of a central spike in the dark matter distribution, the central BH is known to induce an accretion flow of baryonic matter around its event horizon. The accretion flow carries magnetic fields, possibly amplified to near equipartition values due to the strong compression. The distribution of electrons and positrons (hereafter both called electrons) produced by neutralino annihilation at the GC would also be compressed toward the BH radiating through synchrotron and inverse Compton scattering off the local photon background.  By considering the injection of electrons, combined with radiative losses and adiabatic compression, we find the equilibrium spatial and spectral electron distribution and derive the expected radiation signal. We find that the synchrotron emission of electrons from neutralino annihilation ranges from radio and microwave energies up to the optical, in the central more magnetized region of the accretion flow. At low frequencies, synchrotron self-absorption (SSA) reduces the amount of  radiation transmitted outwards. The resulting signal is stronger than the observed emission in the $10$ to $10^{5}$ GHz range for the case of a spiky dark matter profile while for a pure NFW profile the emission is below the observed level. Here, we do not attempt  to explore all possible spike profiles in the GC region, that might originate from the adiabatic compression of different types of {\\it initial} dark matter profiles. Instead, we limit ourselves to the case of a spike arising from adiabatic compression of a NFW profile in the gravitational field of the massive BH at the GC (hereafter, {\\it the spike case}.) in contrast to the pure NFW profile.  Previous calculations of the type presented here were either incomplete or had unphysical values for some of the parameters. In particular, calculations presented in \\cite{gon1} and \\cite{bert} assume a magnetic field  that does not correspond to the equipartition field for  the accretion flow around the black hole. We show here that if this {\\it incorrect} field is adopted, then the main processes responsible for reaching the equilibrium of relativistic electrons are synchrotron self-compton scattering (SSC) and adiabatic compression of the particles advected with the accretion flow, processes which were both ignored in previous calculations. On the other hand, if the {\\it correct} magnetic field is adopted, the equilibrium distribution of the radiating electrons is determined by advection and synchrotron emission.  The paper is organized as follows: in \\S \\ref{sec:spatial} we describe the spatial distribution of dark matter in the GC region. In \\S \\ref{sec:accret} we describe the accretion flow around the BH and its magnetic field. The injection of electrons and positrons from neutralino annihilation is considered in \\S \\ref{sec:inj}. In \\S  \\ref{sec:equi} we solve the transport equation for electron-positron pairs from neutralino annihilation in the presence of energy losses and adiabatic compression. Synchrotron self-absorption is discussed in \\S \\ref{sec:ssa}. Results of our calculations and comparison with observations are shown in \\S \\ref{sec:results} and comparison with previous work is made in \\S \\ref{sec:prev}. We conclude in \\S \\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  In conclusion, we reaffirm that a spike density profile obtained from the adiabatic compression of a NFW profile in the gravitational field of the BH at the GC, using dark matter parameters typical of neutralinos, is ruled out by the radio and NIR observations from  SgA$^*$ while the NFW density profile remains a suitable hypothesis, not strongly constrained by available observations.  We reached this conclusion by a careful consideration of the accretion flow and the loss processes in the transport equations for the neutralino generated electrons and positrons as well as the radiative transfer. At the Galactic Center, the accretion flow onto the central BH sustains strong magnetic fields that induce synchrotron emission by electrons and positrons generated in neutralino annihilations during advection onto the BH. We found that the magnetic fields associated with the flow are significantly stronger than previously adopted. With these equipartition fields, we derive the appropriate equilibrium distribution of electrons and positrons and the resulting radiation considering adiabatic compression in the accretion flow, inverse Compton scattering off synchrotron photons (synchrotron self-Compton scattering), and synchrotron self-absorption of the emitted radiation. We calculate the signal for a NFW dark matter halo profile and a NFW profile with a dark matter spike due to the central BH. We found that the annihilation of neutralino dark matter in the GC results in radio signals that overwhelm observations if the spike is present in the dark matter density profile. However, observed emissions from the GC are consistent with neutralinos following a NFW profile at the Galactic Center."
    },
    "0402/astro-ph0402631_arXiv.txt": {
        "abstract": "Observations indicate that young massive star clusters in spiral and dwarf galaxies follow a relation between luminosity of the brightest young cluster and the star-formation rate (SFR) of the host galaxy, in the sense that higher SFRs lead to the formation of brighter clusters. Assuming that the empirical relation between maximum cluster luminosity and SFR reflects an underlying similar relation between maximum cluster mass ($M_{\\rm ecl,max}$) and SFR, we compare the resulting $SFR(M_{\\rm ecl,max})$ relation with different theoretical models. The empirical correlation is found to suggest that individual star clusters form on a free-fall time-scale with their pre-cluster molecular-cloud-core radii typically being a few~pc independent of mass. The cloud cores contract by factors of~5 to 10~while building-up the embedded cluster. A theoretical $SFR(M_{\\rm ecl,max})$ relation in very good agreement with the empirical correlation is obtained if the cluster mass function of a young population has a Salpeter exponent $\\beta \\approx 2.35$ and if this cluster population forms within a characteristic time-scale of a few-$10$~Myr.  This short time-scale can be understood if the inter-stellar medium is pressurised thus precipitating rapid local fragmentation and collapse on a galactic scale. Such triggered star formation on a galactic scale is observed to occur in interacting galaxies. With a global SFR of $3-5\\,M_\\odot$/yr the Milky Way appears to lie on the empirical $SFR(M_{\\rm ecl,max})$ relation, given the recent detections of very young clusters with masses near $10^5\\,M_\\odot$ in the Galactic disk. The observed properties of the stellar population of very massive young clusters suggests that there may exist a fundamental maximum cluster mass, $10^6 < M_{\\rm ecl,max*}/M_\\odot < 10^7$. ",
        "introduction": "In a series of publications Larsen \\citep{Lar00, Lar01, Lar02} and \\citet{LarRi00} examined star cluster populations of 37 spiral and dwarf galaxies and compared the derived properties with overall attributes of the host galaxy. For this work they used archive HST data, own observations and literature data. They showed that cluster luminosity functions (LFs) are very similar for a variety of galaxies. They also found that the V-band luminosity of the brightest cluster, $M_{\\rm V}$, correlates with the global star-formation rate, SFR, but it is unclear if this correlation is of physical or statistical nature.  According to the statistical explanation there is a larger probability of sampling more luminous clusters from a universal cluster LF when the SFR is higher \\citep{Lar02, BiHuEl02}.  \\citet{Lar01} concluded that all types of star clusters form according to a similar formation process which operates with different masses. Smaller clusters dissolve fast through dynamical effects (gas expulsion, stellar-dynamical heating, galactic tidal field) and only massive clusters survive for a significant fraction of a Hubble time \\citep{Ves98,FZ01,BaMa03}. The notion is that virtually all stars form in clusters \\citep{KB02,LaLa03}, and that a star-formation ``epoch'' produces a population of clusters ranging from about $5\\,M_{\\odot}$ (Taurus-Auriga-like pre-main sequence stellar groups) up to the heaviest star cluster which may have a mass approaching $10^6~M_\\odot$. The time-scale over which such a cluster population emerges within a galaxy defines its momentary SFR.  The aim of this contribution is to investigate if the empirical $M_V(SFR)$ relation may be understood to be a result of physical processes.  In \\S~\\ref{obs} the observational data concerning the correlation between the SFR and $M_V$ of the brightest star cluster are presented, and the empirical and physical models describing this correlation are elaborated in \\S~\\ref{equ}.  \\S~\\ref{condis} contains the discussion and conclusion. ",
        "conclusions": "\\label{condis}  Observations of young star-cluster systems in disk galaxies show that there exists a correlation between the total SFR and the luminosity of the brightest star-cluster in the young-cluster population. This can be transformed to a SFR--heaviest-cluster-mass relation ($SFR(M_{\\rm ecl,max})$, eq.~\\ref{empSFR}).  Very young star-clusters in the MW and the LMC that are deduced to have formed within a few~Myr follow a similar $SFR(M_{\\rm ecl,max})$ relation, although this ``local'' relation is somewhat steeper if it is assumed that the formation time-scale of individual clusters is the same in all cases ($\\approx 1$~Myr, Fig.~\\ref{fig4b}). Taking instead the formation-time-scale to be the free-fall time of the pre-cluster molecular cloud core the correct slope is obtained if the pre-cluster cloud core radius is a few~pc independent of cluster mass (Fig.~\\ref{fig3b}).  This implies that the cluster-forming molecular cloud cores may contract by a factor of~5 to~10 as the clusters form. That the pre-cluster radii appear to not vary much with cluster mass implies the pre-cluster cores to have increasing density with increasing mass. Indeed, \\citet{Lar03} finds young extra-galactic clusters to have only a mild increase of effective radius with mass, and embedded clusters from the local Milky Way also suggest the cluster radii to be approximately independent of cluster mass \\citep{Krou02, KB02}.  A model according to which the total mass of the young-cluster population, $M_{\\rm tot}$, is assumed to be assembled in a star-formation ``epoch'' with an a-priori unknown duration, $\\delta t$, gives the corresponding $SFR=M_{\\rm tot}/\\delta t$ and leads to good agreement with the empirical $SFR(M_{\\rm ecl,max})$ relation for $1\\simless \\delta t/{\\rm Myr}\\simless 10$. A particularly good match with the empirical relation results for $\\delta t\\approx {\\rm few}\\times 10$~Myr and for a power-law CMF with $\\beta\\approx 2.35$. It should be noticed that the slope of this CMF for stellar clusters is virtually the same as for the Salpeter IMF ($\\alpha=2.35$) which applies for the early-type stars in these clusters. Conversely, adopting the empirical $SFR(M_{\\rm ecl,max})$ relation, $\\delta t$ can be calculated for different young-cluster power-law mass functions with exponent $\\beta$. We find that $\\delta t \\simless {\\rm few}\\times 10$~Myr for $\\beta\\simless2.4$. This value is nicely consistent with independent observations. For example, \\citet{HEDM03} find $2 < \\beta < 2.4$ for a sample of 939 LMC and SMC clusters after applying corrections for redding, fading, evaporation and size-of-sample effects.  The same holds true if a fundamental maximum star-cluster mass near $M_{\\rm ecl,max*}=10^7\\,M_\\odot$ is introduced. The existence of such a fundamental maximum cluster mass is supported by ``clusters'' with $M \\simgreat 5\\times 10^6\\,M_\\odot$ having complex stellar populations more reminiscent of dwarf galaxies that cannot be the result of a truly single star-formation event.  The short time-span $\\delta t \\approx {\\rm few} \\times 10$ Myr for completely-populating a CMF up to the maximum cluster mass of the population, $M_{\\rm ecl,max}\\le M_{\\rm ecl,max*}$, can be understood as being due to the high ambient pressures in the inter-stellar medium needed to raise the global SFR high enough for populous star-clusters to be able to emerge.  This short time-scale, which we refer to as a star-formation ``epoch'', does not preclude the star-formation activity in a galaxy to continue for many ``epochs'', whereby each epoch may well be characterised by different total young-star-cluster masses, $M_{\\rm tot}$. According to this notion, dwarf galaxies may experience unfinished ``epochs'', in the sense that during the onset of an intense star-formation activity that may be triggered through a tidal perturbation for example, the ensuing feedback which may include galactic winds may momentarily squelch further star-formation within the dwarf such that the cluster system may not have sufficient time to completely populate the cluster mass function. Squelching would typically occur once the most massive cluster has formed. Dwarf galaxies would therefore deviate notably from the $M_{\\rm ecl,max}(SFR)$ relation (\\S~\\ref{obs}).  The conclusion is therefore that the observed $SFR(M_{\\rm ecl,max})$ data can be understood as being a natural outcome of star formation in clusters and that the SFR at a given epoch dictates the range of star-cluster masses formed given a CMF that appears to be a Salpeter power law. The associated formation time-scales are short being consistent with the conjecture by \\citet{Elme00} that star-formation is a very quick process on all scales. Within about $10^7$~yr a complete cluster system is build (Fig.~\\ref{fig3a},~\\S~\\ref{fft}), while individual clusters form on a time scale of $10^6$ years and stars only in about $10^5$ years. Correspondingly, molecular-cloud life-times are short ($\\approx {\\rm few} \\times 10$ Myr) - suporting the assertion by \\citet{HBB01}.  Applying the empirical $SFR(M_{\\rm ecl,max})$ relation to the MW which has $SFR \\approx 3-5\\,M_\\odot$/yr \\citep{PrAu95} a maximum cluster mass of about $10^5\\,M_\\odot$ is expected from eq.~\\ref{empMmax}. It is interesting that only recently have \\citet{AlHo03} revealed a very massive cluster in our Milky-way with about 100 O stars (similar to R136 in the LMC). \\citet{Kn00} notes that the Cygnus~OB2 association contains $2600\\pm400$ OB stars and about 120~O~stars with a total mass of $(4 - 10) \\times 10^4\\,M_\\odot$, and that this ``association'' may be a very young globular-cluster-type object with a core radius of approximately 14~pc within the MW disk at a distance of about 1.6~kpc from the Sun \\citep[but see][]{BiBoDu03}.  This object may be expanded after violent gas expulsion \\citep{BK03a, BK03b}. The MW therefore does not appear to be unusual in its star-cluster production behaviour."
    },
    "0402/astro-ph0402407_arXiv.txt": {
        "abstract": " ",
        "introduction": "Time-variable emission over the entire observable spectrum is one of the defining characteristics of AGN.  Variability on time-scales of months to years provided the first key evidence that the emitting regions were extremely compact, leading to the suggestion that AGN are powered by massive black holes.  However, although the black hole paradigm has grown stronger due to a variety of subsequent observations, the origin of the variability largely remains a mystery.  In radio-loud AGN, some progress has been made in understanding the broadband variability in terms of jet models of emission\\cite{mch99}, but the situation is less clear in radio quiet AGN, which form the bulk of the AGN population. Because, in the optical waveband, the variability is fairly slow, it can only be studied in detail with long, well-sampled monitoring campaigns which are difficult to organise.  In the X-ray band, where the variability is much more rapid, short-term variability was originally studied using `long-looks' of a day or more duration, by X-ray satellites such as {\\it EXOSAT} and {\\it ASCA}, but longer time-scales were inaccessible due to the constraints of scheduling and pointing these satellites.  In 1995, the launch of the Rossi X-ray Timing Explorer ({\\it RXTE}) revolutionised the study of AGN variability, because the rapid slewing capability and flexible scheduling of {\\it RXTE} allowed well-sampled long-term monitoring of AGN X-ray variability for the very first time. \\\\ With {\\it RXTE}, it has been possible to study X-ray variability of radio-quiet AGN over a very broad range of time-scales for comparison with the (as it turns out) remarkably similar variability properties of stellar mass black holes in X-ray binary systems (BHXRBs).  Also, it has been possible to compare the long-term X-ray variability with the optical variability sampled by a few optical monitoring programs, to examine the relationship between the two bands, which we might expect to be dominated by different emission mechanisms (optically thin versus optically thick).  In this paper we will review our current understanding of the long-term X-ray variability of radio-quiet AGN, and how it relates to the X-ray variability on shorter time-scales. We will also consider the relationship between the X-ray and optical bands, and discuss models which might explain the variability in both bands. ",
        "conclusions": ""
    },
    "0402/astro-ph0402413_arXiv.txt": {
        "abstract": "Top-down models assume that the still unexplained Ultra High Energy Cosmic Rays (UHECR's) are the decay products of superheavy particles. Such particles may have been produced by one of the post-inflationary reheating mechanisms and may account for a fraction of the cold dark matter. In this paper, we assess the phenomenological applicability of the simplest instant preheating framework not to describe a reheating process, but as a mechanism to generate relic supermassive particles as possible sources of UHECR's. We use cosmic ray flux and cold dark matter observational data to constrain the parameters of the model. ",
        "introduction": "The possible observation of Ultra High Energy Cosmic Ray (UHECR's) events with primary energies above $10^{20}$\\,eV \\cite{Agasa} constitute one of the most intriguing puzzles in astroparticle physics (see, for example, \\cite{reviews}), although their origin and composition are not yet understood. The usual bottom-up scenarios in which particles should be accelerated by astrophysical objects do not seem to provide a convincing solution to the puzzle. The arrival direction of the primary particles should point to their sources because at such energies the intergalactic magnetic field does not deviate their direction of propagation. However, the clustering of UHECR events observed in the available data is not statistically significant and therefore there is no evidence that they arise from point sources \\cite{stat}. In addition, it would be necessary to overestimate several parameters of such sources and their acceleration regions in order to reach, marginally, the required energies \\cite{hillas}.  The problem concerning cosmic ray sources is related to the necessity that they must be located in our neighborhood, since particles propagating at high energies suffer a rapid degradation of their energy. For protons or nuclei as primaries, interactions with the cosmic microwave background should cause a loss of their energy due to photopion production. Such an effect should result in a discontinuity in the cosmic ray spectrum for energies above $\\sim 4\\times10^{19}$ eV, the so-called Greisen-Zatsepin-Kuzmin~(GZK) cutoff~\\cite{GZK}. If the Auger Observatory data confirms that this feature is not observed, it can be shown \\cite{Aharonian} that protons must have travelled less than $\\sim 100$ Mpc (attenuation length) in order to arrive at Earth with energies larger than $10^{19}$ eV. The attenuation length for photons depends on their initial energy and it is less than 100 Mpc for energies between $10^{12}$ eV  to $10^{22}$ eV \\cite{Protheroe}. Since neutrinos have a very small cross section with nucleons within the Standard Model, it seems difficult that they could produce air showers in our atmosphere unless they had a yet unknown interaction \\cite{domokos}.  In order to overcome such difficulties, another class of models have been proposed \\cite{berezinsky}. The primary particle would not acquire kinetic energy continuously inside an accelerating region (``bottom-up'' mechanism) as initially thought. Instead, the highly energetic cosmic rays would be originated by the decay products of superheavy particles of cosmological origin (``top down'' mechanism). For simplicity, we will consider that such particles have masses close to the GUT scale and would decay into known particles, as quarks and leptons that evolve following the QCD model \\cite{topdown}. The quarks hadronize producing a small fraction of nucleons and pions that in their turn decay into photons, neutrinos and electrons and their corresponding antiparticles. Therefore, from the decay of such a supermassive particle it is possible to produce energetic photons, neutrinos and leptons, together with a small percentage of nucleons. Depending on which kind of particle is the primary, different attenuation lenghts can be obtained, so that one can establish at what minimum distances the supermassive particles sources should be located.  There are different exotic candidates to play the main role in top-down scenarios, such as decaying topological defects \\cite{vilenkin} or evaporating primordial black holes \\cite{barrau}. The simplest top down models (at least from the particle physics point of view) involve supermassive metastable particles sometimes called WIMPzilla's \\cite{wimpzillas}. Due to their colossal masses, such particles were presumably produced during the post-inflationary epoch and could contribute to a part or to the whole of the dark matter that accounts for about 30\\% of the energy density of the Universe. In order to explain the theoretically estimated UHECR's fluxes \\cite{Kuzmin}, such particles must be decaying now and have to be located in our neighborhood, which is expected, assuming they are concentrated in our galactic halo.  In this work we study the possibility of producing WIMPzillas in the post-inflationary process called instant preheating, suggested by Felder, Kofman and Linde (FKL), originally proposed as an alternative preheating mechanism~\\cite{felder}. This process seems to be essential for particle production in models of quintessential inflation \\cite{HAR,SS}.  In such a scenario, scalar particles $\\chi$ are non-perturbatively produced from the coherent oscillations of the inflaton $\\phi$, have their masses ``boosted'' due to their coupling to the field $\\phi$, and subsequently decay into supermassive metastable fermions $\\psi$. The idea of examining stable supermassive particles in this context was addressed by Felder {\\it et al.}, but its consequences either as dark matter or as cosmic rays primaries were not calculated in detail.  More specifically, there is a relation between the density parameter of these particles and their lifetime. If $\\psi$-particles compose the whole of the dark matter ($\\Omega_{\\psi}\\equiv {m_{\\psi} n_{\\psi} \\over \\rho_{crit}} \\sim 0.3$) \\cite{Freedman}, a maximum lifetime limit will be found. On the other hand, a lower limit on the abundance of such particles can be obtained if lifetimes are constrained to be larger than the age of the Universe ($\\tau_{\\psi} \\geq10^{10}$ yr) \\cite{leticia}. As will be shown later, such limits impose severe constraints on the parameters of the FKL mechanism.  This paper is organized as follows. In the next section we review some features of the nonperturbative processes more directly related to the production of massive scalar particles. In Section III we perform a detailed calculation of superheavy $\\psi$ particle production, extending the previous results in \\cite{felder}. In Section IV we discuss our main results for produced particles considered as dark matter in our current Universe and present the parameter space for this model, which is in accordance with cosmological data. In the last section we present a summary of our main results and discuss their consequences. ",
        "conclusions": "We studied how to generate supermassive fermions that can explain the UHECR's in the context of a particle production mechanism suggested in Ref. \\cite{felder}. We obtained the parameter space for which such a mechanism can take place and concluded that some fine tuning of parameters seems to be necessary. Additionally, the lower limit on the $\\psi$ mass obtained, $m_\\psi \\gtrsim 10^{15} $ GeV, is rather robust. A typical signature of this model would be an unforeseen rise in the flux at the highest energies end of the cosmic ray spectrum, which could be observed by the next generation of experiments, like the Pierre Auger observatory. The cosmic ray spectrum must have a cutoff which is associated to the maximum energy possible to UHECR's and is independent of the GZK feature. If such a cutoff happen to be below $10^{15}$ GeV, this simplest version of the FKL mechanism should be discarded. On the other hand, it is this mechanism that can provide masses of such magnitude more naturally than any other top-down versions, so that if it is at all possible to measure such high energy cosmic rays and no cutoff in the UHECR's spectrum is observed by the next generation experiments, this model can become an attractive candidate. It is also interesting that, despite having perhaps too many free parameters, this model is rather constrained, and such a result is relatively insensitive to wide variations of the relevant cosmological parameters.  It is important to emphasize that we studied the production of non relativistic $\\psi$ particles only, and the scenario can be made more complex by considering the production of $\\psi$ particles that are relativistic at the preheating time but becomes non-relativistic along the Universe evolution. In this case, the energy transfer from the inflaton field to other fields may be not very efficient and it would be necessary to consider the dilution/concentration of $\\psi$ particles throughout the several phases that happened since inflation (coherent oscillations phase, radiation domination and matter domination) and the allowed parameter space may be widened."
    },
    "0402/astro-ph0402139_arXiv.txt": {
        "abstract": "We present results from a deep \\chandra\\ observation of MS1137.5+66, a distant (z=0.783) and massive cluster  of galaxies. Only a few similarly massive clusters are currently known at such high redshifts; accordingly, this observation provides much-needed information on the dynamical state of these rare systems. The cluster appears both regular and symmetric in the X-ray image. However, our analysis of the spectral and spatial X-ray data in conjunction with interferometric Sunyaev-Zel'dovich effect data and published deep optical imaging suggests the cluster has a fairly complex structure. The angular diameter distance we calculate from the \\chandra\\ and Sunyaev-Zel'dovich effect data assuming an isothermal, spherically symmetric cluster implies a low value for the Hubble constant for which we explore possible explanations. ",
        "introduction": "\\label{sec:intro}  Rich galaxy clusters are useful cosmological probes.  The richest clusters are thought to be massive enough to comprise a fair sample of the Universe, \\ie, their mass composition should reflect the universal mass composition \\citep{white1993,evrard1997} and so provide an efficient laboratory for measuring the baryonic to dark matter ratio. The mere existence of massive galaxy clusters at high redshifts can also place powerful constraints on the physical and cosmological parameters of structure formation models (e.g., \\citealt{peebles1989,bahcall1992,luppino1995,oukbir1997,bahcall1998,donahue1998,eke1998,haiman2001,holder2001}). The greatest leverage is provided by the most massive and distant clusters; in fact, to constrain $\\Omega_\\Lambda$ one must use clusters with redshifts greater than 0.5  (e.g., \\citealt{oukbir1997,holder2001}).  For these reasons, we are particularly interested in observations of massive galaxy clusters at high redshift.  A set of observations of distant clusters can be used to constrain the expansion rate and curvature of the universe.  The most noted approach is through the angular diameter distance relation.  The theoretical value of the angular diameter distance $D_A(z)$ is cosmology-dependent; $D_A(z)$ can be calculated from measurements of the X-ray temperature and surface brightness profile of a cluster's intracluster medium (ICM) analyzed in conjunction with a measurement of the cluster's Sunyaev-Zel'dovich effect \\citep[\\cf][]{birkinshaw1991,birkinshaw1994,myers1997,hughes1998,reese2000a,patel2000,grainge2002}, a spectral distortion of the cosmic microwave background (CMB) radiation by the hot, ionized ICM \\citep{sunyaev1970,sunyaev1972}.   Another approach to measuring cosmological constants with clusters uses the idea that the fraction of the total cluster mass contained in this atmosphere, the gas mass fraction $f_g$, traces the universal baryonic mass fraction, under the fair sample assumption. The fair sample assumption then implies also that the gas mass fraction in massive clusters should not vary.   With a well-chosen sample which includes clusters at high-redshifts, both the above determinations of the cluster gas mass fraction \\citep[\\cf][]{sasaki1996,pen1997,rines1999} and cluster distances can be used to constrain the geometry of the universe, since at high redshifts, the calculated cluster distances and gas mass fractions depend significantly on the cosmological parameters \\omegam\\ and \\omegal. Recent measurements of $D_A(z)$ via the X-ray/SZE method \\citep{mason2001,reese2002,jones2003} and of the gas mass fraction \\citep{grego2001} in samples of clusters imply values of the Hubble constant and the curvature of the universe in good agreement with the values determined independently with other methods. In the work of \\cite{reese2002}, the angular diameter distance to 18 clusters, distributed widely in redshift, is calculated.  For \\omegam=0.3, \\omegal=0.7, they find the mean Hubble constant for the sample is 60$^{+4}_{-4} \\ ^{+18}_{-13}$ km s$^{-1}$ Mpc$^{-1}$.  For seven low-redshift clusters, \\cite{mason2001} find a mean Hubble constant of 66$^{+14}_{-11}\\pm$15  km s$^{-1}$ Mpc$^{-1}$ for this same cosmology. \\cite{jones2003} find a mean value of 65$^{+8}_{-7}\\pm$15 km s$^{-1}$ Mpc$^{-1}$ for five moderate-redshift clusters. (The uncertainties, statistical and systematic, are reported at 68\\% confidence.)  These Hubble constant measurements agree within the uncertainties with the value derived from the Hubble Space Telescope \\Ho\\ Key Project \\cite[\\cf][]{mould2000}, 72$\\pm3\\pm7$ km s$^{-1}$ Mpc$^{-1}$, and Type Ia supernovae \\cite[\\cf][]{riess1998}, who find the Hubble constant to be $\\sim$65, in the cosmology considered here.  When determining cosmological parameters from cluster distances and mass fractions, one generally assumes the dominant cluster physics is well-represented by a simple picture: galaxies and an ionized intra-cluster medium in hydrostatic equilibrium in a large dark-matter potential, supported by thermal pressure.  In general, this is descriptive of clusters in the local universe, as evidenced by the strong correlations between X-ray luminosity and temperature \\citep{mushotzky1997,allen1998,arnaud1999,donahue1999} and cluster size and temperature \\citep{mohr2000}; and by the uniformity of the gas mass fraction in massive clusters \\citep[\\cf][]{djf95,mohr1999a}.  One can attempt to ameliorate the effects of departures from this simple picture on the accuracy of the results by determining cosmological parameters from samples of clusters, in which some of these systematic effects will average out, and taking care to choose the sample in an unbiased way.  Indeed, much of the observational cluster work cited above (e.g., \\citealt{mason2001,grego2001,reese2002}) uses such an approach.  In this work, we investigate the validity of the hydrostatic, isothermal, spherically symmetric assumption by studying observations of a massive (and therefore relatively bright) distant cluster.  We analyze a deep observation of the galaxy cluster MS1137.5+6625 taken with the \\chandra\\ Observatory's Advanced CCD Imaging Spectrometer (ACIS). \\ms1137\\ was observed as part of L. VanSpeybroeck's Cycle 1 Guaranteed Time Observations (GTO). These GTO observations, together with observations in subsequent \\chandra\\ observation cycles, form a collection of X-ray observations of galaxy clusters designed for measuring cosmological parameters to high accuracy. Over 1,250 ks of Chandra time were scheduled for this project in observation Cycles 1 and 2, and it included over 40 clusters by the close of Cycle 2. The project was designed to include a sufficient number of clusters to overcome systematic errors in the calculated cosmological parameters arising from cluster ellipticity and orientation \\citep[\\cf][]{sulkanen1999}. The clusters in the sample are massive (kT$_e \\gtrsim$ 5 keV) and distributed widely in redshift.  The clusters were chosen on the basis of having high X-ray luminosity and/or high X-ray temperature.  Those with extremely bright radio point sources at their centers were avoided.  The \\ms1137\\ observation is one of the longest observations scheduled for this project.  The Chandra Observatory carries instruments particularly well-suited to studying the distant, massive clusters in the cosmology project, owing to a combination of unprecedented angular resolution in the X-ray band; low background rates, particularly in the front-illuminated ACIS-I devices; and the ability to detect X-ray photons up to high energies (kT$_e \\sim$10 keV).  The \\xmm\\ X-ray observing facility has substantially more collecting area, permitting investigation of large-scale temperature gradients in distant clusters with observations of moderate length.  The \\chandra\\ facility, however, has significantly better angular resolution than \\xmm, and as is discussed in this paper, this becomes quite important for distant clusters; their compactness does require the higher angular resolution of \\chandra\\ to accurately model the cluster's spectrum and spatial structure.  The two facilities are therefore complementary for study of high-redshift clusters.  The \\chandra\\ observation reveals that while the X-ray image suggests this massive, distant cluster is relaxed, symmetric, and spherical, its story is actually considerably more complicated. We analyze the \\chandra\\ data in conjunction with SZE data and compare it with optical data.  We determine it has a remarkably compact density distribution, find evidence of asymmetric temperature structure, and measure an angular diameter distance which implies a very low value for the Hubble constant.  In this paper, we review previous observations of \\ms1137\\ in Section~\\ref{sec:prevobs} and describe the \\chandra\\ observations and the ICM temperature and density profiles we derive from them in Section~\\ref{sec:chandraobs}.  In Section~\\ref{sec:calculations}, we derive the ICM cooling time, ICM mass and total mass profiles from the results of Section~\\ref{sec:chandraobs}, and calculate the cluster distance and infer the Hubble constant from the X-ray and SZE data in Section~\\ref{sec:dist}.  We discuss these results, evaluate the cluster's suitability for cosmological work, and suggest possible interpretations in Section~\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion} The X-ray and optical images of the distant galaxy cluster MS1137.5+6625 show the X-ray surface brightness and optical light have a regular, symmetric, and quite compact distribution.  However, several lines of evidence suggest that the cluster structure is more complicated: \\begin{itemize} \\item The hardness ratio map and spectral analysis suggest that the cluster has a relatively more complicated temperature and metal abundance structure than density structure. \\item The gas mass fraction calculated from the SZE/\\chandra\\ data is only marginally consistent with the gas mass fraction calculated from the X-ray data. \\item The Hubble constant calculated for this cluster is low by a factor of $\\sim$2. (This is closely related to the discrepancy in the gas mass fraction and gas mass.) \\end{itemize} These calculated quantities are summarized in Table~\\ref{table1}.  The most compelling argument that \\ms1137\\ deviates from isothermal, spherical symmetry is that the Hubble constant value derived from it is so low.  The X-ray and SZE data sets are not sensitive enough individually to allow us to definitively test our assumptions, but taken together they do provide some insight.  Here, we discuss a number of possible explanations for our observations. \\subsection{SZE Data Systematics} The possible contribution to systematic uncertainty to the Hubble constant and gas mass fraction measurement by the SZE data itself is very low.  It is possible that there are undetected radio-bright point sources in the cluster field.  Unless they are unfortunately placed on the sidelobes, point sources will decrease the measured SZE effect and increase the measured \\Ho.  Combining the SZE observations with 1.4 GHz NVSS \\citep{condon1998} observations suggests that possible undetected point sources in the \\ms1137\\ field have a negligible effect on the derived Hubble parameter \\citep{reese2002}.  We find no evidence for a radio halo in this cluster in the NVSS image.  Radio haloes, \\ie, large scale diffuse radio emission, would again depress the measured SZE and therefore would cause an overestimate of the Hubble constant.  Anisotropies of the CMB of other types which might affect the SZE data are at very low levels.  A limit to the primary anisotropies at these angular scales was set by \\cite{dawson2001} and \\cite{holzapfel2000b}. They find the Rayleigh-Jeans temperature of these fluctuations to be less than 14 $\\mu$K, and this in turn should contribute less than 2\\% uncertainty to \\Ho.  Confusion from the kinematic SZE effect (the spectral distortion of the CMB made by the cluster moving with respect to the CMB rest frame) should also be quite small; for an 8 keV cluster moving with a 300 km s$^{-1}$ velocity component in the line of sight direction, the kinetic SZE effect will be $\\sim$4\\% of the thermal SZE (and the sign of the effect will depend on whether the cluster is moving towards or away from the observer) leading to an 8\\% error in the Hubble constant. It is therefore quite unlikely for the SZE to be severely in error due to peculiar velocity of the cluster gas.  \\subsection{Geometry} The Hubble constant is derived under the assumption that the cluster can be characterized by a spherically symmetric distribution.  This will be in error by the ratio of the projected cluster size $L_{PROJ}$ to the line of sight cluster size $L_{LOS}$: $H_{meas}=H_{true}\\times\\, L_{PROJ}/L_{LOS}$.  We have measured the apparent axis ratio to be 0.87$^{+0.06}_{-0.07}$ for the elliptical beta model; the core radius measured by the circular beta model is roughly the geometric mean of the two axes.  For an apparently circular distribution, the line of sight axis could be longer or shorter than the projected axis, leading to a prolate or oblate three dimensional cluster, respectively. For an apparently elliptical mass distribution, the true three-dimensional distribution has another degree of complexity in addition to the degree of prolateness or oblateness; the axis of symmetry could have any angle of inclination with respect to the plane of the sky, allowing the symmetry axis an essentially infinite range of lengths which would produce the same projected image. Additionally, there is no a priori reason that clusters need to be ellipsoidally symmetric; in fact, for axis ratios of less than $1/\\sqrt{2}$, the concentrically ellipsoidal isothermal beta-model for the density distribution is unphysical \\citep[\\cf][]{grego2000a} if the the cluster is indeed in hydrostatic equilibrium.  The gravitational potential supporting such a beta-model requires {\\emph {negative}} dark matter density.  A more sophisticated model for the non-spherical density distribution is clearly required in such cases, but is beyond the scope of this paper. We note that for measuring cosmological parameters from an ensemble of clusters, increasing the complexity of the model is not strictly necessary. For example, \\citep{sulkanen1999} shows that from a sample of triaxial clusters, one can reconstruct an unbiased estimator for the Hubble constant by using a spherical model.  In the simplest picture, the error in the Hubble constant is approximately the  ratio of the apparent cluster size (\\ie, the geometric mean of the two core radii) to the line of sight size.  In such a picture, for \\ms1137's measured Hubble constant to equal the sample mean of \\cite{reese2002}, L(projected)/L(line of sight) must be 0.58.  As a fiducial point, in a flux limited sample of clusters fitted with elliptical beta-models, \\cite{mohr1995} measure an axis ratio for the sample of 0.8$^{+0.14}_{-0.11}$.  For the low Hubble constant to be due solely to an elongated cluster, it would have to be have two unusual characteristics for a cluster chosen from an unbiased sample: it would have a much longer axis ratio than generally observed in projection, and would have its symmetry axis nearly aligned with the line of sight, basically a long cigar pointed at the observer.  If the symmetry axis were not nearly along the line of sight, the true axis ratio would need to be even greater.  Since such an elongation along the line of sight is so unexpected, this may indicate that the criteria by which \\ms1137\\ was detected favored such clusters.  \\ms1137\\ was originally identified in the EMSS \\citep{gioia1990}, using a 2.$'4\\times2.'4$ detection cell, a size which would just fit the cluster emission we observe with \\chandra.  It was detected with a signal-to-noise ratio of 5.3 in a survey with minimum S/N for detection of 4.0.  \\ms1137\\ is the coolest and least luminous of the z $>$ 0.5 clusters in the EMSS survey.  If there does exist a population of similarly luminous, highly elongated clusters at high redshifts, it is likely that only those aligned very nearly along the line of sight would be detected by this survey; a cluster also near the detection cutoff but which was elongated in the plane of the sky would deposit less of its flux into a single detection cell and may not be detected.  For example, if \\ms1137\\ had an intrinsic axis ratio of $\\sim$0.75 but with the major axis in the plane of the sky, it is unlikely that it would have been detected with sufficient significance to be included in the survey.  The selection effects from the original point-source algorithm used for the EMSS were discussed in \\cite{pesce1990}, and reprocessing of the Einstein data led \\cite{lewis2002} to conclude that the EMSS survey systematically missed clusters of low surface brightness.  Similarly, simulations of the EMSS detection algorithm by \\cite{ebeling2000} suggest the algorithm may miss clusters with significant substructure.  \\subsection{Temperature Gradients} Although we have no compelling evidence for a cooling flow at the cluster's center, we do have evidence of structure in the temperature and abundance distribution. Temperature structure could affect both the derived shape parameters and the derived emission-weighted temperature.  Temperature structure will affect the X-ray and SZE data differently. It should not strongly affect the X-ray-derived spatial shape parameters; for a given emission measure, the emitted flux in the band we use for spatial fitting only changes by $\\sim$ 20\\% for gas with kT$_e$ from 4 keV to 10 keV. Also, for the X-ray spatial fitting, the data are azimuthally averaged, so that such structure would have little effect.   The SZE measurements, by virtue of the spatial filtering inherent in interferometry, are insensitive to temperature gradients on scales larger than a few arcminutes, but can be quite sensitive to smaller scale gradients.  As the SZE effect is proportional to $n_e \\times T_e$, the density inferred at any particular point will be over (under) estimated from an under (over) estimate of the temperature. Temperature structure could result in the beta-model fitted to the SZE differing from that fitted to the X-ray data. We do not expect this poor fit to the beta-model to strongly affect the calculated Hubble constant, because the joint-fit to the model is dominated by the X-ray imaging data in this case.  It may, however, affect strongly the gas mass fraction determined from the SZE data, as this is derived from the SZE spatial fits.  If \\ms1137\\ had a non-hydrostatic pressure structure, then we may expect to see evidence of this in the SZE data.  The low surface brightness of the SZE-effect makes it difficult to image on small scales, however, and with the current data, we are unable to make a conclusion about non-hydrostatic pressure.  Future generations of SZE instruments may have such capability.  The temperature structure of the cluster could contribute to the low value of the Hubble constant in another way.  The temperature we use for the \\Ho\\ calculation is the emission-weighted temperature, and the relevant quantity for the calculation is actually the density-weighted temperature.  As a fiducial point for the discussion, we compare the emission-weighted temperature and the density-weighted temperature for a cluster with the following structure: \\be\\ equal to the best-fit value from our model fit, and gas within two core radii ($\\sim$200 kpc) at half the temperature of the gas from two core radii to ten: $T_e(r < 2 r_c) = 0.5\\times T_e(2 r_c < r < 10 r_c)$. Temperature structure of this magnitude, while unusually pronounced, has been observed by \\chandra\\ in the cluster Abell 1835 \\citep{schmidt2001}, a massive cooling-flow cluster with highly luminous optical emission-line nebulae.  In such a scheme, the emission-weighted temperature would be 0.80 times the density-weighted temperature, and the Hubble constant would be 0.64 times the true Hubble constant.  For the underestimation of the Hubble constant to be due {\\em solely} to an underestimate of the temperature, the emission-weighted temperature must be $\\sim$0.65-0.70 times the density-weighted temperature.  We stress that this is a simple calculation; the exact effect depends also on the specifics of the instruments used to make the measurements. However, even if \\ms1137\\ had as extreme a temperature gradient as the atypical Abell 1835, this could not explain the entire Hubble constant underestimate.  In conclusion, the massive, distant galaxy cluster \\ms1137\\ appears compact and regular in its optical and X-ray images, but evidence suggests that it is not an isothermal, spherical cluster.  The Hubble constant calculated for the cluster from X-ray imaging and spectra and interferometric SZE data is unusually low.  Many of the known systematic effects in the X-ray/SZE Hubble constant measurement tend to make the Hubble constant higher, and few effects make it lower.  (A thorough discussion of these systematic effects can be found in \\cite{reese2002}.)  Two effects which {\\em can} make the measured Hubble constant artificially low are elongation along the line of sight and temperature structure in the ICM.  We have evidence suggestive of a line of sight elongation.  The optical lensing data indicate the surface mass density is quite high, as evidenced by strongly lensed arcs, and that the mass distribution is very compact. The X-ray data show the ICM density distribution is unusually compact, as well.  We see structure in the hardness ratio map which could be interpreted as temperature structure, with cooler or less metal-enriched gas at the center of the cluster.  We also compare other calculated quantities besides the Hubble constant to investigate the systematic effects.  The SZE-derived gas mass is completely insensitive to geometry, though it is quite sensitive to temperature structure; if the emission-weighted temperature were biased low, the SZE gas mass would be an overestimate of the true gas mass.  The X-ray-derived gas mass is essentially insensitive to temperature structure, but has a moderate dependence on geometry; if the cluster were elongated along the line of sight, the X-ray gas mass would be a slight overestimate. The X-ray-derived gas mass is 2.1$^{+0.1}_{-0.1}\\ h_{65}^{-5/2}\\times 10^{13} M_\\circ$  and the SZE-derived mass is $3.0^{+1.0}_{-1.0}\\ h_{65}^{-2}\\times 10^{13} M_\\circ$.  Within the uncertainty the masses agree and we cannot place strong limits on temperature structure and elongation with the current datasets.  We conclude with two points: one, that cluster appearance is a poor criterion for choosing sources for cosmological tests.   \\ms1137\\ appeared as a possibly ideal cluster for cosmology, with its regular, compact, round, X-ray and optical images, and no evidence of a strong cooling flow.  However, it yields a Hubble constant measurement which is strikingly low.  We stress that a cluster sample must be selected via well-defined, objective criteria.  And two, that selection effects in the parent sample from which the clusters are drawn should be well-understood, \\ie, although for cosmological work we generally select clusters from existing X-ray surveys on the basis of luminosity rather than surface brightness, any selection for surface brightness in the X-ray survey itself will still be present in the luminosity-selected sample.  This situation can be ameliorated by future cluster surveys, such as the SZE survey discussed in \\cite{holder2001b}, which have selection functions different from the X-ray and optical cluster surveys."
    },
    "0402/astro-ph0402625_arXiv.txt": {
        "abstract": "We present a new and homogeneous set of explosive yields for masses $\\rm 13,~15,~20,~25,~30~and~35~M_\\odot$ and metallicities $Z=0$, $10^{- 6}$, $10^{- 4}$, $10^{- 3}$, $6\\cdot 10^{- 3}$, $2\\cdot 10^{- 2}$. A wide network extending up to Mo has been used in all the computation. We show that at low metallicities ($\\rm Z \\le 10^{- 4}$) the final yields do not depend significantly on the initial chemical composition of the models so that a scaled solar distribution may be safely assumed at all metallicities. Moreover, no elements above Zn are produced by any mass in the grid up to a metallicity $\\sim 10^{- 3}$. These yields are available for any choice of the mass cut upon request. ",
        "introduction": "A proper understanding of the chemical evolution of our galaxy and of the universe in general requires a good knowledge of the chemical composition of the matter ejected by stars of different masses and initial composition. Massive stars certainly play a pivotal role in the chemical enrichment of the interstellar medium because they are very probably responsible for the production of at least most of the intermediate mass elements (O through Ca). In spite of their central role in the general comprehension of the chemical evolution of the matter, only one extended set of models has been published so far: the one computed and discussed by Woosley \\& Weaver (1995), hereinafter WW95, and Timmes, Woosley \\& Weaver (1995), hereinafter TWW95. Their yields are based on presupernova models computed by assuming, among the others, no mass loss, no rotation, a moderate amount of overshooting and semiconvection, a value of the $\\rm ^{\\rm 12}C(\\alpha,\\gamma)^{\\rm 16}O$ calibrated on preexplosive yields and a network extending up to Ge. The explosions were computed in spherical symmetry and the yields eventually obtained by imposing the ejecta to have a specific final kinetic energy (their cases A, B and C). The initial chemical composition of the models at intermediate metallicities was obtained by means of a galactic chemical evolution model (described by TWW95).  Unfortunately, the present simulations of both the presupernova evolution and the explosion are still far from being robustly established. Qualitatively (and partly quantitatively) we know how and where the various nuclei are synthesized (see, e.g., WW95, Arnett 1996, Thielemann, Nomoto \\& Hashimoto 1996, Limongi, Straniero \\& Chieffi 2000), but still large uncertainties connected to both the hydrostatic evolution and the explosion of massive stars prevent a rigorous computation of the yields. Such uncertainties are mainly related to the efficiency of the convection (see, e.g., Chiosi \\& Maeder 1986, Woosley \\& Weaver 1988, Bazan \\& Arnett 1994), the determination of the cross section of a few nuclear processes (first of all the $\\rm ^{12}C(\\alpha,\\gamma)^{16}O$ - see, e.g., Weaver \\& Woosley 1993 and Imbriani et al. 2001), the time delay between the collapse of the core and the rejuvenation of the shock wave and the precise location of the mass cut (which is the mass coordinate that separates the part of the star that collapses in the remnant from the one that is ejected outward), even in spherical symmetry. To further complicate the situation, also rotation, mass loss, magnetic field and asymmetric explosions may also produce large variations in the final yields (see, e.g., Heger, Langer \\& Woosley 2000 and Maeda \\& Nomoto 2003).  Some years ago we started a long term project devoted to the study of the evolution of massive stars and their associated explosive yields (Chieffi, Limongi \\& Straniero 1998, Limongi, Straniero \\& Chieffi 2000, Limongi \\& Chieffi 2002, Chieffi \\& Limongi 2002a, Limongi \\& Chieffi 2003 - LC03). Since the beginning we made a strong effort to avoid the use of the various kinds of statistical equilibrium usually adopted to determine the chemical evolution of the matter at temperatures larger than, roughly, 3 billions degrees. Moreover we made an effort to fully couple the integration of the physical equations to the ones describing the evolution of the nuclear species in order to increase the numerical accuracy. Over the years we increased progressively the nuclear network that now extends up to Molybdenum. However, similarly to WW95, also our models are still computed by neglecting both mass loss and rotation. In our latest paper (LC03) of the series we presented our most updated version of the hydrostatic code (FRANEC) together to our new hydrodynamic code needed to follow the propagation of the blast wave. We also showed that the yields produced by a given stellar mass depend mainly on the location of the mass cut rather than from the explosion energy. This means that, as a first approximation, the yields corresponding to the ejection of different amounts of $\\rm ^{\\rm 56}Ni$ may be obtained by assuming an explosion strong enough to eject the full mantle and imposing by hand the mass cut at the desired $\\rm ^{\\rm 56}Ni$ abundance. Such a finding means that one can easily explore different choices for the mass cut without the necessity of recomputing many times the explosion of the models.  By making use of the latest versions of the two codes (hydrostatic and hydrodynamic) described in LC03, here we present a wide database of yields. In particular, we present the explosive yields produced by a grid of six masses ($\\rm 13$, $\\rm 15$, $\\rm 20$, $\\rm 25$, $\\rm 30$ and $\\rm 35~M_\\odot$) and six metallicities ($\\rm Z=0$, $10^{- 6}$, $10^{- 4}$, $10^{- 3}$, $6\\cdot 10^{-3}$, $2\\cdot 10^{-2}$).  The paper is organized as follows. The evolutionary code and the input physics adopted to compute the grid are briefly summarized in section 2. Section 3 is devoted to the discussion of the initial chemical composition used to compute the models in the intermediate metallicity range between the primordial and the solar one. A final discussion and conclusions follow. ",
        "conclusions": "The final explosive isotopic yields in solar masses of all the computed models are reported in Table 1, available only in electronic format, once all the unstable isotopes have decayed into their stable isobars. The yields of selected radioactive isotopes at $\\rm 10^{7}~s$ after the explosion are collected in the same table. For obvious reasons we could not present different sets of yields for different choices of the mass cut; hence we chose to present just one case, i.e. the one in which all masses eject $\\rm 0.1~M_\\odot$ of $\\rm ^{\\rm 56}Ni$. Any other choice is promptly available upon request.  The full set of elemental production factors (PFs) is shown in Figure 2. Each panel refers to a specific metallicity and each symbol refers to a given mass (see Figure caption). Let us remind that, in our case, the PF of any given isotope/element is defined as the ratio of each isotope's/element's mass fraction in the total ejecta divided by its corresponding initial mass fraction, i.e.,  ${\\rm PF}=X_{ejected}/X_{ini}$. Note that this definition is different from the one usually adopted by WW95, where ${\\rm PF}=X_{ejected}/X_{\\odot}$.  Some basic properties of the present yields may be seen by looking at Figure 2. First of all the PFs of all the elements from C to Zn significantly decrease as the metallicity increases, almost independently on the initial mass - the only exceptions being N and F. The reason for this is that, regardless of the mass of the star, the yields of the elements do not vary by more than an order of magnitude within the entire range of metallicities (see Table 1), while the $X_{ini}$ obviously scale directly with the initial global metallicity and hence they vary by several order of magnitudes. Such a strong dependence of the PFs on the metallicity simply means that the larger the metallicity, the more difficult is the further chemical enrichment.  A second feature is the well known {\\em odd-even effect}, i.e. that the difference between the PFs of the odd (Na to Sc) and the even nuclei (Ne-Ca) decreases as the metallicity increases: at the solar metallicity most of the elements show a roughly scaled solar distribution (see LC03 for a more detailed discussion of this topic). It is worth noting that, with the $\\rm ^{12}C(\\alpha,\\gamma)^{16}O$ rate adopted in the present calculations, Ne, Mg, Si, S, Ar and Ca preserve a scaled solar distribution at all the metallicities (see Imbriani et al. 2001 for a more comprehensive discussion of this topic).  A last feature worth mentioning here is that below $Z=10^{-3}$ there is a cutoff in the PFs at the level of Zn, i.e., no elements heavier than Zn are produced. On the contrary, above $Z=10^{-3}$ such a cutoff progressively reduces so that a consistent production of elements beyond Zn is obtained. Elements above Sr are not produced in a significant amount even at solar metallicity. This means that the observed abundances of elements above Zn in very metal poor stars must be attribute to stars (or, in general, to processes) outside the range presently analyzed.  Since the only other paper presenting a full set of yields is the WW95 one, we show in Figure 3 the comparison between the WW95 and the present yields, for two masses and three metallicities. Only elements up to Ge are shown because the nuclear network adopted by WW95 does not extend above this element. For this comparison we chose, for each stellar model, the mass cut that provides the ejection of the same amount of $\\rm ^{\\rm 56}Ni$ as in the corresponding WW95 model. Note that, since the grid of metallicities computed by WW95 does not coincide exactly with the ones presented here, the comparison shown in Figure 3 refers to models having a slightly different initial metallicity. We selected the $\\rm 20$ and the $\\rm 25~M_\\odot$ because are the ones that dominate the yields of a stellar generation having a Salpeter like IMF (see, e.g., LC03). The lower right panel in Figure 3 shows that there is a very good agreement between ours and the WW95 yields for the $\\rm 25~M_\\odot$ of solar metallicity. On the contrary, all other panels disclose significant (and not systematic!) differences between the two sets of yields. In particular there are a few things worth noting: a) both sets of models produce O and C in similar amounts (within a factor of two) while the N yields tend to be similar only for $\\rm Z\\ge10^{-3}$, b) the light elements Ne, Na and Mg tend to be significantly more produced in our models than in the WW95 ones while Al is produced in quite similar amounts, c) we tend to systematically underproduce the products of the explosive oxygen burning and incomplete Si burning, i.e. Si, S, Ar and Ca by roughly a factor of two with respect to WW95 (even if the relative scaling among these elements is remarkably similar), d) also the odd elements P, Cl and K, tend to be quite largely underproduced in our models with respect to WW95 and e) the Iron peak nuclei show a quite contradictory behavior because, while Ti is always in good agreement, Co and Ni are generally overproduced and Sc often underproduced relative to WW95.  A proper understanding of the sources of such differences, though of overwhelming interest, is extremely difficult because the chemical yields are, in general, the result of a complex interplay among the various hydrostatic evolutionary phases plus the subsequent passage of the shock wave (Chieffi, Limongi \\& Straniero 2000). For example, elements like N and Mg are not significantly affected by the passage of the shock wave and hence their final differences will mainly reflect a different presupernova evolution (but note that, e.g., O, that it is also a product of the hydrostatic burnings, is produced in very similar amount). Other elements are produced, viceversa, only by the explosive burnings and therefore one could think that playing with the mass cut could significantly improve the comparison. This is not the case. First of all let us note that the mass cut must be located within the region undergoing complete explosive Si burning because appreciable amounts of Sc, Co and Ni must be ejected. Hence the abundances of the elements produced by the explosive oxygen burning and/or incomplete explosive Si burning (Si, S, Ar, K, Ca, V, Cr and Mn) would not be modified by a changing of the mass cut. But also the comparison of the elements mainly produced by the complete explosive Si burning would not be improved by a changing of the mass cut because a better fit to any of the elements like Sc, Ti, Co and Ni would worsen the fit to the others.  A deeper comparison between these two sets is virtually impossible because either the two sets of models have been computed by adopting different choices for both the treatment of the convective layers and the rate of the $\\rm ^{12}C(\\alpha,\\gamma)^{16}O$ nuclear process, and also because the models on which the WW95 yields are based have never been published. The only possible comparisons between our presupernova models and the ones that are at the base of the WW95 yields have been presented in Limongi, Straniero \\& Chieffi (2000) and hence we refer the reader to that paper for such a comparison.  The differences between the WW95 and our yields are large enough that they should produce visible differences in GCE simulations and hence we strongly suggest the use of both sets of yields in the GCE modeling so to understand how alternative sets of yields influence our current understanding of the chemical evolution of the universe.  In conclusion, we provide in this paper a brand new set of yields in a wide range in both mass and initial metallicity. All the yields are freely available to the community for any choice of the mass cut (upon request). We have shown for the first time that the initial chemical composition does not affect significantly the final yields up to at least a metallicity of the order of $Z=10^{-4}$. We have also shown that a metallicity larger than $Z=10^{-3}$ is necessary to begin to produce elements beyond Zn up the neutron magic number $\\rm N=50$. The present yields are quite different from the WW95 ones and the observed differences cannot be simply explained in terms of one or few causes but are certainly due to the complex interplay among various aspects of both the hydrostatic evolution and the explosion itself that are very difficult to disentangle at the moment."
    },
    "0402/astro-ph0402555_arXiv.txt": {
        "abstract": "{ We self-consistently estimate the outflow rate from the accretion rates of an accretion disk around a black hole in which both the Keplerian and the sub-Keplerian matter flows simultaneously. While Keplerian matter supplies soft-photons, hot sub-Keplerian matter supplies thermal electrons. The temperature of the hot electrons is decided by the degree of inverse Comptonization of the soft photons. If we consider only thermally-driven flows from the centrifugal pressure-supported  boundary layer around a black hole, we find that when the thermal electrons are cooled down, either because of the absence of the boundary layer (low compression ratio), or when the surface of the boundary layer is formed very far away, the outflow rate is negligible. For an intermediate size of this boundary layer the outflow rate is maximal. Since the temperature of the thermal electrons also decides the spectral state of a black hole, we predict that the outflow rate should be directly related to the spectral state. ",
        "introduction": "Most of the galactic black hole candidates are known to undergo spectral state transitions (Tanaka \\& Lewin, 1995; Chakrabarti \\& Titarchuk, 1995, hereafter CT95; Ebisawa et al. 1996). Two common states are the so-called hard state and the soft state. In the former, soft-X-ray luminosity is low and the energy spectral index $\\alpha \\sim 0.5$ ($E_\\nu \\propto \\nu^{-\\alpha}$) in the 2-10keV range. In the latter state, the soft-X-ray luminosity is very high, and hard-X-ray intensity is negligible. There is also a weak power-law hard-tail component with an energy spectral slope $\\alpha \\sim 1.5$. In the two component advective flow (TCAF) model (CT95), the viscous Keplerian disk resides in the equatorial plane, while the weakly viscous sub-Keplerian flow flanks the Keplerian component both above and below the equatorial plane. The two components merge into a single component when the Keplerian disk also become sub-Keplerian. It is suggested (Chakrabarti, 1990) that close to a black hole, at around $10-15~r_g$, ($r_g = 2GM_{BH}/c^2$ is the Schwarzschild radius, $M_{BH}$ and $c$ are the mass of the black hole and the velocity of light respectively) the sub-Keplerian flow slows down due to the centrifugal barrier and becomes hotter. Chakrabarti (1999, hereafter Paper I) shows that this centrifugal pressure-supported boundary layer (CENBOL for short) region could be responsible for the generation of thermally-driven outflowing winds and jets and computed the ratio of the outflow to the inflow rate assuming a simple conical accretion disk model.  In the present {\\it paper}, we compute the {\\it absolute} value of the outflow rate as a function of the rates of the two inflow components, Keplerian and sub-Keplerian. This we do analytically following the recently developed procedure of obtaining shock locations (Das, Chattopadhyay and Chakrabarti, 2001). By dynamically mixing these two components using solutions of the viscous transonic flows we obtain the specific energy and angular momentum of the sub-Keplerian region. We use these pair of parameters to locate shocks in the flow, compute the compression ratio and from this, the outflow rate. We note that as Keplerian matter is increased in the mixture, the shock compression ratio goes down, and the outflow rate decreases. This is also the case even from a radiative transfer point of view --  when the Keplerian rate is high, the CENBOL region is completely cooled and the shock compression ratio $R\\sim 1$. Hence in the soft state, which is due to increase of the Keplerian rate, outflow should be negligible.  In the next Section, we present the governing equations to compute the outflow rates using a purely analytical method. We compute results for both the isothermal and adiabatic outflows. In \\S 3, we present our results for a single component sub-Keplerian flow. We also produce examples of realistic disks with Keplerian and sub-Keplerian components and obtain outflow rates as functions of the inflow parameters.  In \\S 4, we discuss our results and draw conclusions. ",
        "conclusions": "CT95  pointed out that the centrifugal pressure-supported boundary layer (CENBOL) of a black hole accretion flow is responsible for the spectral properties of a black hole candidate. In this {\\it Paper}, we present analytical results to show that this CENBOL is also responsible for the production of the outflows,  and the outflow rate is strongly dependent on the inflow parameters, such as specific energy and angular momentum.  We showed that in general, the outflow rate is negligible when the shock is absent and very small when the shock is very strong.  In intermediate strength, the outflow rate is the highest. As the specific angular momentum is increased, the outflow rate is also increased. This conclusion is valid when the flow is either isothermal or adiabatic.  We also demonstrated how a realistic two-component flow  (TCAF) consisting of Keplerian and sub-Keplerian components produces a significant amount of outflow. Since matter close to a black hole is sub-Keplerian by nature, the two components must mix to form a single sub-Keplerian flow which has positive specific energy and almost constant specific angular momentum. We showed that as the Keplerian rate of the disk is increased, the outflow rate is decreased as the shock compression ratio approaches unity. This conclusion, drawn from a dynamical point of view, is also corroborated by the spectral behavior as well --- as the Keplerian rate is raised, the post-shock region is cooled due to inverse Comptonization and the shock disappears. This reduces the thermal pressure drive and the resulting  outflow rate is reduced."
    },
    "0402/astro-ph0402280_arXiv.txt": {
        "abstract": "{We surveyed the central 4 x 4 degrees of the Galactic center for planetary nebulae in the light of [S~III] $\\lambda9532$ and found 94 PNe that were not previously known, plus 3 that were previously identified as possible candidates. For 63 of these 97 objects, we obtained spectra that are consistent with highly reddened PN while the other 34 could not be recovered spectroscopically and remain unverified. Of the 94 candidates, 54 and 57 were detected via radio at 3 and 6 cm, respectively. An additional 20 PNe candidates were found during follow-up H$\\alpha$ imaging but have not yet been verified spectroscopically. Based on the properties of IRAS sources in this region of the Galaxy, and on the total luminosity of the Galactic bulge, the expected number of PNe is $\\sim250$, only 50\\% more than the 160 PNe candidates now known. Thus, surveys for PNe in the bulge are approximately two-thirds complete with the remainder likely hidden behind dust. ",
        "introduction": "Estimates of the number of planetary nebulae (PNe) in the Galaxy have always been subject to large uncertainties, ranging, for example, from 6,000 to 80,000 (Peimbert 1993). The principal obstacle in deriving an accurate count is the very high level of extinction close to the galactic plane, especially when looking toward the Galactic center where a large fraction of the population is expected. Interest in the population count stems from studies of the chemical enrichment rates (e.g., Peimbert 1987) and comparisons between the populations in our Galaxy and those in external systems (e.g., Jacoby 1980, Peimbert 1993).  For example, when looking at M31, the PN system is strongly concentrated toward the nucleus (Ciardullo \\etal 1989). That is, the PN density follows the radial increase of starlight inward. Yet, the Galactic distribution fails to rise accordingly, in contrast to the rapidly increasing population of IRAS sources and OH/IR stars that have IR colors of PNe (see Figure 2 of Pottasch \\etal 1988). If, in fact, the PNe were to follow the same distribution as these sources, then a relative lifetime argument implies that there are $\\sim320$ PNe within 2 degrees of the Galactic center (see Figure 3 of Pottasch \\etal 1988).  Yet, prior to this survey, initiated in 1994, only 34 PNe, or 10\\% of the expected number, were known in this region.  The expectation of 320 PNe, though, is highly approximate because not all color-selected IRAS candidates are true PNe. Van de Steene \\& Pottasch (1995) found that roughly 25\\% of their IRAS sources could be recovered in a radio survey, implying a minimum estimate of 80 PNe for the central 2 degrees. On the other hand, many PNe are not detectable by IRAS and so, IRAS counts underestimate the true number of PNe. Unfortunately, there is no reliable estimate for this factor because neither the IRAS survey nor PN surveys are complete. We return to this question in Section 3 where we derive another estimate for the number of PNe in the Galactic Bulge.  Beyond the simple counting statistics of PNe, the greater value of their study in and near the Galactic center is to use them as probes of the kinematic and chemical history of the bulge population of stars. If it were possible to identify many hundreds of PNe in the bulge region, one could map out the matter distribution, including any dark matter contributions, as was done in the galaxy NGC 5128 by Hui \\etal (1995).  Perhaps most importantly, follow-up spectroscopy of PNe can tell us the rate at which the alpha elements were enhanced near the Galactic center. In particular, the elements O, Ne, S, and Ar survive the stellar evolution and nucleosynthesis process unaffected from when the star was formed (Forestini \\& Charbonnel 1997). By analyzing a PN spectrum, one can deduce the time of formation of the progenitor star, as well as its initial chemical composition. By analyzing many PNe, a timeline for elemental enhancements can be constructed, as was done for the LMC by Dopita \\etal (1997).  Once the chemical compositions are known for many PNe, they can complement the stellar compositions (McWilliam \\& Rich 1994) by providing information on different elements; or, PNe can be used in place of the more observationally challenging stellar composition analysis.  For these reasons, the interest in Galactic center PNe has grown since we began our survey. The most productive general survey for Galactic PNe has been described by Parker \\& Phillipps (1998) and Parker \\etal (2003) netting 1214 newly identified PNe in H$\\alpha$. Searches for PNe specifically toward the Galactic center and bulge have also been carried out by G\\'omez, Rodriguez, \\& Mirabel (1997), Beaulieu, Dopita, \\& Freeman (1998), Kohoutek (2002), and Boumis \\etal (2003).  In this paper, we report on the identification of many additional PNe within a few degrees of the Galactic center that were not previously known, thanks to a wide-field near-IR survey initially motivated by discussions with Stuart Pottasch. Van de Steene \\& Jacoby (2001) have already reported the results of synthesis radio observations at 3 and 6 cm that were carried out at the Australian Telescope Compact Array to improve the determination of extinction measures for these highly obscured PNe. In a subsequent paper, we will report the results of our spectroscopic follow-up survey that provides the validation for most of these objects as PNe, as well as radial velocities for kinematic studies, and in a few cases, chemical composition estimates and ages for the PNe and their progenitors to explore the chemical enrichment history of the Galactic bulge. ",
        "conclusions": "With this near-IR survey for PNe close to the Galactic center, we have shown that:  1. Highly extincted PNe can be found effectively by surveying in the emission line of [SIII] $\\lambda$9532, especially if augmented by a survey in H$\\alpha$.  2. The census of PNe within 2 degrees of the Galactic center is about two-thirds complete.  3. There are now many target PNe for follow-up chemical composition analysis but most of these are so highly reddened that their blue emission lines are likely to be unobservable. Thus, optical studies will be severely hampered except for a few of the brighter, less extincted, objects. We will present our spectroscopic results for several of these PNe in a subsequent paper. The remaining PNe to be found in the inner Galactic bulge will be even fainter and/or more seriously reddened."
    },
    "0402/astro-ph0402249_arXiv.txt": {
        "abstract": "We present the computation of effective refractive coefficients for inhomogeneous two-component grains with 3 kinds of inclusions with ${\\rm m_{incl}=3.0+4.0i, 2.0+1.0i, 2.5+0.0001i}$ and a matrix with ${\\rm m_m=1.33+0.01i}$ for 11 volume fractions of inclusions from 0\\% to 50\\% and wavelengths ${\\rm\\lambda}$=0.5, 1.0, 2.0 and 5.0 ${\\rm \\mu m}$. The coefficients of extinction for these grains have been computed using a discrete dipole approximation (DDA). Computation of the extinction by the same method for grains composed of a matrix material with randomly embedded inclusions has been carried out for different volume fractions of inclusions. A comparison of extinction coefficients obtained for both models of grain materials allows to choose the best mixing rule for a mixture. In cases of inclusions with ${\\rm m_{incl}}$=2.0+1.0i and 2.5+0.0001i the best fit for the whole wavelengths range and volume fractions of inclusions from 0 to 50\\% has been obtained for Lichtenecker mixing rule. In case of ${\\rm m_{incl}=3.0+4.0i}$ the fit for the whole wavelength range and volume fractions of inclusions from 0 to 50\\% is not very significant but the best has been obtained for Hanai rule. For volume fractions of inclusion from 0 to 15\\% a very good fit has been obtained for the whole wavelength range for Rayleigh and Maxwell-Garnett mixing rules. ",
        "introduction": "The primary goal of the research we present here has been to use computational electrodynamics methods to choose the best mixing rule for different materials (dielectric, semiconductor and metal) of inclusions in dielectric matrix. Light scattering computation for composite particles can be made using the discrete dipole approximation (DDA). \\citet{vaidya2001} applied this method to calculate extinction and scattering efficiencies for silicate sphere with embedded graphite inclusions and evaluated interstellar extinction. \\citet{andersen2003} tested the programs DDSCAT by \\citet{draine2000, draine2000b} and MarCoDES by \\citet{markel1998}. The program by \\citet{markel1998} is much faster than the one by \\citet{draine2000, draine2000b} which is very time consuming, however, it is ment to be used mainly for calculating the extinction by sparse clusters, whereas for compact clusters the accuracy of this method is very low. For this reason the DDA methods are not used for mass study of inter- or circumstellar extinction. This is why, many researchers still consider various models of grains such as host material in which other materials are embedded (for example \\citet{mathis1989} or \\citet{maron1989}) or core-mantle (\\citet{jones1988}) and multilayer spherical particles (\\citet{voshchin1999}). Besides, the use of a particular model compared to observations can provide information about the inhomogeneous grains. \\citet{perrin1990} computed cross section of extinction using: 1$^{0}$ DDA method assuming cubic lattice with elementary cell treated as inclusion of 10 \\AA\\ for spherical grain of radius 100 \\AA\\ and 2$^{0}$ Mie theory for grain with the same radius. They considered two mixtures: first a matrix of silicate with complex index of refraction proposed by \\citet{draine1985} and inclusions of thaolin (organic dielectric with imaginary part of refractive index ${\\rm k<0.001}$) and the second case adesite containing inclusions of water ice. Both cases concern only a mixture of dielectrics for two values of volume fraction of inclusions in the infrared region. In their paper the effective optical constants applied to the Mie theory were obtained from the most popular rules in astrophysics: Maxwell-Garnett's and Bruggeman's. \\citet{perrin1990} compared cross sections of extinction obtained by both methods and stated that the application of effective-medium theories (EMT) to the problem of light scattering by inhomogeneous grains is not straightforward. Other authors (e.g. \\citet{chylek2000}) compared the DDA with EMT method but only for very limited range of volume fractions of inclusions, wavelengths and materials. ",
        "conclusions": "The figures \\ref{fig2}-\\ref{fig13} show the dependence of mean values of efficiency factors for extinction ${\\rm Q_{l,j}^{rand}}$ and ${\\rm Q_{l,j,p}^{homog}}$ on the volume fractions of inclusions for the given mixing rule and wavelength and the best fitted curves.  \\begin{table*} \\centering \\begin{minipage}{116mm} \\caption[]{Goodness of fit for each mixing rule and each wavelength $\\Delta_{l,p}$ - ${\\rm m_{inc}=3.0+4.0i}$, ${\\rm \\sigma=10\\%~ of~ Q_{ext}^{rand}}$} \\label{tab1} \\begin{tabular}{@{}lccccl@{}} \\hline  Mixing rule   & $\\Delta_{l,p}(0.5\\mu m)$ & $\\Delta_{l,p}(1.0\\mu m)$ & $\\Delta_{l,p}(2.0\\mu m)$ & $\\Delta_{l,p}(5.0\\mu m)$ & $\\Delta_{p}$\\\\ \\hline  Bruggeman & 0.616E$+$00 & 0.643E$-$40 & 0.000E$+$00 & 0.000E$+$00 & 0.000E$+$00 \\\\ Hanai & 0.389E$+$00 & 0.162E$-$06 & 0.911E$-$02 & 0.449E$-$01 & 0.231E$-$06 \\\\ Lichtenecker & 0.996E$+$00 & 0.750E$-$22 & 0.000E$+$00 & 0.000E$+$00 & 0.000E$+$00 \\\\ Looyenga & 0.121E$+$00 & 0.000E$+$00 & 0.000E$+$00 & 0.000E$+$00 & 0.000E$+$00 \\\\ Maxwell$-$Garnett & 0.240E$-$01 & 0.595E$-$01 & 0.535E$-$21 & 0.616E$-$29 & 0.000E$+$00 \\\\ Rayleigh & 0.132E$+$00 & 0.433E$-$01 & 0.120E$-$11 & 0.123E$-$18 & 0.129E$-$26 \\\\  \\hline \\end{tabular} \\end{minipage} \\end{table*}  \\begin{table*} \\centering \\begin{minipage}{116mm} \\caption[]{Goodness of fit for each mixing rule and each wavelength $\\Delta_{l,p}$ - ${\\rm m_{inc}=3.0+4.0i}$, volume fraction of inclusions {\\bf 0-15\\%}, ${\\rm \\sigma=10\\%~ of~ Q_{ext}^{rand}}$} \\label{tab2} \\begin{tabular}{@{}lccccl@{}} \\hline  Mixing rule   & $\\Delta_{l,p}(0.5\\mu m)$ & $\\Delta_{l,p}(1.0\\mu m)$ & $\\Delta_{l,p}(2.0\\mu m)$ & $\\Delta_{l,p}(5.0\\mu m)$ & $\\Delta_{p}$\\\\ \\hline  Bruggeman & 0.826E$+$00 & 0.383E$-$12 & 0.518E$-$34 & 0.000E$+$00 & 0.000E$+$00 \\\\ Hanai & 0.592E$+$00 & 0.182E$-$02 & 0.184E$-$05 & 0.597E$-$09 & 0.950E$-$13 \\\\ Lichtenecker & 0.999E$+$00 & 0.158E$-$37 & 0.000E$+$00 & 0.000E$+$00 & 0.000E$+$00 \\\\ Looyenga & 0.157E$-$09 & 0.000E$+$00 & 0.000E$+$00 & 0.000E$+$00 & 0.000E$+$00 \\\\ Maxwell$-$Garnett & 0.100E$+$00 & 0.997E$+$00 & 0.599E$+$00 & 0.399E$+$00 & 0.994E$+$00 \\\\ Rayleigh & 0.100E$+$00 & 0.997E$+$00 & 0.624E$+$00 & 0.423E$+$00 & 0.996E$+$00 \\\\  \\hline \\end{tabular} \\end{minipage} \\end{table*}  \\begin{table*} \\centering \\begin{minipage}{116mm} \\caption[]{Goodness of fit for each mixing rule and each wavelength $\\Delta_{l,p}$ - ${\\rm m_{inc}=2.0+1.0i}$, ${\\rm \\sigma=2.5\\%~ of~ Q_{ext}^{rand}}$} \\label{tab3} \\begin{tabular}{@{}lccccc@{}} \\hline  Mixing rule   & $\\Delta_{l,p}(0.5\\mu m)$ & $\\Delta_{l,p}(1.0\\mu m)$ & $\\Delta_{l,p}(2.0\\mu m)$ & $\\Delta_{l,p}(5.0\\mu m)$ & $\\Delta_{p}$\\\\ \\hline  Bruggeman & 0.929E$+$00 & 0.245E$-$26 & 0.000E$+$00 & 0.000E$+$00 & 0.000E$+$00 \\\\ Hanai & 0.643E$+$00 & 0.216E$-$15 & 0.420E$-$38 & 0.000E$+$00 & 0.000E$+$00 \\\\ Lichtenecker & 0.256E$+$00 & 0.237E$+$00 & 0.987E$+$00 & 0.999E$+$00 & 0.896E$+$00 \\\\ Looyenga & 0.988E$+$00 & 0.000E$+$00 & 0.000E$+$00 & 0.000E$+$00 & 0.000E$+$00 \\\\ Maxwell$-$Garnett & 0.693E$-$01 & 0.103E$-$05 & 0.946E$-$11 & 0.450E$-$13 & 0.317E$-$26 \\\\ Rayleigh & 0.172E$+$00 & 0.158E$-$07 & 0.131E$-$14 & 0.327E$-$17 & 0.736E$-$35 \\\\  \\hline \\end{tabular} \\end{minipage} \\end{table*}  \\begin{table*} \\centering \\begin{minipage}{116mm} \\caption[]{Goodness of fit for each mixing rule and each wavelength $\\Delta_{l,p}$ - ${\\rm m_{inc}=2.5+0.0001i}$, ${\\rm \\sigma=2.5\\%~ of~ Q_{ext}^{rand}}$} \\label{tab4} \\begin{tabular}{@{}lccccl@{}} \\hline  Mixing rule   & $\\Delta_{l,p}(0.5\\mu m)$ & $\\Delta_{l,p}(1.0\\mu m)$ & $\\Delta_{l,p}(2.0\\mu m)$ & $\\Delta_{l,p}(5.0\\mu m)$ & $\\Delta_{p}$\\\\ \\hline  Bruggeman & 0.259E$-$06 & 0.371E$-$19 & 0.483E$-$03 & 0.342E$-$05 & 0.143E$-$29 \\\\ Hanai & 0.177E$+$00 & 0.645E$-$04 & 0.277E$-$01 & 0.999E$+$00 & 0.138E$-$02 \\\\ Lichtenecker & 0.682E$+$00 & 0.991E$+$00 & 0.100E$+$01 & 0.367E$-$01 & 0.884E$+$00 \\\\ Looyenga & 0.217E$-$20 & 0.455E$-$32 & 0.110E$-$03 & 0.966E$-$12 & 0.000E$+$00 \\\\ Maxwell$-$Garnett & 0.995E$+$00 & 0.822E$+$00 & 0.307E$+$00 & 0.263E$-$10 & 0.780E$-$05 \\\\ Rayleigh & 0.998E$+$00 & 0.375E$+$00 & 0.132E$+$00 & 0.397E$-$03 & 0.240E$-$01 \\\\  \\hline \\end{tabular} \\end{minipage} \\end{table*} The best fit of extinction for homogeneous grains consisting of material of averaged refractive index calculated from the mixing rule and extinction of grains consisting of DDA elements of two materials (matrix and inclusions) with different refractive indices (random) was obtained in the following way:\\\\ a) for the given mixing rule (subscript p) and the given wavelength (subscript l) the $\\chi_{p,l}^{2}$ was calculated from \\begin{equation} \\chi_{l,p}^{2}=\\frac{1}{10}\\sqrt{\\sum_{j=1}^{10}\\left(\\frac{ Q_{l, j}^{rand}-Q_{l,j,p}^{homog}}{\\sigma_{l,j,p}}\\right)^2}, \\end{equation} where summing is done for volume fractions of inclusions $\\rm j$, $Q_{l, j}^{rand}$ is averaged extinction coefficient for randomly located inclusions, $Q_{l,j,p}^{homog}$ is extinction coefficient for homogeneous grains for a given mixing rule. Because the values of $Q_{l, j,i}^{rand}$ for different random locations of inclusions (i) was not treated as equivalent of measured values, therefore $\\sigma_{l,j,p}$ is not the standard deviation. $\\sigma_{l,j,p}$ is the fraction $\\alpha$ of the averaged extinction coefficient $Q_{l, j}^{rand}$. \\begin{equation} \\sigma_{l,j,p}=\\alpha \\cdot Q_{l,j}^{rand}. \\end{equation} The calculations of $\\chi_{l,p}^{2}$ have been made for ${\\rm \\alpha=}$0.005, 0.01, 0.025, 0.05 and 0.1. The distribution of $Q_{l, j,i}^{rand}$ values has been assumed to be normal,\\\\ b) The goodness of fit $\\Delta_{l,p}$ has been calculated using the incomplete gamma function ${\\rm \\Delta_{l,p}=gammq(0.5\\nu, 0.5\\chi_{l,p}^{2})}$ described in \\citet{nr}, where ${\\rm \\nu=N-M}$ is the number of degrees of freedom, N is the nuber of points in the curve (number of volume fractions of inclusions), M are adjustable parameters (M=l number of mixing rules for calculating the $\\chi_{l,p}^{2}$. The obtained values of fitting coefficients $\\Delta_{l,p}$ for each mixing rule and each wavelength for ${\\rm \\alpha=0.025}$ are presented in Tables~\\ref{tab1} and \\ref{tab2}, and for the metallic inclusions - ${\\rm \\alpha=0.1}$ in Tables~\\ref{tab3} and \\ref{tab4} because for smaller values of ${\\rm \\alpha}$ the goodness-of-fit coefficients were very small. For the mixture of given materials determined by a mixing rule different fitting coefficients have been obtained for each wavelength. For different wavelength ranges different mixing rules may be used based on the given values of $\\Delta_{l,p}$.\\\\ c) In order to choose the best mixing rule in the whole considered range of wavelengths the value  \\begin{equation} \\chi_{p}^{2}=\\sum_{l=1}^{4}\\chi_{l,p}^{2} \\end{equation} has been calculated and the procedure described in b) carried out leading to the obtained goodness-of-fit coefficient $\\Delta_{p}$.  Summing up, we have examined 3 different materials as inclusions in a matrix of ${\\rm m_{m}=1.33+0.01i}$: \\begin{enumerate} \\renewcommand{\\theenumi}{(\\arabic{enumi})} \\item {\\bf ${\\rm \\bf m_{inc}=3.0+4.0i}$ (Figures~\\ref{fig2}-\\ref{fig5})} \\begin{enumerate} \\item The values of goodness-of-fit coefficients for ${\\rm \\alpha=}$0.025 and 0.05 are too small to be considered for analysis of the fit, therefore Table~\\ref{tab1} shows the values of ${\\rm \\Delta_{l,p}}$ and ${\\rm \\Delta_{p}}$ for ${\\rm \\alpha=}$0.1; \\item The goodness-of-fit of the extinction curve for each mixing rule strongly depends on the wavelength. For ${\\rm \\lambda=0.5\\mu m}$ the best fit has been obtained for the Lichtenecker rule. However, for the whole range of wavelengths the Hanai rule gives the best fit for this material. It is necessary to point out that the value of the fitting coefficient ${\\rm \\Delta_{p}}$ is very small. It may be necessary to investigate other mixing rules in order to obtain a better fit; \\item For a high volume factor of inclusions they may be placed so close that, on one hand, they may create inclusions of much bigger sizes and therefore more vulnerable to skin effect, on the other hand they may exceed the percolation threshold. Therefore, we have considered the fit for the volume factor of inclusions in the range from 0 to 15 \\%. The results are presented in Table~\\ref{tab2} from which we point out that for the whole wavelength range the best fit is obtained for Rayleigh and Maxwell-Garnett mixing rules. \\end{enumerate} \\item {\\bf ${\\rm \\bf m_{inc}=2.0+1.0i}$ (Figures~\\ref{fig6}-\\ref{fig9})} \\begin{enumerate} \\item The smallest value of ${\\rm \\alpha}$ which gives the goodness-of-fit coefficients suitable for further analysis is 0.025 and for this value the coefficients ${\\rm \\Delta_{l,p}}$ and ${\\rm \\Delta_{p}}$ have been calculated; \\item For ${\\rm \\lambda=0.5\\mu m}$ the best fit has been obtained for the Looyenga mixing rule (Fig~\\ref{fig6}). However, for the whole range of considered wavelengths the Lichtenecker rule gives the best fit (Table~\\ref{tab3}). \\end{enumerate} \\item {\\bf ${\\rm \\bf m_{inc}=2.5+0.0001i}$ (Figures~\\ref{fig10}-\\ref{fig13})} \\begin{enumerate} \\item The smallest value of ${\\rm \\alpha}$ which gives the goodness-of-fit coefficients suitable for further analysis is 0.025 and for this value the coefficients ${\\rm \\Delta_{l,p}}$ and ${\\rm \\Delta_{p}}$ have been calculated; \\item For ${\\rm \\lambda=0.5\\mu m}$ the best fit has been obtained for the Rayleigh and Maxwell-Garnett mixing rules. However, for the whole range of considered wavelengths the Lichtenecker rule gives the best fit (Table~\\ref{tab4}). \\end{enumerate} \\end{enumerate}   Considering the above results it is possible that different mixing rules should be applied for the same mixture of materials for different wavelengths, especially with higher volume fractions of inclusions. In spite of their deficiencies like simplifying idealisations while deriving them or basing only upon experiments, the mixing rules are able to provide important information about the inhomogeneous materials. In astrophysics generally the models by Maxwell-Garnett and Bruggeman are widely used because they are justified by theory. They are based on different topology of inclusions. Other mixing rules, less theoretically justified, are often neglected although they show better agreement with experimental data. However, \\citet{zakri} on the basis of effective medium theory found physical grounds for the Lichtenecker rule which, as shown by the calculations, gives better fits for inclusions with refractive indices ${\\rm m_{inc}=2.0+1.0i}$ and ${\\rm m_{inc}=2.5+0.0001i}$ in the whole considered wavelength range and for ${\\rm m_{inc}=3.0+4.0i}$ in the short wavelength region.   \\begin{figure} \\centering \\includegraphics[angle=0,width=84mm]{fig2.ps} \\caption{The dependence of efficiency factors for extinction on the volume fraction of inclusions for different mixing rules and wavelengths. Refractive index of the matrix is ${\\rm 1.33+0.01i}$ and for inclusion ${\\rm m_{inc}=3.0+4.0i}$, the wavelength $\\lambda=0.5{\\rm \\mu}$m. The solid curve \"random\" shows the best fit for 10 extinction dependencies for randomly distributed DDA elements with refractive index corresponding to inclusion among elements with refractive index of a matrix. The grain radius in all cases is ${\\rm 0.15 \\mu}$m.} \\label{fig2}  \\end{figure}  \\begin{figure} \\centering \\includegraphics[angle=0,width=84mm]{fig3.ps} \\caption{Same as Fig.~\\ref{fig2} but for $\\lambda=1.0{\\rm \\mu}$m.} \\label{fig3} \\end{figure}  \\begin{figure} \\centering \\includegraphics[angle=0,width=84mm]{fig4.ps} \\caption{Same as Fig.~\\ref{fig2} but for $\\lambda=2.0{\\rm \\mu}$m.} \\label{fig4} \\end{figure}  \\begin{figure} \\centering \\includegraphics[angle=0,width=84mm]{fig5.ps} \\caption{Same as Fig.~\\ref{fig2} but for $\\lambda=5.0{\\rm \\mu}$m.} \\label{fig5} \\end{figure}   \\begin{figure} \\centering \\includegraphics[angle=0,width=84mm]{fig6.ps} \\caption{Same as Fig.~\\ref{fig2} but for ${\\rm m_{inc}=2.0+1.0i}$ and $\\lambda=0.5{\\rm \\mu}$m.} \\label{fig6} \\end{figure}  \\begin{figure} \\centering \\includegraphics[angle=0,width=84mm]{fig7.ps} \\caption{Same as Fig.~\\ref{fig2} but for ${\\rm m_{inc}=2.0+1.0i}$ and $\\lambda=1.0{\\rm \\mu}$m.} \\label{fig7} \\end{figure}  \\begin{figure} \\centering \\includegraphics[angle=0,width=84mm]{fig8.ps} \\caption{Same as Fig.~\\ref{fig2} but for ${\\rm m_{inc}=2.0+1.0i}$ and $\\lambda=2.0{\\rm \\mu}$m.} \\label{fig8} \\end{figure}  \\begin{figure} \\centering \\includegraphics[angle=0,width=84mm]{fig9.ps} \\caption{Same as Fig.~\\ref{fig2} but for ${\\rm m_{inc}=2.0+1.0i}$ and $\\lambda=5.0{\\rm \\mu}$m.} \\label{fig9} \\end{figure}  \\begin{figure} \\centering \\includegraphics[angle=0,width=84mm]{fig10.ps} \\caption{Same as Fig.~\\ref{fig2} but for ${\\rm m_{inc}=2.5+0.0001i}$ and $\\lambda=0.5{\\rm \\mu}$m.} \\label{fig10} \\end{figure}  \\begin{figure} \\centering \\includegraphics[angle=0,width=84mm]{fig11.ps} \\caption{Same as Fig.~\\ref{fig2} but for ${\\rm m_{inc}=2.5+0.0001i}$ and $\\lambda=1.0{\\rm \\mu}$m.} \\label{fig11} \\end{figure}  \\begin{figure} \\centering \\includegraphics[angle=0,width=84mm]{fig12.ps} \\caption{Same as Fig.~\\ref{fig2} but for ${\\rm m_{inc}=2.5+0.0001i}$ and $\\lambda=2.0{\\rm \\mu}$m.} \\label{fig12} \\end{figure}  \\begin{figure} \\centering \\includegraphics[angle=0,width=84mm]{fig13.ps} \\caption{Same as Fig.~\\ref{fig2} but for ${\\rm m_{inc}=2.5+0.0001i}$ and $\\lambda=5.0{\\rm \\mu}$m.} \\label{fig13} \\end{figure}"
    },
    "0402/hep-th0402162_arXiv.txt": {
        "abstract": "We show that eternal inflation is compatible with holography. In particular, we emphasize that if a region is asymptotically de Sitter in the future, holographic arguments by themselves place no bound on the number of past $e$-foldings. We also comment briefly on holographic restrictions on the production of baby universes. ",
        "introduction": "New inflation models typically suffer from severe fine tuning problems associated with the choice of initial state. In the eternal inflation scenario, these problems are avoided because inflating bubbles persist along any time-slice, and these inflating regions self-reproduce leading to a fractal multiverse spacetime structure \\cite{Linde:1994xx}. From the point of view of a local observer the details of this multiverse structure are irrelevant, except to set up the initial conditions for the observer's own inflating bubble.  In this paper, we examine constraints on the eternal inflation scenario arising from holographic entropy bounds. Historically, the idea of holographic bounds \\cite{tHooft,LS} and their cousins \\cite{Bek73,Bek2000a,Bek2000b,BV} emerged from the study of black hole entropy (but see \\cite{MS1,MS2,ObsEnt}) and some researchers find motivation for such bounds in certain results from string theory. Such bounds are equivalent to the assumption that black holes are maximally entropic objects of a given size; they state that the entropy residing inside the relevant region is bounded by its surface area in Planck units \\cite{tHooft,LS,Easther:1999gk}.  On the other hand, for at least one proposed form \\cite{Bousso} of the holographic bound, it was argued in \\cite{FMW} that when i) the number of fields is small, ii) the matter $T_{\\mu\\nu}$ is not too anisotropic locally, and iii) temperatures are below the Planck scale, the bound follows as a consequence of the Einstein equations. Similarly, one generally does not expect holographic bounds to impose additional restrictions on thermodynamics at temperatures below the Planck scale. However, Banks and Fischler argued that holography (of a somewhat different form than that used in \\cite{Bousso}), together with certain additional assumptions, requires any late-time observer entering a region dominated by a small value of the cosmological constant to observe a bounded number of $e$-foldings \\cite{Banks:2003pt}. See \\cite{Wang:2003qr,Cai:2003zs} for subsequent related works.  Here we wish to emphasize one of the additional assumptions of \\cite{Banks:2003pt}.  In particular, \\cite{Banks:2003pt} considers a scenario where the universe is inflating at early times, passes through a matter-dominated regime, and then becomes asymptotically-de Sitter in the future.  The assumption of interest is that {\\it the total entropy of the universe in the early-time inflationary region can be computed by local field theory methods even when no observer can directly measure all of this entropy.}  In particular, we will see in section \\ref{bound} that most of this entropy lies outside the past light cone of any observer.  We are motivated to question this assumption by the observation that a similar assumption in the late-time de Sitter region would already violate any holographic bound on the entropy of the system. This is just the observation that de Sitter space expands to infinite size in the far future, so that any field theory with any finite cutoff contains an infinite number of degrees of freedom.  This observation is not new, and is well known to proponents of holography who propose that nevertheless de Sitter space is associated only with a finite number of states.  The usual resolution (see e.g. \\cite{Mobs}) is to note that this calculation does not contradict the experience of any observer in the spacetime, as such observers have access to only a small part of the entropy -- small enough, it turns out, to satisfy a holographic bound.  One then supposes that the true entropy of the system is comparable to the maximum entropy measurable by any given observer, and that field theory breaks down on scales large enough that it would predict violations of the holographic entropy bounds (see e.g., \\cite{CKN}).  Our goal here is to show that applying similar reasoning to the system considered in  \\cite{Banks:2003pt} yields a similar conclusion.  That is, in contrast to \\cite{Banks:2003pt} we assume that holographic bounds restrict only the field theory entropy in any past light cone, as field theory may generally acquire holographic corrections on larger scales. In this context we show that holography imposes no restrictions on inflation.  In particular, the number of $e$-foldings can be arbitrarily large. To distinguish our assumption from that of \\cite{Banks:2003pt} we refer to it as ``light-cone holography'' below\\footnote{Bousso's covariant entropy bound \\cite{Bousso} a special case of light-cone holography, but we allow much more general settings here.   We emphasize that some form of light-cone holography is essential for {\\it any} consistent holographic description of de Sitter space.}. We also comment briefly on claims \\cite{BanksFalseVac} of holographic restrictions on baby Universe formation. ",
        "conclusions": "We have established that light-cone holography places no bounds on the number of $e$-foldings to the past of a late-time observer and is thus consistent with the eternal inflation scenario.   This conclusion differs from that of \\cite{Banks:2003pt} because we do not share their assumption that local field theory can correctly compute the entropy of a volume larger than that contained in the past light-cone of any observer.  Again, we note that assuming local field theory to correctly describe the entropy of similar large volumes at late times would also contradict holographic bounds.  In particular, the corresponding calculation in pure de Sitter space would contradict the idea that asymptotically de Sitter space has a finite number of states, on which the discussion of \\cite{Banks:2003pt} also rests\\footnote{The authors of \\cite{Banks:2003pt} express their skepticism of the existence of a consistent theory which approximates local field theory in the inflating $\\Lambda$-region and leads to similar predictions for the CMB, but yields a smaller total entropy in the inflating regime.  We have no such example to offer, but see no reason why creating such a model is fundamentally more difficult than achieving the same goal for de Sitter space itself, a task not yet completed for which the authors of \\cite{Banks:2003pt} expect success. See \\cite{Guijosa,Verlinde} for some steps toward a model for de Sitter with a finite dimensional Hilbert space. Unfortunately, until a model exists for the perhaps simpler pure de Sitter case, there will be few solid grounds on which to resolve this difference of opinion.}.  Thus, at the current level of holographic understanding, we see no reason to suppose a contradiction between holography and a large number of $e$-foldings such as would arise in eternal inflation.  For eternal inflation to be self-reproducing, the inflaton must be able to fluctuate up its potential with some finite probability, giving rise to inflating regions with an increasing rate of expansion. One may also ask if there are holographic constraints on this process.  Discussions of related issues have appeared in \\cite{Coule}.  Let us begin with a clear example that illustrates how this mechanism can be compatible with holography. Consider a region of spacetime with some effective Hubble parameter $H$, homogeneous over many horizon volumes. Suppose a bubble with effective Hubble parameter $H'>H$ is nucleated inside this region with a size larger than the horizon size set by $H$. This process occurs with finite probability in the eternal inflation scenario \\cite{Linde:1994xx}. Applying the holographic bound to this situation \\cite{Easther:1999gk} one finds that the generalized second law yields no constraint on the evolution of this super-horizon size bubble, as it is unable to collapse. The bubble is then free to expand in a manner compatible with holography, and no contradiction is later reached when inflation ends in this bubble and a vast amount of entropy is produced, despite the fact the bubble started out near GUT scale size, with low entropy. From a quantum mechanical viewpoint, the system starts in a special state of low entropy, but as time evolves the state explores a larger subspace of the full Hilbert space of states, corresponding to the $H'$ bubble expanding into the ambient $H$ region. Clearly we have ignored inhomogeneities, but we believe this example suffices to illustrate the essential compatibility of the seeding mechanism of eternal inflation with holography.  Another oft discussed situation occurs when the initial radius of the bubble $H'$ is smaller than the ambient spacetime's inverse Hubble scale $H^{-1}$. For $H'>H$ this bubble will collapse and can form a black hole whose interior becomes a baby universe that undergoes inflation. For uncharged bubbles, Farhi and Guth \\cite{Farhi:1987ty,Farhi:1990yr} concluded the initial conditions for the formation of such a bubble are always singular. This may prevent the formation of such bubbles in the first place. Even if they are formed, a curvature singularity separates any external observer from the inflating interior of the bubble, so the application of holography is not entirely clear. The case of charged bubbles was analyzed in \\cite{Alberghi:1999kd}. There it was found these problems can be avoided, but a new difficulty appears because the inflating region lies inside a Cauchy horizon, which is unstable \\cite{SimpsonPenrose,ChandraHartle,PoissonIsrael,BurkoOri,Dafermos1,Dafermos2}.  Let us nonetheless assume that such problems are somehow solvable and that sub-horizon scale bubbles do play a role in seeding eternal inflation. In this scenario, we wish to analyze possible constraints of holography. If semi-classical physics is valid in the appropriate regions of spacetime, the bubbles will collapse and form horizons. Banks has argued the entropy of such bubbles should be bounded by the Bekenstein-Hawking entropy of the resulting horizon \\cite{BanksFalseVac}. In particular, the argument is that universes such as ours today have $S\\approx10^{85}$, which requires an event horizon of radius $10^{8}m$ , somewhat larger than the radius of the earth. The probability of nucleating such a large bubble in the early universe is extremely small.  However there are a number of subtleties in the above argument. Let us at least enumerate some of the possibilities, assuming we start with a GUT scale bubble that collapses to form a black hole. Such a GUT bubble might have formed during quantum fluctuations in the eternal inflation foam, or perhaps through interactions in a very high energy particle accelerator.  \\begin{enumerate} \\item One interpretation of the Bekenstein-Hawking entropy $S_{BH}$ is that $e^{S_{BH}}$ bounds the dimension of the Hilbert space of states associated with the region inside the horizon. Let us denote this Hilbert space by $\\mathcal{H}$. This interpretation is supported by string theory calculations of black hole entropy via D-brane methods. At first glance, the idea that a black hole with initial Bekenstein-Hawking entropy of order $S_{BH}\\approx10^{6}$ expands to give a large universe appears to conflict with this idea. In particular, suppose one assumes that time evolution maintains a sharp distinction between those states in $\\mathcal{H}$ and those orthogonal to $\\mathcal{H}$. By this we mean both a) that $\\mathcal{H}$ and its complement do not significantly mix under time evolution and b) that that the two classes of states appear quite different to local observers which experience them. In this case, a local observer inside the bubble can estimate the dimension of the Hilbert space of states similar to what she observes and compare this with $\\mathcal{H}$ for an initial GUT size bubble. While it has been argued \\cite{Banks:2002wr} that precise measurements are impossible for this observer, it is clear that there are certain classes of scattering observables that can be measured with very high precision. The observer in our present universe would then be able to conclude with a high degree of certainty that her Hilbert space is larger than that inherited from a GUT scale bubble and rule out creation of her universe via such a black hole . \\item However, it is not clear to what extent the assumption in (1) above is physically justified. In particular, consider assumption (1a), that the bubble remains in the space $\\mathcal{H}$ under time evolution. Certainly the original black hole Hawking radiates, and may well disappear in the distant future. Black hole complementarity suggests the state of the Hawking particles is actually equivalent to a state inside the black hole, and in particular to any baby universe so created. Since these Hawking particles explore a much larger Hilbert space, it is conceivable then that the entropy of the bubble is not constrained by holographic bounds. From the external point of view, the late time de Sitter phase with large entropy could have a complementary description as a high entropy state in the Hilbert space $\\mathcal{H}\\times\\mathcal{O}$, where $\\mathcal{O}$ is the Hilbert space of states outside the horizon. \\item Let us also consider assumption (1b), that $\\mathcal{H}$ and its compliment appear quite different (for all time) to local observers which experience them.  Such observers are unable to measure the exact late-time state of the full system, and so end up measuring the entropy of a locally accessible subsystem. It is not clear to us whether observations of such subsystems can indeed distinguish between universes produced via black holes and those which arose from other initial conditions. Let us suppose now that they cannot. Let us also note that the von Neumann ($Tr\\rho\\ln\\rho$) entropy of such a subsystem may well \\emph{exceed} the entropy of the full system, because the observer is unable to measure correlations with causally disconnected regions of the asymptotically de Sitter space. Indeed the exact von Neumann entropy will vanish if the system is in a pure state. It is thus conceivable that a late time observer could see a vastly larger entropy than the Bekenstein-Hawking entropy associated with the horizon of an initial black hole from which our universe somehow emerged. \\end{enumerate} We see that in order to arrive at a contradiction, one would need to prove the existence of more than $e^{S_{BH}}$ states which a) are macroscopically indistinguishable from our universe and b) could have been formed from a GUT-scale black hole. We conclude that successful production of a de Sitter region with large apparent entropy must produce some fine-tuning of the universe, but not that it is otherwise ruled out without additional assumptions. Such a fine tuning is not a surprise, as the instability of the charged black hole's Cauchy horizon and the resulting singularity already indicate that successful production of a universe via a black hole is either far from generic or is dependent on high energy effects not currently understood and associated with the singularity. If one believes that the process is possible at all, it is plausible that any fine-tunings required by holography are a natural result.  In summary, we see that holography appears quite compatible with eternal inflation. In particular, a late time observer sees no bound on the number of $e$-foldings or on any other parameters in the model of figure \\ref{cap:Spacetime-diagram}. Furthermore, the mechanisms of self-reproduction in eternal inflation survive holographic constraints. Holography may place strong constraints on branches of the eternal inflation spacetime that somehow emerge from black hole interiors, but even here such a conclusion follows only if one introduces additional assumptions. Because quantum gravitational processes are necessarily involved in the production of such regions, any such assumptions are necessarily difficult to test and must remain inherently speculative."
    },
    "0402/astro-ph0402005_arXiv.txt": {
        "abstract": "Core-collapse in massive stars is believed to produce a rapidly spinning black hole surrounded by a compact disk or torus. This forms an energetic MeV-nucleus inside a remnant He-core, powered by black hole-spin energy. The output produces a GRB-supernova, while most of the energy released is in gravitational radiation. The intrinsic gravitational-wave spectrum is determined by multipole mass-moments in the torus. Quadrupole gravitational radiation is radiated at about twice the Keplerian frequency of the torus, which is non-axisymmetric when sufficiently slender, representing a ``{\\em black hole-blob}\" binary system or a ``{\\em blob-blob}\" binary bound to the central black hole.  We here discuss line-broadening in the observed spectrum due to Lense-Thirring precession, which modulates the orientation of the torus to the line-of-sight. This spectral feature is long-lived, due to weak damping of precessional motion. These events are believed to occur perhaps once per year within a distance of 100Mpc, which provides a candidate source for Advanced LIGO. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0402/astro-ph0402233_arXiv.txt": {
        "abstract": "{We have measured the spectrum of the Cosmic X-ray Background (CXB) in the 2-8 keV range with the high throughput EPIC/MOS instrument onboard XMM-Newton. A large sample of high galactic latitude observations was used, covering a total solid angle of 5.5 square degrees. Our study is based on a very careful characterization and subtraction of the instrumental background, which is crucial for a robust measurement of the faintest diffuse source of the X-ray sky. The CXB spectrum is consistent with a power law having a photon index $\\Gamma=1.41\\pm0.06$, with a 2-10 keV flux of (2.24$\\pm$0.16)$\\times10^{-11}$ erg cm$^{-2}$ s$^{-1}$ deg$^{-1}$ (90\\% confidence level, including the absolute flux calibration uncertainty). Our results are in excellent agreement with two of the most recent CXB measurements, performed with BeppoSAX  LECS/MECS data (Vecchi et al. 1999) and with an independent analysis of XMM-Newton EPIC/MOS data (Lumb et al. 2002), providing a very strong constrain to the absolute sky surface brightness in this energy range, so far affected by a $\\sim$40\\% uncertainty. Our measurement immediately implies that the fraction of CXB resolved by the recent deep X-ray observations in the 2-10 keV band is of $80\\pm7$\\% (1$\\sigma$), suggesting the existence of a new population of faint sources, largely undetected within the current sensitivity limits of the deepest X-ray surveys. ",
        "introduction": "\\label{intro}  The discovery of a diffuse background radiation in the X-ray sky dates back to the birth of X-ray astronomy: the first evidence was obtained by during the same rocket experiment which led to the discovery of Sco X-1, the first extra-solar X-ray source (Giacconi et al. 1962). Later observations have demonstrated that the bulk of Cosmic X-ray Background (CXB) above energies of $\\approx$2 keV is of extragalactic origin, due to sources below the detection threshold. The first wide band measures of the CXB were made by HEAO$-$1 (1977): the CXB spectrum in the 2$\\div$10 keV range was well described by a simple power law with photon index $\\approx$ 1.4 (Marshall et al. 1980). More recently, several measurements of the CXB spectrum have been obtained at energies below 10 keV. While the results on the spectral shape confirmed a power law with $\\Gamma\\sim1.4$, the normalization of the CXB remained highly uncertain as a consequence of large discrepancies (well beyond the statistical errors) among the different determinations. A difference as large as $\\sim40$\\% is found from the highest measured value (Vecchi et al. 1999, using SAX data) to the lowest one (the original HEAO-1 experiment, Marshall et al. 1980).\\\\ Barcons et al. (2000) showed that cosmic variance cannot account for the differences among the previous measures of the CXB intensity, concluding that systematic errors and cross-calibration differences must have some role. Measurements published after the analysis of  Barcons et al. (2000), namely Lumb et al. (2002) with XMM-Newton EPIC and Kushino et al. (2002) with ASCA/GIS, while differing by $\\sim$15\\% only, either because of the small covered solid angle (Lumb et al. 2002), or because of large uncertainties in the stray light assessment (Kushino et al. 2002), do not allow to constrain the value of the CXB normalization to a much narrower range.\\\\ Such an uncertainty on the intensity represents a severe limit to the definition of a big picture, explaining which are the sources of the CXB and constraining their cosmological properties, in spite of the wealth of information coming from the deep pencil-beam X-ray surveys (Chandra Deep Field North, Brandt et al. 2001; Chandra Deep Field South, Giacconi et al. 2002), medium-deep wide angle X-ray surveys (HELLAS2XMM, Baldi et al. 2002) and their multiwavelength follow-up Campaigns. Even a basic information such as the resolved fraction of the CXB cannot be firmly evaluated.   \\\\ We have obtained a new measurement of the CXB spectrum in the 2-8 keV range with the EPIC/MOS\\footnote{ the {\\em pn} camera, having different characteristics, will require a different approach} cameras onboard XMM-Newton. Our study of the CXB is based on a large sample of high galactic latitude pointings for a total solid angle of $\\sim$5.5 square degrees, reducing the effects of cosmic variance. Our analysis includes a very robust characterization of the instrumental background properties, which is crucial for a robust measurement of the faintest diffuse source of the X-ray sky. Here we shall briefly report on some highlights of our study. The reader is referred to De Luca \\& Molendi (2004) for more details. ",
        "conclusions": "\\begin{figure*} \\centering \\includegraphics[angle=-90,width=9.5cm]{deluca_fig2.ps} \\caption{CXB intensity measurements. The flux in the 2-10 keV band is represented as a function of the epoch of the experiment. The plot is an update of Fig.3 of Moretti et al. (2003) including the results of the present work.  From left to right the CXB values are from Marshall et al. (1980) with HEAO-1 data; McCammon et al. (1983) with a rocket measurement; Gendreau et al. (1995) with ASCA SIS data; Miyaji et al. (1998) with ASCA GIS; Ueda et al. (1999) with ASCA GIS and SIS; Vecchi et al. (1999) with BeppoSAX LECS/MECS; Lumb et al. (2002) with XMM-Newton EPIC/MOS; Kushino et al. (2002) with ASCA GIS. Finally, the hollow star mark our own CXB measurement with the EPIC MOS cameras onboard XMM-Newton. All the uncertainties are at 1$\\sigma$ level. The horizontal dotted lines mark the range of CXB intensity constrained by our measurement after including an extra 3.2\\% uncertainty (1$\\sigma$ confidence level, rescaled from the 5\\% uncertainty quoted in the text, which corresponds to 90\\% confidence level) on the absolute flux calibration of the EPIC cameras (it is indeed a measure of the cross-calibration accuracy among different instruments). Such range represents the most robust estimate of the true absolute sky surface brightness in the 2-8 keV band. } \\end{figure*}   We have plotted in Figure 2 our measurement of the CXB intensity, together with previous determinations. Our value is fully compatible with the results of Vecchi et al. (1999) on BeppoSAX LECS/MECS data (total solid angle $\\Omega \\sim 0.77$ square degrees) and with the independent analysis of XMM-Newton EPIC/MOS data by Lumb et al. (2002) ($\\Omega \\sim1.2$ square degrees).\\\\ The ASCA/GIS measure of Kushino et al. (2002) is marginally consistent with ours. Such study was performed over a rather large solid angle ($\\Omega \\sim50$ square degrees) but the absolute flux determination accuracy was limited to 10\\% by the large stray light component ($\\sim40$\\% of the collected flux) affecting ASCA data. Very recently, Revnivtsev et al. (2003) published a new measurement of the CXB spectrum performed with the PCA instrument onboard RXTE, reporting an intensity consistent with the results of Kushino et al. (2002). They used data from nearly all the sky ($\\Omega \\sim 2.2\\,10^4$ square degrees); the instrumental background spectrum was extracted from only 25 ksec of dark Earth observation. The original HEAO-1 measurement (Marshall et al., 1980) is significantly lower than ours. Such measurement is very robust as for solid angle coverage ($\\Omega \\sim10^4$ square degrees); however, all of the subsequent determinations yielded invariably higher values of the CXB intensity, casting some doubts on the absolute flux calibration of the HEAO-1 instruments.\\\\ In our analysis, we used a compilation of sky fields covering a solid angle of $\\sim5.5$ square degrees, reducing the cosmic variance below the overall quoted uncertainty. Our measurement, performed with the well-calibrated EPIC/MOS instrument, relies on a very careful study and subtraction of the NXB, as well as on a detailed analysis of all the possible sources of errors. In conclusion, we believe that our measure, fully consistent with two of the most recent CXB determinations, set a very strong constrain on the CXB intensity in the 2-8 keV range, significantly higher than the former result from HEAO-1 data, assumed more than 20 years ago as a reference.\\\\ We can now compare our value of the CXB intansity with the source number counts derived from recent X-ray surveys, both deep/pencil-beam and medium-deep/wide-angle (Moretti et al., 2003). Our measurement of the CXB intensity, F$_{CXB}$=(2.24$\\pm$0.10)$\\times10^{-11}$ erg cm$^{-2}$ s$^{-1}$ deg$^{-1}$ ($1\\sigma$ error, including the absolute flux uncertainty), implies that 80$^{+7}_{-6}$\\% of the cosmic X-ray background has been resolved into discrete sources in the 2-10 keV band. The resolved fraction rises only to $\\sim$84\\% when extrapolating the LogN/LogS  down to fluxes of $\\sim10^{-17}$ erg cm$^{-2}$ s$^{-1}$, a factor of 10 below the sensitivity limits of the deepest X-ray surveys. Such a result suggests the existence of a new class of X-ray sources (possibly heavily absorbed AGNs, or star-forming galaxies, see Moretti et al. 2003 and references therein) emerging at fluxes fainter than $10^{-16}$ erg cm$^{-2}$ s$^{-1}$ and accounting for the remaining part of unresolved CXB."
    },
    "0402/astro-ph0402469_arXiv.txt": {
        "abstract": "We study the accretion flow of a hot gas captured by the black hole gravity in the presence of a thin cold accretion disk. Such geometrical arrangement is expected in Active Galactic Nuclei (AGN) and in galactic X-ray binary systems because both hot and cold gases are present in the black hole vicinity. Previous astrophysical literature concentrated on the evaporation of the cold disk in the classical heat conduction limit. Here we consider the inverse process, i.e. condensation of the hot gas onto the cold disk. We find two distinct condensation regimes. (i) In the classical thermal conduction limit, the radiative cooling in the hot gas itself force condensation above a certain critical accretion rate. Most of the flow energy in this case is re-emitted as X-ray radiation.  (ii) Below a certain minimum accretion rate, the hot electrons are collisionless and the classical heat flux description becomes invalid.  We use the ``non-local'' heat flux approach borrowed from the terrestrial laser heated plasma experiments. Due to their very large mean free path, the hot particles penetrate deep into the cold disk where the radiative losses are significant enough to enable condensation. In this case the hot flow energy is inconspicuously re-radiated by the transition layer in many UV and especially optical recombination lines (e.g., Ly$\\alpha$, $H\\alpha$, H$\\beta$) as well as via the optically thick disk emission.  We derive an approximate analytical solution for the dynamics of the hot condensing flow. If the cold disk is inactive, i.e.  accumulating mass for a future accretion outburst, then the two-phase flows appear radiatively inefficient. These condensing solutions may be relevant to \\sgra, low luminosity AGN, and transient binary accreting systems in quiescence. ",
        "introduction": "An optically thick accretion disk (e.g., Shakura \\& Sunyaev 1973) appears to be an important part of the accretion flow of gas into the black hole (BH) or a compact object. This conclusion is supported by spectral energy distributions (SED; see Elvis et al. 1994; Ho 1999; and Fig. 1 in Gierlinski et al. 1999), double peaked emission line profiles (e.g. McClintock et al. 2003; and references in Ho 2003), and eclipse mapping of binary systems (e.g. Wood et al. 1986). Even for \\sgra, a very dim source, there is now a suspicion that a cold {\\em inactive}, i.e. not accreting, disk may be present (Nayakshin, Cuadra \\& Sunyaev 2004).  There is also a significant amount of hot ($T > 10^7$ K) gas at large distances from the BH in galactic nuclei. In the best studied cases this hot gas is observed as close as its capture radius (e.g. in the giant eliptical galaxy and a LLAGN M87 [Di Matteo et al. 2003]; and in \\sgra\\ [Baganoff et al. 2003]), meaning that accretion of this hot gas on the BH is unavoidable. In galactic binary systems, a diffuse hot gas may be present near the disk outer rim due to a shock in the hot spot, if the radiative cooling time is longer than dynamical time. A shocked captured stellar wind from the secondary is another source of hot gas (and it may also form the cold disk itself; Kolykhalov \\& Sunyaev 1980).  Thus the accretion flow at large distances is often a two-flow problem (see Figure \\ref{fig:geometry}).  Thermal conduction in a multi-phase medium leads to either evaporation of the cold gas, or vice versa, condensation of the hot gas (e.g., Zel'dovich \\& Pickel'ner 1969; Penston \\& Brown 1970).  In reference to accretion flows, Meyer \\& Meyer-Hofmeister (1994; MMH94 hereafter) were the first to study the mass exchange between the cold disk and the corona. Their pioneering study, extended later by, e.g., Liu, Meyer \\& Meyer-Hofmeister (1997); Dullemond (1999); \\rozanska \\& Czerny (2000a,b), showed that the cold disk evaporates at low coronal accretion rates. At large accretion rates the radiative cooling suppresses the corona (\\rozanska \\& Czerny 2000a). The magnitude of the viscosity coefficient, $\\alpha$ (Shakura \\& Sunyaev 1973), was shown to be extremely important. In all the models the evaporation rate decreases very strongly with decreasing $\\alpha$ (e.g.  Meyer-Hofmeister \\& Meyer 2001). Recently, Liu, Meyer \\& Meyer-Hofmeister (2004) found condencing solutions for a sufficiently small value of the viscosity parameter.   Here we aim to study the condensation process systematically, i.e. attempting to cover a broad range of accretion rates in the hot flow. We first review the classical thermal conduction case. As expected, we recover qualitatively the results previously obtained by the above referenced authors. Namely, when the heat flux is classical, we find condensation at high and evaporation at low coronal accretion rates (\\S \\ref{sec:clas}). The critical accretion rate (at which no mass exchange between the disk and the hot corona takes place) is a strong function of $\\alpha$ (\\S\\S \\ref{sec:clascond} \\& \\ref{sec:paramspace}).  We then point out that for every value of $\\alpha$, there exists an accretion rate below which the classical thermal conduction treatment becomes invalid since the hot electron mean free path, $\\lambda$, is long compared with the flow height scale, $H$. To study the problem in this limit, we use a modified (see \\S\\S \\, \\ref{sec:basic} \\& \\ref{sec:vertical}) prescription for the heat flux that is borrowed from the laser-heated plasma experiments.  The heat flux is then proportional to the saturated heat flux coefficient, $\\phi\\le 1$. While this case still requires a future physical kinetics treatment, we show that there may exist an additional condensation regime. In this regime the hot particles penetrate deep into the cold gas. The cold gas density turns out to be great enough to re-radiate the deposited energy away, enabling condensation. With the help of certain approximations, we build an analytical solution (\\S\\S\\, \\ref{sec:dynamics} \\& \\ref{sec:analytical}) for the radial structure of the hot flow in this case.    We put our results into the context of the previous work in \\S \\ref{sec:paramspace}, where we determine the type of solution for a given combination of parameters (accretion rate, radius, $\\alpha$ and $\\phi$). We discuss the expected spectra from the two types of the condensing flows in \\S \\ref{sec:obs}. The classical thermal conduction condensing flow emits mainly in X-rays. Such solutions apply at relatively high accretion rates, and therefore they are likely to be relevant to bright or medium-bright AGN, such as Seyfert Galaxies. Due to these high accretion/condesation rates, the cool disk is probably active, meaning that the SED of these sources would be dominated by the emission from the small scale accretion disk in the vicinity of the last stable orbit. Such flows would obviously be radiatively efficient.   The second (non-local) condensation regime is applicable at lower accretion rates. It is therefore possible that accretion disks fed by condensation at these low rates would be inactive, as in quiescent states of transient binary sources. In this case the hot gas, settling onto the cold disk, effectively stops accreting. The hot accretion flow is thus terminated at some large distance away from the central object, which implies that the radiation emitted by such a flow will be much less luminous than expected from a flow onto the black hole. Furthermore, these flows may be quite dim in X-rays since the raditive cooling in the corona itself is small. The cold disk serves as a huge cooling plate (radiator) for the hot flow. The SED of such flows would then be dominated by their infra-red/optical thermal-like bumps rather than UV or X-ray frequency regions. We suggest that this radiative ``inefficiency'' (rather a time-delayed accretion) of condensing flows may be one of the reason why some LLAGN, and the galactic black hole \\sgra, appear underluminous. ",
        "conclusions": "\\label{sec:conclusions}  Here we studied the physics by which a hot flow above a cold accretion disk could condense. Such a condensation process is of a large practical importance for accretion in AGN and X-ray binaries. Without means to cool, the hot gas is an extremely inefficient kind of fuel for the black hole (e.g. BB99). Not only the gas radiates little, it also {\\em accretes} little (for recent numerical simulations see Proga \\& Begelman 2003). In contrast, if the hot gas condenses onto a cold accretion disk, the gas may temporarily loose its viscosity, i.e. the ability to accrete efficiently, but it remains tightly bound to the black hole. With time, when accretion through the cold disk is restarted, the hot gas will accrete onto the black hole. Thus our condensing scenario for the accretion of the hot gas onto the black hole is much more efficient in feeding the hot gas into the black hole than the non-radiative hot flows (BB99, Narayan 2002).   We have found two distinct condensation regimes: (i) radiative or classical, which takes place when the radiative cooling term in the corona is comparable with the viscous heating term (see \\S \\ref{sec:clas}) and the thermal conductivity is classical; and (ii) non-radiative. When the flow density is very low and the radiative losses are negligible. However the hot particles have very long mean free paths that allow them to penetrate the cold gas directly. Their energy is then re-radiated by the dense cold layers (see \\S \\ref{sec:vertical}). The cold disk thus serves as a cooling plate, radiator, for the hot flow. This second condensation regime is not expected based on the previous literature that employed the classical or the saturated heat flux formulations.  We also presented a simple analytical solution describing the radial structure of the hot flow for the non-radiative condensation. The solution is obtained under the assumption of hot flow temperature constant with radius and is meant to be an example only. The detailed structure of the flow is not as important for the flow energetics as long as one can clearly define a radius $R'$ where most of the corona condenses. As the hot gas condenses, the gravitational energy released per unit mass of hot gas is $\\sim G\\mbh/2R'$. If the cold disk is inactive, then most of the condensed mass is stored in the disk until the disk becomes massive enough (e.g. Siemiginowska, Czerny, \\& Kostyunin 1996; but note that galaxy mergers, etc., may be another ``direct'' way of triggering accretion outbursts in the inner parsec from the BH).  Thus the luminosity of the condensing flow (while the disk is inactive) is quite small compared with that expected for standard accretion flows (\\S \\ref{sec:inactive}).   Nuclei of LLAGN is an example where the two-flow geometry is natural -- e.g. see Fig. 2 of Ho (2003) and note that in reality the hot gas probably fills the whole available space because it is too hot to be confined (BB99). Currently, the bump-like infra-red feature seen in the SED of LLAGN is thought to be due to the {\\em accretion} through the disk. Assuming that the flow in the disk is continuous down to the last stable orbit however leads to a paradox since then one would expect a far brighter source. In contrast, if the cold disk is powered by the condensing corona and there is no accretion in the disk itself, there is then no problem with the source being too dim.  The current ``non-radiative'' condensation mechanism is thus somewhat inconspicuous in that the radiative output of such a flow is dominated by the emission of a warm atomic and ionized gas and not by the energy source -- the hot gas.   A similar picture may be relevant for quiescent disks in transient binary systems if not all of the accretion stream radiatively condenses in the hot spot (the hot spot emission is quite weak in many systems; see references in Lasota 2001, his \\S 7). Thus in some systems the ``disk'' continuum and line emission may be excited by shocks in the hot spot, whereas in others (with a weaker hot spot), the emission may be produced by the diffuse hot flow condensing onto the disk.   Since the accretion flow may span a broad range of radii, it is possible that different types of mass exchange solutions realize at different radii. For example, at a high enough accretion rate, the hot flow may be condensing radiatively at large radii, whereas at small radii, where the accretion rate is significantly reduced, it may condense by the non-radiative mechanism. It is equally possible for the accretion flow to be evaporating at large radii and condensing at small ones. Therefore we should keep in mind that in general a mix of the solutions is possible. This ``unfortunate'' multi-phase complexity of accretion flows is not something unique to the accretion process; after all the interstellar medium has a multi-phase structure too.   Finally note that due to processes not taken into account in the standard disk instability model (e.g. Cannizzo 1998), for example photo-ionization of the disk upper layers by starlight and X-rays; stellar impacts (Nayakshin et al. 2004), etc., there will always be a finite (weak) accretion onto the BH. This weaker accretion may power the jet emission.  Our results also indicate that a left alone cold disk will {\\em not} necessarily have to evaporate into a hot corona at a low $\\mdot$ via the classical thermal conductivity (MMH94). Indeed, when a hot corona just starts developing above the disk, the coronal accretion rate initially is very low and it is thus in the collisional, not classical, regime."
    },
    "0402/astro-ph0402143_arXiv.txt": {
        "abstract": "Observation of  cooling neutron stars can potentially provide information about the states of matter at supernuclear densities. We review physical properties important for cooling such as neutrino emission processes and superfluidity in the stellar interior, surface envelopes of light elements due to accretion of matter and strong surface magnetic fields. The neutrino processes include the modified Urca process, and the direct Urca process for nucleons and exotic states of matter such as a pion condensate, kaon condensate, or quark matter. The dependence of theoretical cooling curves on physical input and observations of thermal radiation from isolated neutron stars are described.  The comparison of observation and theory leads to a unified interpretation in terms of three characteristic types of neutron stars:  high-mass stars which cool primarily by some version of the direct Urca process; low-mass stars, which cool via slower processes; and medium-mass stars, which have an intermediate behavior. The related problem of thermal states of transiently accreting neutron stars with deep crustal burning of accreted matter is discussed in connection with observations of soft X-ray transients. ",
        "introduction": "\\label{introduc}  Neutron stars are the most compact stars in the Universe.  They have masses $M \\sim 1.4$ \\msun\\ and radii $R \\sim 10$ km, and they contain matter at supernuclear densities in their cores.  Our knowledge of neutron star interiors is still uncertain and, in particular, the composition and equation of state of matter at supernuclear densities in neutron star cores cannot be predicted with confidence.  Microscopic calculations are model dependent and give a range of possible equations of state (e.g., Lattimer \\& Prakash 2001; Haensel 2003), from stiff to soft ones, with different compositions of the inner cores (nucleons, pion or kaon condensates, hyperons, quarks).  One of the strong incentives for studying the thermal evolution of neutron stars is the promise that, by confronting observation and theory, one may learn about matter in the stellar interior. The foundation of the theory of neutron star cooling was laid by Tsuruta \\& Cameron (1966). The development of the theory was reviewed in the 1990s, e.g., by Pethick (1992), Page (1998a, b), Tsuruta (1998), and Yakovlev et al.\\ (1999). The latter authors presented also a historical review covering earlier studies. Some recent results have been summarized by Yakovlev et al.\\ (2002b, 2004a). In this paper we shall discuss the current state of the cooling theory and compare it with observations of thermal radiation from isolated neutron stars. We shall also outline the related problem of accreting neutron stars and its application to the quiescent radiation from soft X-ray transients. We shall describe mainly results that have been obtained since the middle of 1990s. The emphasis in this review is on discussing the dependence of the cooling behavior of neutron stars on physical properties of the matter and comparing the results with observation, rather than giving a detailed account of the results of microscopic theory. ",
        "conclusions": "As a consequence of improved measurements of thermal emission from cooling neutron stars in recent years, it has become very clear that the observations cannot be explained on the basis of a single universal cooling curve.  If thermal radiation from neutron stars in soft X-ray transient sources is due to nuclear burning processes deep in the crust, the observations of isolated neutron stars and X-ray transients can be analyzed within a common theoretical framework.  Moreover, observations may be explained in terms of physically reasonable models.  The basic ingredients of such a model are:  ({\\it a}) In the cores of massive neutron stars, a neutrino emission process faster than the modified Urca one operates. If one disregards the observations of SAX J1808.4--3658, it is not possible to pin down which of the faster processes (direct Urca processes for nucleons and hyperons, a pion condensate, a kaon condensate, or quark matter) is responsible, but if one includes the data from SAX J1808.4--3658, the nucleon or hyperon direct Urca process would be favored, and the other possibilities would be excluded.  ({\\it b}) In the cores of low-mass stars, neutrino emission is slower than that produced by the modified Urca process. For instance, this emission may be provided by neutron-neutron bremsstrahlung while other potentially efficient neutrino processes may be suppressed by strong superfluidity of protons.  ({\\it c}) Medium-mass stars show cooling intermediate between slow and fast. In particular, they may cool via enhanced neutrino emission partly suppressed by proton superfluidity. The mass range for these stars is determined by the density range over which the transition in the neutrino emission rate from slow to fast occurs. Some physical models of neutron star interiors contradict observations, for instance, the model of mild $^3$P$_2$ neutron superfluidity in the stellar cores with a maximum superfluid transition temperature $T_{\\rm cn}(\\rho)$ in the range from $2 \\times 10^8$ to $2 \\times 10^9$~K.   It is unlikely that advances in understanding the nature of the interiors of neutron stars will come from a single piece of evidence, but rather from a systematic appraisal of a variety of different sorts of evidence, just as in many legal cases. Directions for future study include:  $\\bullet$ Further observations of thermal radiation from neutron stars. A search for new very cold or very hot stars would be useful. Very cold neutron stars would rule out the possibility of not too fast neutrino emission produced by exotic matter in neutron star cores.  $\\bullet$ Further theoretical investigations of the effects of correlations in dense matter.  In particular, the role of tensor correlations needs to be reexamined following the work of Akmal \\& Pandharipande (1997), which found a strong increase of tensor correlations, a sign of incipient pion condensation at relatively low densities, and the recent study by Schwenk \\& Friman (2004) which pointed to the strong modification of the tensor force by the nuclear medium.  $\\bullet$ Information about neutron stars obtained from studies of cooling needs to be integrated with what has been learned by other means.  Examples are other observations of neutron stars, for instance, measurements of their radii or gravitational redshifts. Of special importance are observations of neutron stars in binary systems, which can be used to determine neutron star masses.  Even a firm lower bound on a neutron star mass obtained from, e.g., radio observations of compact binaries containing pulsars (either binary neutron stars or pulsar--white-dwarf binaries, such as J0751+1807 reported recently by Nice et al.\\ 2004), could rule out a number of theoretical equations of state.  It is also important to ensure that the physical input to neutron star calculations is consistent with experimental nuclear physics data on correlations between nucleons, hyperons and other degrees of freedom in dense matter.  {\\bf Acknowledgment.} We are grateful to O.\\ Gnedin, P.\\ Haensel, A.D.\\ Kaminker, K.\\ Levenfish, A.\\ Potekhin, and A.\\ Shibanov, DY's coauthors on papers discussed in this review, and to G.G.\\ Pavlov for enlightening remarks on observational data. We are also grateful to Olga Burstein and Daniel Cordier for critical comments which improved the presentation of the text and Table \\ref{tab:param}. This work has been supported partly by the RFBR, grants 02-02-17668 and 03-07-90200."
    },
    "0402/astro-ph0402375_arXiv.txt": {
        "abstract": "A sample of 19 low redshift (0.03$<$z$<$0.07) very luminous infrared galaxy (VLIRG: $10^{11}L_\\odot< $ L[8-1000 $\\mu$m] $ < 10^{12} L_\\odot$) systems (30 galaxies) has been imaged in $B$, $V$, and $I$ using ALFOSC with the Nordic Optical Telescope. These objects cover a luminosity range that is key to linking the most luminous infrared galaxies with the population of galaxies at large.  As previous morphological studies have reported, most of these objects exhibit features similar to those found in ultraluminous infrared galaxies (ULIRGs), which suggests that they are also undergoing  strong interactions or mergers. We have obtained photometry for all of these VLIRG systems, the individual galaxies (when detached), and their nuclei, and the relative behavior of these classes has been studied in optical color-magnitude diagrams.  The observed colors and magnitudes for  both the systems and the nuclei lie parallel to the reddening vector, with  most of the nuclei having redder colors than the galaxy disks.  Typically, the nuclei comprise 10 percent of the total flux of the system in $B$, and 13 percent in $I$. The photometric properties of the sample are also compared with previously studied samples of ULIRGs. The mean observed optical colors and magnitudes agree well with those of cool ULIRGs. The properties of the nuclei also agree with those of warm ULIRGs, though the latter show a much larger scatter in both luminosity and color. Therefore, the mean observed photometric properties of VLIRG and ULIRG samples, considered as a whole, are indistinguishable at optical wavelengths. This suggests that not only ULIRG, but also the more numerous population of VLIRGs, have similar rest-frame optical photometric properties as the submillimeter galaxies (SMG), reinforcing the connection between low-{\\it z} LIRGs -- high-{\\it z} SMGs. When the nuclei of the {\\it young} and {\\it old} interacting systems (classified according to a scheme based on morphological features) are considered separately, some differences between the VLIRG and the ULIRG samples are found. In particular, although the young VLIRGs and ULIRGs seem to share similar properties, the old VLIRGs are less luminous and redder than old ULIRG systems. If confirmed with larger samples, this behavior suggests that the late-stage evolution is different for VLIRGs and ULIRGs. Specifically, as suggested from spectroscopic data, the present photometric observations support the idea that the activity during the late phases of VLIRG evolution is dominated by starbursts, while a higher proportion of ULIRGs could evolve into a QSO type of object. ",
        "introduction": "Luminous infrared galaxies have been the subject of numerous studies over the past years (see, for instance, Sanders \\& Mirabel 1996; Veilleux et al. 2002, and references therein). Many of these works, both from the ground and space-based, have been focussed on the most energetic objects: the Ultraluminous Infrared Galaxies (ULIRGs: $L_{ir}=L[8-1000\\mu m]  > 10^{12} L_\\odot$) (e.g. Sanders et al. 1988; Melnick and Mirabel, 1990; Leech et al. 1994; Murphy et al. 1996; Clements et al. 1996; Surace et al. 1998 and 2000; Borne et al. 2000; Colina et al. 2001; Farrah et al. 2001; Kim et al. 2002, Bushouse et al. 2002 and references therein). These studies have found that the vast majority of ULIRGs are strongly interacting or advanced merger systems. The high IR luminosities are attributed to dust emission, with the heating source being varying combinations of interaction-induced starbursts and active galactic nuclei (AGN). ULIRGs also appear to be forming moderately massive ({\\it L*}) field ellipticals (e.g. Genzel et al. 2001 and references therein).  At a lower energy, the Very Luminous Infrared Galaxies (VLIRGs: $10^{11}L_\\odot<  L_{ir} < 10^{12} L_\\odot$, see Section 2.1) have not been the subject of as much scrutiny (although see, for instance, Lawrence et al. 1989; Kim et al. 1995; Wu et al. 1998 a and b). However, this subclass of lower luminosity objects has become increasingly interesting for several reasons. First, the above mentioned studies on ULIRGs have shown that many of the fundamental properties of these objects, such as the frequency of interactions, the interaction phase, the frequency of AGN, etc., all appear to correlate with the IR luminosity (Sanders \\& Mirabel 1996, and references therein).  VLIRG studies offer the possibility to analyze these properties over a wider luminosity range. Also, VLIRGs represent a much larger fraction of galaxies in the local universe than ULIRGs. In particular, the local density of VLIRGs is $\\sim$2 orders of magnitude higher than the density of ULIRGs (see, for example, Soifer {\\it et al}. 1987; Saunders {\\it et al}.  1990). Moreover, the infrared luminosities of VLIRGs are between those of the ULIRGs and local field spirals, for which typically $L_{ir} = 10^{10}-10^{11} L_\\odot$ (Rieke and Lebofsky, 1986). Therefore, VLIRGs represent a key link between the ULIRGs and the population of galaxies at large. VLIRGs are also believed to be low-redshift analogs of the galaxies that give rise to the far-IR background (Hauser et al., 1998) (e.g. high-redshift sub-mm galaxies, SMGs or SCUBA sources;  Smail et al. 1998). The nature of the SMGs is not yet clear. Some of them could be ULIRGs, as has recently been suggested by Frayer et al. (2003) from their near infrared (rest frame optical) magnitudes and colors. However, the majority could be related to the significantly more numerous class of VLIRGs. Therefore, well studied local samples of VLIRGs are valuable when establishing the relationship with the high redshift populations. This is especially true in light of the ongoing deep surveys such as, for instance, the Great Observatories Origin Deep Survey (GOODS, Dickinson et al. 2003; Giavalisco et al. 2004), and the Ultra Deep Field (Beckwith et al. 2003).  This paper presents recently obtained ground-based optical $B$, $V$, and $I$ images for a sample of VLIRGs. In Section 2, we comment on previous optical imaging surveys of VLIRGs, describe the characteristics of our sample, and give details about the observations. Section 3 describes the data reduction process. In Section 4, the photometric properties are presented and compared with previously studied  samples of ULIRGs. In Section 5, we present our conclusions. An appendix gives additional comments on individual objects. ",
        "conclusions": "We have obtained multi-wavelength optical imaging and photometry of a sample of 19 low redshift (0.03$<$z$<$0.07) very luminous infrared galaxy (VLIRG) systems. These objects have morphological characteristics similar to those found in ULIRGs suggesting strong interactions and/or mergers. The main conclusions of the present study are:  1)  In the color-magnitude diagram, the nuclei and the systems have well defined distributions which run along the reddening line, suggesting that part of the observed scatter is due to extinction effects. Most of the nuclei have redder colors than the galaxy disks, though there are a few exceptions. Typically, the nuclei comprise 10 percent of the total flux of the system in B, and 13 percent in I. Considering the mean extinction values from the Balmer decrement given by Veilleux et al., the mean intrinsic magnitudes of the nuclei in the VLIRG sample are: $< M_B>_o$ = -22.75 and $<M_I>_o$ = -21.90 (i.e. $(B-V)_o$ = - 0.85).  2) The optical colors and magnitudes for the nuclei and the systems of the VLIRGs agree well with those for the cool ULIRGs. The VLIRG nuclei also agree with those of the warm ULIRGs, though the latter show a much larger scatter in both luminosity and color.  Despite the difference in bolometric/infrared luminosity, the photometric properties of the VLIRG and ULIRG samples, considered as a whole, are indistinguishable at optical wavelengths (i.e. morphologies, compactness, magnitudes, and colors).  3) The recent suggestion by Frayer et al., based on near infrared photometry, that nearby ULIRGs could be the local analogue of submillimeter galaxies (SMGs) can also be extended to the more numerous population of VLIRGs. In fact, the previous conclusion suggests that the optical luminosities and colors of VLIRGs are also in good agreement with the corresponding rest-frame values of the SMGs, reinforcing the connection low-{\\it z} LIRGs -- high-{\\it z} SMGs.  4) When considering {\\it young} and {\\it old} systems separately, according to the interaction classification scheme proposed by Surace (see also Veilleux et al. 2002), some differences are found between the VLIRG and ULIRG samples. In particular, although the young VLIRGs and ULIRGs seem to share similar properties, the old VLIRGs are less luminous and redder than old ULIRG systems. If confirmed with larger samples, this behavior suggests that the late evolution of VLIRGs and ULIRGs is different.  5) Under the starburst scenario, the activity in the nuclei of the VLIRGs could be, in principle, explained by young bursts of less than $10^7$ years in age with a continuous rate of star formation of $\\sim$ 300 $M_\\odot$/yr), or with an instantaneous creation of $\\sim$ 10$^9M_\\odot$. However, the observed behavior, according to which old systems have lower luminosities and redder colors, is in better agreement with the evolution of an instantaneous starburst. Our data are also consistent with the idea that VLIRG activity tends to be dominated by starbursts, while a higher proportion of ULIRGs may evolve into a QSO-like object."
    },
    "0402/astro-ph0402579_arXiv.txt": {
        "abstract": "We report the results of an effort to measure the low frequency portion of the spectrum of the Cosmic Microwave Background Radiation (CMB), using a balloon-borne instrument called ARCADE (Absolute Radiometer for Cosmology, Astrophysics, and Diffuse Emission). These measurements are to search for deviations from a thermal spectrum that are expected to exist in the CMB due to various processes in the early universe.  The radiometric temperature was measured at 10 and 30~GHz using a cryogenic open-aperture instrument with no emissive windows. An external blackbody calibrator provides an {\\it in situ} reference. A linear model is used to compare the radiometer output to a set of thermometers on the instrument. The unmodeled residuals are less than 50~mK peak-to-peak with a weighted RMS of 6~mK. Small corrections are made for the residual emission from the flight train, atmosphere, and foreground Galactic emission. The measured radiometric temperature of the CMB is 2.721$\\pm .010$~K at 10~GHz and 2.694$\\pm 0.032$~K at 30~GHz. ",
        "introduction": "Since the discovery of the CMB a key question has been: How does it deviate from a perfect uniform black body spectrum? The FIRAS (Far InfraRed Absolute Spectrophotometer) instrument (Fixsen \\& Mather 2002, Brodd \\etal\\ 1997) has shown the spectrum is nearly an ideal black body spectrum from $\\sim60$~GHz to $\\sim600$~GHz with temperature 2.725$\\pm$.001 K. At lower frequencies the spectrum is not so tightly constrained; plausible physical processes (reionization, particle decay) could generate detectable distortions below 10 GHz while remaining undetectable by the FIRAS instrument. The ARCADE (Absolute Radiometer for Cosmology, Astrophysics and Diffuse Emission) experiment observes the CMB spectrum at frequencies a decade below FIRAS to search for potential distortions from a blackbody spectrum.  The frequency spectrum of the cosmic microwave background (CMB) carries a history of energy transfer between the evolving matter and radiation fields in the early universe. Energetic events in the early universe (particle decay, star formation) heat the diffuse matter which then cools via interactions with the background radiation, distorting the radiation spectrum away from a blackbody. The amplitude and shape of the resulting distortion depend on the magnitude and redshift of the energy transfer (Burigana \\etal\\ 1991, Burigana \\etal\\ 1995).  The primary cooling mechanism is Compton scattering of hot electrons against a colder background of CMB photons, characterized by the dimensionless integral \\begin{equation} y = \\int^z_0 ~\\frac{ k[T_e(z) - T_{\\gamma}(z)] }{ m_e c^2} \\sigma_T n_e(z) c \\frac{dt}{dz^\\prime} dz^\\prime, \\label{compton_y_definition} \\end{equation} of the electron pressure $n_e k T_e$ along the line of sight, where $m_e$, $n_e$ and $T_e$ are the electron mass, spatial density, and temperature, $T_{\\gamma}$ is the photon temperature, $k$ is Boltzmann's constant, $z$ is redshift, and $\\sigma_T$ denotes the Thomson cross section (Sunyaev \\& Zeldovich 1970). For recent energy releases $z < 10^4$, the gas is optically thin, resulting in a uniform decrement $ \\Delta T_{\\rm RJ} = T_{\\gamma} (1 - 2y) $ in the Rayleigh-Jeans part of the spectrum where there are too few photons, and an exponential rise in temperature in the Wien region with too many photons. The magnitude of the distortion is related to the total energy transfer \\begin{equation} \\frac{\\Delta {\\rm E}}{\\rm E} = {\\rm e}^{4y} - 1 \\approx 4y \\label{delta_e_vs_y_eq} \\end{equation} Energy transfer at higher redshift $10^4 < z < 10^7$ approaches the equilibrium Bose-Einstein distribution, characterized by the dimensionless chemical potential $ \\mu_0 = 1.4 \\frac{\\Delta {\\rm E}}{\\rm E}. $ Free-free emission thermalizes the spectrum at long wavelengths. Including this effect, the chemical potential becomes frequency-dependent, \\begin{equation} \\mu(x) = \\mu_0 \\exp(- \\frac{2x_b}{x}), \\label{mu_vs_freq_eq} \\end{equation} where $x_b$ is the transition frequency at which Compton scattering of photons to higher frequencies is balanced by free-free creation of new photons. The resulting spectrum has a sharp drop in brightness temperature at centimeter wavelengths (Burigana \\etal\\ 1991). A chemical potential distortion would arise, for instance, from the late decay of heavy particles produced at much higher redshifts.  Free-free emission can also be an important cooling mechanism. The distortion to the present-day CMB spectrum is given by \\begin{equation} \\Delta T_{\\rm ff} = T_{\\gamma} \\frac{Y_{\\rm ff}}{x^2} \\label{FF_distortion_eq} \\end{equation} where $x$ is the dimensionless frequency $h \\nu / k T_{\\gamma}$, $Y_{\\rm ff}$ is the optical depth to free-free emission \\begin{equation} Y_{\\rm ff} = \\int^z_0 ~\\frac{ k[T_e(z) - T_{\\gamma}(z)] }{ T_e(z) } \\frac{ 8 \\pi e^6 h^2 n_e^2 g }{ 3 m_e (kT_{\\gamma})^3 \\sqrt{6\\pi m_e k T_e} } \\frac{dt}{dz^\\prime} dz^\\prime, \\label{Yff_definition} \\end{equation} and g is the Gaunt factor (Bartlett \\& Stebbins 1991). The distorted CMB spectrum is characterized by a quadratic rise in temperature at long wavelengths. Such a distortion is an expected signal of the reionization of the intergalactic medium by the first collapsed structures.  \\begin{figure} \\includegraphics[angle=90,width=3.25in]{arcade_f1.ps} \\caption{Current 95\\% confidence upper limits to distorted CMB spectra. Measurements at short wavelengths (Fixsen \\etal\\ 1996) do not preclude detectable signals at wavelengths longer than 1 cm. \\label{arcade_vs_firas} } \\end{figure}  Figure \\ref{arcade_vs_firas} shows current upper limits to CMB spectral distortions. Measurements at wavelengths shorter than 1~cm are consistent with a blackbody spectrum, limiting $y < 14 \\times 10^{-6}$ and $\\mu < 9 \\times 10^{-5}$ at 95\\% confidence (Fixsen \\etal\\ 1996, Gush \\etal\\ 1990). Direct observational limits at longer wavelengths are weak. Reionization is expected to produce a cosmological free-free background with amplitude of a few mK at frequency 3 GHz (Haiman \\& Loeb 1997, Oh 1999). The most precise observations (Table \\ref{measure}) have uncertainties much larger than the predicted signal from reionization. Existing data only constrain $|Y_{\\rm ff}| ~< ~1.9 \\times 10^{-5}$, corresponding to temperature distortions $\\Delta T < 19 $ mK at 3 GHz (Bersanelli \\etal\\ 1994).  Uncertainties in previous measurements have been dominated by systematic uncertainties in the correction for emission from the atmosphere, Galactic foregrounds, or warm parts of the instrument. ARCADE represents a long-term effort to improve measurements at cm wavelengths using a cryogenic balloon-borne instrument designed to minimize these systematic errors. This paper presents the first results from the ARCADE program. ",
        "conclusions": "Since the sky data are compared directly to the external calibrator, systematic effects from the instrument are all eliminated to first order. Instead the burden falls on the external calibrator. The key issues are the emissivity (or blackness) of the calibrator, the accuracy of the thermometers and the relationship between the temperature of the thermometers and the temperature of the emitting surfaces.  The ARCADE external calibrator has been measured to be $>99.97$\\% emissive at 10~GHz using the flight horn antenna. The emissivity measurement was done at 295~K but the change from 295~K to 3~K should be small because the resistance changes less than 50\\% from 300~K to 1~K (Hemmati \\etal 1985). But the reflected radiation is almost entirely radiation from the horn which was at $\\sim2.8$~K. The residual uncertainty from this effect is 0.1~mK.  The remaining issue is the relationship between the temperature of the thermometers and the temperature of the emitting surface of the external calibrator. If there were no gradients the issue would vanish, but typically the gradient is 700~mK during the time of the observations. Some of these gradients are measured showing that the main gradient is from the front to the back of the external calibrator. The mean temperature of the emitting surface is modeled by a linear combination of the seven thermometers on the external calibrator. The changing temperature gradients themselves determine the fit to the combination of thermometers that best models the data. As long as the variations reflect the actual mean temperatures this is a good model. Tests fitting the sky temperature after dropping individual thermometers demonstrate that the calibrator has enough thermometers to adequately sense the thermal gradients.  The engineering tests done during this flight provide a base for the design and operation of the next ARCADE mission which will have 6 frequencies extending from 3~GHz to 90~GHz.  The ARCADE instrument has measured the radiometric temperature of the CMB to be ~2.721~K \\err\\ at 10 GHz, and 2.994~K$\\pm 32$~mK at 30~GHz."
    },
    "0402/astro-ph0402053_arXiv.txt": {
        "abstract": "Regular monitoring of the SMC with RXTE has revealed a huge number of X-ray pulsars. Together with discoveries from other satellites at least 45 SMC pulsars are now known.  One of these sources, a pulsar with a period of approximately 7.8 seconds, was first detected in early 2002 and since discovery it has been found to be in outburst nine times.  The outburst pattern clearly shows a period of 45.1 $\\pm$ 0.4 d which is thought to be the orbital period of this system. Candidate outburst periods have also been obtained for nine other pulsars and continued monitoring will enable us to confirm these. This large number of pulsars, all located at approximately the same distance, enables a wealth of comparative studies.  In addition, the large number of pulsars found (which vastly exceeds the number expected simply by scaling the relative mass of the SMC and the Galaxy) reveals the recent star formation history of the SMC which has been influenced by encounters with both the LMC and the Galaxy. ",
        "introduction": "The first known X-ray pulsar in the SMC was the persistent supergiant system SMC X-1. Two luminous transients (SMC X-2, SMC X-3) were discovered with SAS-3. \\cite{C78}. These were thought to be transient Be/neutron star systems although pulsations were not detected due to the low sensitivity of SAS-3. It was hypothesized that SMC pulsars were exceptionally luminous, possibly related to the low metallicity of the SMC. \\cite{Westerlund90} This was later to be disproved and an alternative explanation found for the high luminosity of the first few SMC X-ray pulsars to be discovered. Over the years a few pulsars were also found with satellites such as ROSAT. \\cite{Hughes94} ",
        "conclusions": ""
    },
    "0402/astro-ph0402265_arXiv.txt": {
        "abstract": "Recent astronomical observations suggest that the bulk of energy in the Universe is repulsive and appears like a dark component with negative pressure ($\\omega \\equiv p_x/\\rho_x < 0$). In this work we investigate thermodynamic and statistical properties of such a component. It is found that its energy and temperature grow during the evolution of the Universe since work is done on the system. Under the hypothesis of a null chemical potential, the case of phantom energy ($\\omega < -1$) seems to be physically meaningless because its entropy is negative. It is also proved that the wavelengths of the $\\omega$-quanta decrease in an expanding Universe. This unexpected behavior explains how their energy may be continuously stored in the course of expansion. The spectrum and the associated Wien-type law favors a fermionic nature with $\\omega$ naturally restricted to the interval $-1 \\leqslant \\omega < -1/2$. Our analysis also implies that the ultimate fate of the Universe may be considerably modified. If a dark energy dominated Universe expands forever, it will become increasingly hot. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0402/astro-ph0402609_arXiv.txt": {
        "abstract": "We elaborate further on the possibility that the inflationary primordial power spectrum contains superimposed oscillations. We study various effects which could influence the calculation of the multipole moments in this case. We also present the theoretical predictions for two other cosmological observables, the matter power spectrum and the $\\mathrm{EE}$ polarization channel. ",
        "introduction": "The possibility that the multipole moments $C_{\\ell }$ characterizing the angular distribution of the Cosmic Microwave Background Radiation (CMB) anisotropy on the celestial sphere possess superimposed oscillations has recently been investigated in Ref.~\\cite{MR1}. In that article, the superimposed oscillations originate from trans-Planckian wiggles~\\cite{tpl} in the primordial inflationary power spectrum but the study of Ref.~\\cite{MR1} was meant to be as independent as possible from the details of the underlying model. The presence of oscillations in the primordial spectrum has also been envisaged in Refs.~\\cite{osci}. Then, it has been demonstrated that the superimposed oscillations can cause a significant improvement of the fit to the first year Wilkinson Microwave Anisotropy Probe (WMAP) data~\\cite{wmap}, thanks to the presence of the cosmic variance outliers around the first Doppler peak. Moreover, it has also been shown that the corresponding drop in the $\\chi ^2$ is statistically significant according to the so called $F$ test. We have since received various inquiries about different effects that could modify the result obtained in Ref.~\\cite{MR1} like the influence of the splinning, the consideration of the lensing~\\cite{lim} and the way of estimating the statistical significance of the drop in the $\\chi ^2$. In this addendum, we clarify these issues and, in addition, present new results which are important for the completeness of Ref.~\\cite{MR1}, namely a fit with an improved $\\chi ^2$ which does not suffer from the back reaction problem and the prediction for two other cosmological observables, the matter power spectrum and the $\\mathrm{EE}$ polarization channel. Finally, we would like to emphasize that the general questions analyzed in this Brief Report are important irrespective of the data set used to study them. Therefore, although we use the first year WMAP data, the conclusions reached in this addendum are also valid for the future CMB data releases. ",
        "conclusions": "In this addendum, we have investigated various effects that could influence the presence of wiggles in the CMB multipole moments. We have shown that the splinning is an important parameter which, in the presence of high frequency superimposed oscillations, can influence the shape of the of CMB angular spectrum at intermediate scales. We have shown that the model presented in Ref.~\\cite{MR1} remains favored by the data (up to a variation of $\\Delta \\chi ^2\\simeq 2$). Two other cosmological observables have also been presented, namely the matter power spectrum and the $\\mathrm{EE}$ polarization multipole moments. With the available data, we have concluded that these two observables do not affect the results reached in Ref.~\\cite{MR1}.  \\par  A last remark is in order here. In Ref.~\\cite{MR1}, the statistical significance of the wiggles has been discussed by means of the $F$ test. This was motivated by the fact that the $F$ test only requires the knowledge of the likelihood maxima for the different models under consideration. However, by only varying the ``fast parameters'', one may expect to efficiently probe the corresponding subspace. Therefore, an interesting improvement of the present work would be to compute the statistical evidence in this subspace, thus providing us with a different test of the statistical significance of the wiggles."
    },
    "0402/astro-ph0402323_arXiv.txt": {
        "abstract": "Gravitational microlensing events with high peak magnifications provide a much enhanced sensitivity to the detection of planets around the lens star. However, estimates of peak magnification during the early stages of an event by means of $\\chi^2$ minimization frequently involve an overprediction, making observing campaigns with strategies that rely on these predictions inefficient.  I show that a rudimentary Bayesian formulation, incorporating the known statistical characteristics of a detection system, produces much more accurate predictions of peak magnification than $\\chi^2$ minimisation. Implementation of this system will allow efficient follow-up observing programs that focus solely on events that contribute to planetary abundance statistics. ",
        "introduction": "Following the suggestion of \\citet{Paczynski86}, several collaborations, notably MACHO and EROS, began to search for gravitational microlensing towards the Magellanic Clouds as an indicator of compact objects in the halo of the Milky Way \\citep{Alcock93, Aubourg93}. At about the same time, the OGLE collaboration began a survey in the direction of the Galactic bulge \\citep{Udalski92,Udalski93}. It was soon found that a much higher event rate occurred in fields towards the Galactic bulge relative to the rate towards the Magellanic Clouds \\citep{Udalski94a, Alcock95,Alcock97a}. Since 1990, approximately 1000 such events have been detected \\citep{Alcock00,Udalski00}.   Several groups including PLANET (Probing Lensing Anomalies NETwork, \\citealt{Albrow98,Albrow01,Dominik02,Gaudi02}), MPS (Microlensing Planet Search, \\citealt{Rhie99}) and $\\mu$FUN (Microlensing Follow-Up Network, \\citealt{Yoo03}) monitor events much more intensively than the survey groups in order to identify anomalous behavior that can signal the presence of a planet associated with the lens star. High-magnification events in particular (those with $A_{\\rm 0} \\gtrsim 10$) attract the attention of follow-up groups since it is these that are most likely to give detectable planetary signals \\citep{Griest98,Gaudi98}. In addition, for high magnification events the angular size of the source star may be non-negligible in comparison to the lens-source angular separation. In these cases the lightcurves of the events can provide the possibility to determine the lens-source relative proper motion \\citep{Gould94,Alcock97b} and atmospheric properties of the source \\citep{Heyrovsky03}.   In the first years of operation, when microlensing alerts came primarily from the MACHO collaboration, detected event rates were low enough that PLANET could monitor almost all potentially interesting events with ease. For the last two years (the 2002 and 2003 Bulge seasons), this has not been the case, due to the much improved alert rate since the advent of the OGLE III early warning system (EWS), \\url{http://www.astrouw.edu.pl/~ogle/ogle3/ews/ews.html} \\citep{Udalski94b,Udalski03}. In excess of 400 events were alerted by the EWS in each of these years. In addition, approximately 75 events were alerted in 2003 by the MOA collaboration \\citep{Bond02} although some of these were duplicates of EWS events.  We are now in an era in which a careful selection of events is necessary to optimize planet detection and exclusion productivity. For this reason, follow-up groups require accurate predictions of eventual maximum amplications in the early days following a detection. For the remainder of this paper I will focus exclusively on events detected by the OGLE III EWS. ",
        "conclusions": "High magnification events provide the best opportunity for detecting signals of planets around lens stars and for obtaining upper limits on their abundances. Intensive photometric monitoring programs are hampered currently by difficulties in identifying high magnification events well before peak. Systems that use $\\chi^2$ minimization to fit PSPL models to early data are prone to exagerated predictions of peak magnification. Such prections induce observers to spend their time monitoring events that ultimately have little statistical power.  I have shown here that a predictive system based on a Bayesian formalism that takes account of the characteristics of a detection system is immune to such behavior. Although such a Bayesian system tends to initially underpredict the peak for high magnification events, accurate prediction occurs as soon as sufficient data accumulate to justify the assertion. In all cases examined, this occurs well ahead of peak in their associated lightcurves and early enough for the events to be targeted for observation. Implementation of such a system based on the OGLE Early Warning System should result in much improved observing productivity for the 2004 season."
    },
    "0402/hep-ph0402200_arXiv.txt": {
        "abstract": "We study the $D$-term inflation scenario with a nonperturbative K\\\"ahler potential of the dilaton field. Although the FI term which leads an inflationary expansion is given by the derivative of the K\\\"ahler potential with respect to the dilaton in heterotic string models with anomalous $U(1)$, the too large magnitude is problematic for a viable $D$-term inflation. In this paper, we point out that the K\\\"ahler potential with a nonperturbative term can reduce the magnitude of FI term to desired values while both the dilaton stabilization and $D$-term domination in the potential are realized by nonperturbative superpotential. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0402/hep-th0402072_arXiv.txt": {
        "abstract": "{ Within the tachyon condensation approach, we find that a D$(p-2)$-brane is stable inside D$p$-branes when the bulk is compactified. It is a codimension-2 soliton of the D$p$-brane action with coupling to the bulk $(p-1)$-form RR field. We discuss the properties of such solitons. They may appear as detectable cosmic strings in our universe. }  \\begin{document} ",
        "introduction": "The evidence that the early universe has gone through an inflationary epoch has become very strong. In the brane world scenario, where the standard model fields (i.e., photons, electrons, quarks, etc., except graviton) are open string modes on branes, brane inflation \\cite{Dvali:1998pa} is quite natural. A particularly simple version of brane inflation involves the slow motion of a D3-brane towards an anti-D3-brane \\cite{Burgess:2001fx,Dvali:2001fw,Alexander:2001ks,Jones:2002cv, Buchan:2003gx}. Recently, based on a realistic superstring compactification where all moduli of the vacuum are stabilized \\cite{Giddings:2001yu,Kachru:2003aw}, it is shown that this D3-brane pair inflationary scenario may be realized in superstring theory \\cite{Kachru:2003sx,Silverstein:2003hf,Hsu:2003cy,Firouzjahi:2003zy}. Inflation ends as the D3-brane pair annihilates. Suppose the standard model fields live in a stack of (anti-)D3-branes. To allow the energy released to heat up the universe to start the hot radiation dominated big bang era \\cite{Shiu:2002xp,Cline:2002it}, the D3-brane pair annihilation should happen close to this stack of branes.  Towards the end of the inflationary epoch, cosmic strings are produced \\cite{Jones:2002cv,Sarangi:2002yt,Jones:2003da,Copeland:2003bj, Dvali:2003zh,Dvali:2003zj,Binetruy:2004hh,Halyo:2003uu}. In the above scenario, the annihilation of the D3-brane pair produces cosmic strings that are D1-strings (i.e., D1-branes) \\cite{Sen:1999mg}. The tension of such D1-strings are estimated to be around $G \\mu \\simeq 10^{-9}$ to $10^{-10}$ \\cite{Sarangi:2002yt,Copeland:2003bj}. If they survive long enough to evolve as a cosmic string network \\cite{Vilenkin:2000}, they will produce distinct signatures that should be detected in the near future, in particular by gravitational wave detectors such as LIGO II/Virgo or LISA \\cite{Damour:2001bk}.  If the D3-brane collides with a stack of anti-D3-branes and annihilates one of them, D1-strings are produced inside (or very close to) the (anti-)D3-branes \\cite{Sen:1999mg}. However, in string theory, D1-strings and D3-branes are not BPS with respect to each other. In fact, it is generally believed that the D1-strings will dissolve inside a D3-brane \\cite{Gava:1997jt,Hashimoto:1997gm}; that is, their energy will spread throughout the D3-brane. In this case, the cosmic strings would have dissolved almost immediately after they were formed and no observable signature will be left. So the (in)stability of such D1-strings inside a D3-brane is a very important phenomenological issue. Of course, this question in string theory is interesting in its own right.  The possibility of the stability of D1-strings inside a D3-brane was recently pointed out by Copeland, Myers and Polchinski (CMP) \\cite{Copeland:2003bj}. The coupling of the RR 2-form field $C_2$ to the Abelian gauge field $A_1$ inside the D3-brane is finite when the extra 6 dimensions are compactifed. This leads to spontaneous symmetry breaking and a D1-string inside a D3-brane becomes a topologically stable vortex (a D1-vortex) with localized energy (tension) density inside a D3-brane. However, this vortex is no longer BPS and has a net zero RR charge (but non-zero charge density) as measured by $C_2$ inside the D3-brane, since the winding number contribution to this charge is canceled by the magnetic flux contribution. Because of the conservation of the winding number, this same vortex becomes a D1-string when moved outside a D3-brane, as expected.  To justify the action used here and in CMP, we present a topological argument on the stability of this D1-vortex in the context of tachyon condensation, where a D1-string also appears as a vortex. Although our approach uses boundary superstring field theory, the D1-vortex stability is based on topological reasonings and so is insensitive to the details of the particular framework used. This approach also reveals the relation between a D1-string outside a D3-brane (a BPS vortex due to tachyon condensation) and a D1-vortex inside a D3-brane (a vortex due to the $C_2$ coupling). Although the dynamics is somewhat involved, this transition as a D1-string moves in/out a D3-brane is expected to be smooth. In the limit of vanishing $C_2$ coupling (e.g., some of the extra dimensions decompactfy), the size of such a D1-vortex grows to infinity. In effect, the magnetic flux spreads and a D1-string dissolves inside a D3-brane. It is the $C_2$ coupling that stabilizes the D1-string inside a D3-brane.  More generally, a D($p-2$)-brane is stable inside a stack of D$p$-branes. In a brane world where $(p-3)$ dimensions of the D$p$-branes are compactified while the remaining 3 dimensions span the universe, D$p$-anti-D$p$-brane inflation would generically create stable cosmic strings that are D($p-2$)-brane defects (with $(p-3)$ dimensions compactified). Their individual stability allows the cosmic string network evolution and implies that the detection of signatures of cosmic strings should be a good test of the brane inflationary scenario; furthermore, it gives an eagerly sought window to the superstring theory itself. Of course, the details of the phenomenology may be quite sensitive to the specific inflationary scenario.  The organization of this paper is as follows. In Sec. 2, we review the old argument why a D1-brane (i.e., a D1-string) was believed to dissolve inside a D3-brane. We also review the argument why there should be a domain wall enclosed by a D1-string. These two features constitute something of a puzzle, as pointed out in CMP. We then summarize our resolution to this puzzle. In Sec. 3, we solve analytically the vortex solution of the model involving the gauge field $A_1$ and $C_2$ inside a D3-brane. The $C_2$ coupling to $A_1$ leads to the Green-Schwarz mechanism where the gauge field becomes massive. We show that this model admits a topologically stable vortex solution, which is identified as a D1-vortex, that is, a D1-string inside a D3-brane, even though this vortex is not BPS. We show it has localized energy density, not spread throughout the D3-brane. However, the tension of this vortex is logarithmically divergent. In Sec. 4, we give the physical picture. By comparing to the Abelian Higgs model, we argue that the divergence in tension is expected in the approximate nature of the model, and is easily cured by adding the massive mode (namely the radial Higgs field associated to the axion) to the low energy effective supergravity action. In Sec. 5, we study this problem in the framework of tachyon condensation, where we give a topological argument why the D1-vortex is stable inside a D3-brane. We also use the boundary superstring field theory models to draw the connection between the two approaches. ",
        "conclusions": "In this paper we have shown that there exists a consistent vortex solution of the D3-brane action with RR 2-form bulk terms, provided the $C_2$ coupling is non-zero. This requires the compactification of the extra dimensions, as is the case in any brane world scenario. This vortex solution is not BPS, but it has a localized energy density, not spread throughout the D3-brane. As we pull this vortex out of a D3-brane, it becomes a BPS D1-string; that is, the D1-string is never broken, although its properties do change as it moves inside/outside a D3-brane. The transition from a D1-string to a D1-vortex may be followed in tachyon condensation.  In terms of tachyon condensation, all D$p$-branes (maybe with the exception for $p=9$) may be viewed as solitonic defects. So there are only odd $p$-defects in Type IIB theory. If a D1-string is a boundary of a domain wall, or if a D1-string breaks so that it has ends, then such domain walls or string ends will appear as even dimensional defects (2-dim. and 0-dim. respectively), which presumably should not exist in Type IIB theory. The resolution proposed in this paper is consistent with this belief.  It will be interesting to study the interaction and intercommutation properties of these D1-vortices. In cosmology, these defects may become cosmic strings in our universe. We need to determine their properties to find any distinct signature that would differentiate them from the usual quantum field theory vortices. The importance of this study cannot be overstated since this is, at the moment, a most promising window into superstring theory and the inflationary scenario."
    },
    "0402/astro-ph0402445_arXiv.txt": {
        "abstract": "In previous work we found that the spectral state switch and other spectral properties of both neutron star (NS) and galactic black hole candidates (GBHC), in low mass x-ray binary systems could be explained by a magnetic propeller effect that requires an intrinsically magnetic central compact object. In later work we showed that intrinsically magnetic GBHC could be easily accommodated by general relativity in terms of magnetospheric eternally collapsing objects (MECO), with lifetimes greater than a Hubble time, and examined some of their spectral properties. In this work we show how a standard thin accretion disk and corona can interact with the central magnetic field in atoll class NS, and GBHC and active galactic nuclei (AGN) modeled as MECO, to produce jets that emit radio through infrared luminosity $L_R$ that is correlated with mass and x-ray luminosity as $L_R \\propto M^{0.75 - 0.92}L_x^{2/3}$ up to a mass scale invariant cutoff at the spectral state switch. Comparing the MECO-GBHC/AGN model to observations, we find that the correlation exponent, the mass scale invariant cutoff, and the radio luminosity ratios of AGN, GBHC and atoll class NS are correctly predicted, which strongly implies that GBHC and AGN have observable intrinsic magnetic moments and hence do not have event horizons. ",
        "introduction": "In earlier work (Robertson \\& Leiter 2002, hereafter RL02) we extended analyses of magnetic propeller effects (Campana et al. 1998, Zhang, Yu \\& Zhang 1998) of neutron stars (NS) in low mass x-ray binaries (LMXB) to the domain of GBHC. From the luminosities at the low/high spectral state transitions, accurate rates of spin were found for NS and accurate quiescent luminosities were calculated for \\textit{both} NS and GBHC. The NS magnetic moments were found to be consistent with $\\sim 10^{8 - 9}$G magnetic fields, in good agreement with those others have found (e.g. Bhattacharya 1995) from spin-down rates for similarly spinning 200 - 600 Hz millisecond pulsars. GBHC spins were found to be typically 10 - 50 Hz. Their magnetic moments of $\\sim 10^{29}$ gauss cm$^3$ are $\\sim 200$ times larger than those of `atoll' class NS (e.g. Burderi et al. 2002, DiSalvo \\& Burderi 2003). The implied magnetic fields of GBHC are in good agreement with fields of $\\sim 10^8$G that have been found at the base of the jets of GRS 1915+105 (Gliozzi, Bodo \\& Ghisellini 1999, Vadawale, Rao \\& Chakrabarti 2001) and in the accretion disk of Cygnus X-1 (Gnedin et al. 2003). At accretion disk inner radii corresponding to the low/high spectral state switch, the magnetic fields of both `atoll' class NS and GBHC are $\\sim 5\\times10^7$G, which may account for some of their strong similarities (e.g. Yu et al. 2003, Tanaka \\& Shibazaki 1996, van der Klis 1994).  In later work (Leiter \\& Robertson 2003, Robertson \\& Leiter 2003, hereafter RL03) we have described how the Einstein field equations of General Relativity applied to compact plasmas with equipartition magnetic fields permit the existence of magnetic, eternally collapsing objects (MECO) that can have lifetimes in excess of a Hubble time. These highly redshifted, faintly (as distantly observed in quiescence) radiating objects can produce `ultrasoft' thermal spectral peaks and the magnetic propeller effects found in RL02. Here we examine the accretion disk - magnetosphere interaction and show how the magnetosphere can drive jets. Our model should be applicable for any jet producing objects with sufficiently large magnetic moments, whether T-Tauri stars or NS, or GBHC and active galactic nuclei (AGN) modeled as MECO. In this context, the scaling of the magnetic moments of MECO with mass will be an important consideration. As shown in RL03, the MECO is dominated by a photon-photon collision generated pair plasma which is stabilized at high redshift deep inside the photon orbit by an Eddington limit radiation pressure generated by an equipartition magnetic field intrinsic to the MECO. The surface value of the MECO intrinsic magnetic field is calculated by equating the synchrotron generated photon pressure ($\\propto B^4$) to the gravitational force per unit area, which is proportional to the density. Since the density is inversely proportional to the square of the MECO mass, $M$, the internal magnetic field scales as $M^{-1/2}$ and the MECO magnetic moment, $\\mu$, scales as $M^{-1/2}(2GM/c^2)^3 \\propto M^{5/2}$.  In the following, for NS, GBHC and AGN, we will assume the existence of a gas pressure dominated, geometrically thin accretion disk (Shakura \\& Sunyaev 1973). For gas pressure dominance, it has been shown (e.g., Merloni \\& Fabian 2002) that the hard x-ray spectral tail and reflection features of the low state spectrum can be adequately explained by reprocessing of the soft thermal disk photons in an accretion disk corona (ADC). The physical size of a corona is consistent with limits found for the source of the power-law x-ray emissions of LMXB (Church \\& Baluci\\'{n}ska-Church 2003).  It has been suggested, however, (Markoff, Falcke \\& Fender 2001, Falcke, K\\\"{o}rding \\& Markoff 2003) that the power-law x-ray emissions might originate in a jet. Flat or inverted spectrum synchrotron radio- infrared emissions are generally believed to originate in jets and low state jets have been resolved (Stirling et al. 2001) and studied over a wide range of luminosity variation (Corbel et al. 2000, 2003). As a result of these outflows, it has been pointed out (Fender, Gallo \\& Jonker 2003) that the low quiescent luminosities of GBHC cannot be taken as evidence of advective accretion flows (ADAF) through event horizons and as noted by Abramowicz, Kluzniak and Lasota (2002) there is presently no other observational evidence of event horizons.  Whether or not the x-rays originate in the jet, there is a strong coupling between x-ray and radio emissions that must be related to the accretion flow and jet structure. A universal low state radio / X-ray correlation ($L_R \\propto L_x^{0.7}$) (Gallo, Fender \\& Pooley 2003) with a cutoff at the low/high state transition (Fender et al. 1999, Tannenbaum et al. 1972, Corbel et al. 2003) has been found for GBHC and NS ( Migliari et al. 2003). A similar radio / x-ray correlation (Merloni, Heinz \\& Di Matteo 2003, Falcke, K\\\"{o}rding \\& Markoff, 2003) and its suppression at the transition to the high/soft state (Maccarone, Gallo \\& Fender 2003) have been shown to hold for AGN as well. These radio / X-ray luminosity correlations have been examined for scale invariant jets (Heinz \\& Sunyaev 2003, hereafter HS03), yielding constraints on the accretion processes. In the context of HS03, Merloni, Heinz \\& Di Matteo (2003) (Hereafter MHD03) have examined their correlation for compatibility with various accretion flow models and found better consistency with an ADAF / jet model than with radiatively efficient disk / jet or pure jet models.  An ADAF/jet model (Meier 2001) can also account for the low/high spectral state transition as a transition from an ADAF to a standard thin disk. It relies on a rapid black hole spin to provide energy to drive the jet. The model predicts that stable high/soft states would not exist for AGN more massive than $7 \\times 10^4 M_\\odot$ (Meier 2001) or, with more generous allowance for hysteresis effects, $\\sim 4\\times 10^6 M_\\odot$ (Maccarone, Gallo \\& Fender 2003). The theoretical mass limit occurs because the Eddington scaled luminosity at which a thin disk (constrained to match a radiatively inefficient ADAF accretion rate) becomes radiation dominated is mass dependent. Since the high/soft state nevertheless appears to exist in AGN more massive than $6 \\times 10^7 M_\\odot$, the ADAF transition model cannot be regarded as established. Understanding the origin of the mass limit error of the model remains an open question (Maccarone, Gallo \\& Fender 2003).  Black hole models that rely entirely on the jet to produce the power-law x-ray emissions may have difficulties with constraints on the physical size of a jet.  For dipping sources the size of the region of the low state power-law production has been found (Church \\& Baluci\\'{n}ska-Church 2003) to be $\\leq 10^9$ cm. In addition, there is the enigma of the size of the hard spectral producing region increasing while the jet dies in the high state. It is also unclear how black hole and NS behaviours could be so similar with the magnetic fields of even the weakly magnetized atoll class NS being capable of disrupting the inner accretion disk. On the other hand, we will show that our MECO model, with a radiatively efficient disk, will provide a superior fit to the radio / X-ray correlations and provide a mass scale invariant cutoff at the high/soft state transition while permitting radio-infrared and some of the x-ray luminosity to originate in a jet. ",
        "conclusions": "In a previous paper (RL02) we found that the spectral state switch and other spectral properties of low mass x-ray binaries, including both NS and GBHC, could be explained by a magnetic propeller effect that requires an intrinsically magnetized central object. Subsequently (RL03) we applied the Einstein field equations of General Relativity to the case of a highly compact, Eddington limited, pair dominated plasma with an intrinsic equipartition magnetic field. We found that the Einstein equations permit the existence of intrinsically magnetic, highly red shifted, extremely long lived, collapsing, radiating MECO objects that can produce the required propeller effects. In addition to accounting for the strong spectral similarities of NS and GBHC, the magnetosphere-accretion disk interaction associated with the MECO model has provided explanations for radio / x-ray luminosity correlations, the mass scale invariant spectral state switch phenomenon with its suppression of the radio jet outflow in the high/ soft state, the \"ultrasoft\" thermal peak and hard spectral tail of the high state, and, finally, the quiescient luminosities described as spin-down driven radiations.  In conclusion, we have shown here how a standard, thin, gas pressure dominated accretion disk and corona can interact with the central intrinsic magnetic moments of MECO-GBHC/AGN and NS in x-ray biniaries to drive low state jets. In the case of the MECO-GBHC/AGN the radio-infrared emissions of the jets have been found to correlate with the x-ray luminosity up to a mass scale invariant cutoff $L_c / L_{Edd}$ at the spectral state switch. In this context we obtained radio-infrared luminosities for MECO that vary as $M^{0.75-0.92}L_x^{2/3}$, consistent with observations of GBHC and AGN, and correctly predicted the observed relative radio luminosities of NS, GBHC, and AGN. While much detailed work remains to be done, the successful comparison of the MECO model predictions with observations strongly suggests that GBHC and AGN may have observable intrinsic magnetic moments anchored within them and hence they do not have event horizons.\\\\  \\noindent {\\bf Acknowledgements}\\\\ We thank the anonymous referee for many comments and suggestions that have substantially improved this paper. We thank Elena Gallo for providing data for Figure 1. Useful information has been generously provided by Mike Church, Heino Falcke and Thomas Maccarone. We are very grateful to Abhas Mitra for many helpful discussions of gravitational collapse and pertinent astrophysical observations."
    },
    "0402/astro-ph0402390_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\subsubsection*{Observations} There are several hints indicating that satellite galaxies orbiting within our own Milky Way are interacting with each other.  Zhao (1998), for instance, proposed a scenario where the Sagittarius Dwarf galaxy had an encounter with the Magellanic Cloud system some 2--3 Gyrs ago, something that has also been speculated and noted by Ibata~\\& Lewis (1998).  Moreover, the two Magellanic Clouds themselves are another example of an interacting pair of substructure galaxies. It has also been noted by Moore~\\ea (1996) that ``galaxy harrasment'' in cosmological simulations of galaxy cluster evolution will lead to a morphology change of satellite galaxies.  However, the literature to date lacks a statistical analysis of interacting satellite galaxies orbiting within the potential of a common dark matter host halo. How frequent are satellite-satellite encounters and where in the galaxy cluster do they happen? Furthermore, observations of the Local Group Dwarfs indicate a clear correlation between star formation activity and the distance of the respective Dwarf to the centre of the Milky Way (van den Bergh 1994) with satellites farther away showing stronger activity. Can this be ascribed to satellite-satellite interactions? The aim of this study is to quantify such interactions in galaxy clusters derived from fully self-consistent cosmological \\nbody\\ simulations within the framework of the currently accepted Cold Dark Matter (CDM) structure formation scenario.  \\subsubsection*{Is Cold Dark Matter still feasible?} There is mounting, if not overwhelming, evidence that CDM provides the most accurate description of our Universe. Observations point towards a \\LCDM\\ Universe comprised of 28\\% dark matter, 68\\% dark energy, and luminous baryonic matter (i.e. galaxies, stars, gas, and dust) at a mere 4\\% (cf. Spergel~\\ea 2003). This so-called ``concordance model'' induces hierarchical structure formation whereby small objects form first and subsequently merge to form progressively larger objects (e.g. White \\& Rees 1978; Davis \\ea 1985).  Hence, galaxies and galaxy clusters are constantly fed by an accretion stream of smaller entities starting to orbit within the encompassing dark matter potential of the host. While generally successful, the \\LCDM\\ model does face several problems, one such problem actually being the prediction that one-to-two orders of magnitude more satellite galaxies should be orbiting within galactic halos than are actually observed (Klypin~\\ea 1999; Moore~\\ea 1999).  However, there are also indications that the CDM model is in fact correct and does \\textit{not} have a problem with an overabundant population of satellite galaxies. For instance, Benson~\\ea (2002) carried out a semi-analytical study of satellites in the Local Group and found that an earlier epoch of reionisation was sufficient to suppress star formation in many of the subhalos and thus produce a significant population of ``dark galaxies''.  Therefore, if the CDM model is in fact correct and the (overabundant) population of (dark) satellites predicted by it really does exist, it is imperative to understand the discrepancy by investigating the orbital evolution of these objects and their deviation from the background dark matter distribution.  \\subsubsection*{The story, so far} To date, typical satellite properties such as orbital parameters and mass loss under the influence of the host halo have primarily been investigated using \\textit{static} potentials for the dark matter host halo (Johnston \\ea 1996; Hayashi \\ea 2003).  We stress that each of these studies have provided invaluable insights into the physical processes involved in satellite disruption; our goal is to augment those studies by relaxing the assumption of a static host potential as, in practice, realistic dark matter halos are neither static nor spherically symmetric.  \\subsubsection*{The story continues} The work presented here is based upon a set of numerical simulations of structure formation within said concordance model, analysing in detail the temporal and spatial properties of satellite galaxies residing within host dark matter halos that formed fully self-consistently within a cosmological framework. We focus on interactions between satellite galaxies orbiting within a larger dark matter halo and especially if there is a relation between mutual interplay and distance to the host.  The outline of the paper is as follows. In \\Sec{Identify} we present our new halo finding algorithms based upon the \\nbody\\ code \\mlapm.  We then apply it to our set of eight cosmological dark matter halos in \\Sec{Application} with a summary of ours results given in \\Sec{Summary}. ",
        "conclusions": "\\label{Summary} We used a set of eight high-resolution cosmological simulations to investigate and quantify interactions between satellite galaxies orbiting within a common dark matter halo. Using our definition for encounter, which is based upon the mutually induced tidal radius, we showed that on average 30\\% of the substructure population had had more than one encounter per orbit with another satellite galaxy orbiting within the same host halo. There is, however, a clear trend for interactions to be more common in young galaxy clusters.  We furthermore showed that satellite galaxies closer to the centre of the host halo had had more interactions with companion satellites, not because they simply orbited for longer in the underlying host potential but most likely because of the universal radial distribution of satellite galaxies found in cosmological dark matter halos (Gill~\\ea 2004). Even though satellite-satellite interactions are unimportant for the majority of satellite galaxies, there exists a sub-population for which this needs to be investigated in more detail and more carefully, respectively.  We also noted that there is a degeneracy between the influence of the host halo and the interactions with the companion satellites which can only be disentangled with an appropriate resolution for both the actual \\nbody-simulation and the halo finding technique. We therefore applied a new method for identifying gravitationally bound objects in cosmological \\nbody\\ simulations. This new technique is based upon the adaptive grid structures of the open source adaptive mesh refinement code \\mlapm\\ (Knebe, Green~\\& Binney 2001). The halo finder is called \\mhf\\ and acts on the same accuracy level as the actual simulation. A more thorough study of the functionality of \\mhf\\ is presented in Gill, Knebe~\\& Gibson (2004a). A detailed analysis of the degeneracy between influence of the host halo and interactions with companion satellites can be found in a companion paper, too (Gill, Knebe~\\& Gibson 2004b)."
    },
    "0402/astro-ph0402359_arXiv.txt": {
        "abstract": "Two years of microwave background observations with the Cosmic Background Imager (CBI) have been combined to give a sensitive, high resolution angular power spectrum over the range $400 < \\ell < 3500$. This power spectrum has been referenced to a more accurate overall calibration derived from the {\\it Wilkinson Microwave Anisotropy Probe}.  The data cover $90 \\, {\\rm deg^2}$ including three pointings targeted for deep observations.  The uncertainty on the $\\ell >2000$ power previously seen with the CBI is reduced. Under the assumption that any signal in excess of the primary anisotropy is due to a secondary Sunyaev-Zeldovich anisotropy in distant galaxy clusters we use CBI, Arcminute Cosmology Bolometer Array Receiver, and Berkeley-Illinois-Maryland Association array data to place a constraint on the present-day rms mass fluctuation on $8 \\, h^{-1} \\, {\\rm Mpc}$ scales, $\\sigma_8$.  We present the results of a cosmological parameter analysis on the $\\ell < 2000$ primary anisotropy data which show significant improvements in the parameters as compared to {\\it WMAP} alone, and we explore the role of the small-scale cosmic microwave background data in breaking parameter degeneracies. ",
        "introduction": "\\label{sec:intro}  The Cosmic Background Imager (CBI) is a planar synthesis array designed to measure cosmic microwave background (CMB) fluctuations on arcminute scales at frequencies between $26$ and $36$ GHz.  The CBI has been operating at its site at an altitude of 5080 m in the Chilean Andes since late 1999. Previous results have been presented by \\citet{Padin01}, \\citet{Mason03}, and \\citet{Pearson03}. The principal observational results of these papers were: (i) the first detection of anisotropy on the mass scale of galaxy clusters---thereby laying a firm foundation for theories of galaxy formation; (ii) the clear delineation of a damping tail in the power spectrum, best seen in the mosaic analysis of \\citeauthor{Pearson03}; (iii) the first determination of key cosmological parameters from the high-$\\ell$ range, independent of the first acoustic peak; and (iv) the possible detection, presented in the deep field analysis of \\citeauthor{Mason03}, of power on small angular scales in excess of that expected from primary anisotropies.  The interpretation of these results has been discussed by \\citet{Sievers03} and \\citet{Bond04}.  The CBI data, by virtue of their high angular resolution, were able to place constraints on cosmological parameters which are largely independent of those derived from larger-scale experiments; for instance,  10\\% measurements of $\\Omega_{\\rm tot}$ and $n_s$ using only CBI, DMR and a weak $H_0$ prior.  The small-scale data also play an important role in improving results on certain key parameters ($\\Omega_b h^2$, $n_S$, $\\tau_C$) which are less well-constrained by large-scale data.  Theoretical models predict the angular power spectrum of the CMB \\begin{equation} C_{\\ell} = \\langle |a_{\\ell m}|^2 \\rangle \\end{equation} where the $a_{\\ell m}$ are coefficients in a spherical harmonic expansion of temperature fluctuations in the CMB, $\\Delta T/T_{\\rm CMB}$, where $T_{\\rm CMB} \\approx 2.725$ K is the mean temperature of the CMB, and the angle brackets denote an ensemble average.  These theories also predict a series of acoustic peaks in the angular power spectrum on scales $\\lesssim 1^{\\circ}$ ($\\ell \\gtrsim 200$), and a decline in power towards higher $\\ell$ due to photon viscosity and the thickness of the last scattering surface. Early indications of the first acoustic peak were presented by \\citet{Miller99}; definitive measurements of the first and second peaks were reported by \\citet{boomnature}, \\citet{Lee01}, \\citet{Netterfield02}, \\citet{Halverson02}, \\citet{Scott03}, and \\citet{Grainge03}\\footnote{In the parameter analysis of \\S4 we use the latest VSA data \\citep{dickinson}, which was released shortly after this paper was first submitted.}. The last of these experiments reached $\\ell \\sim 1400$.  The CBI \\citep{Padin02} has complemented these experiments by covering an overlapping range of $\\ell$ extending to $\\ell \\sim 3500$. The Arcminute Cosmology Bolometer Array Receiver (ACBAR) \\citep{kuo} has recently covered a similar range of $\\ell$ as the CBI at higher frequency; the Berkeley-Illinois-Maryland Association array (BIMA) has also made 30 GHz measurements at $\\ell \\sim 5000$ which probe the secondary Sunyaev-Zeldovich effect (SZE) anisotropy \\citep{dawson}.  These experiments---which employ a wide variety of instrumental and experimental techniques---present a strikingly consistent picture which supports inflationary expectations (see \\citealt{bondproc} for a review).  However the results at intermediate angular scales ($500 < \\ell < 2000$) currently have comparatively poor $\\ell$--space resolution, and the high-$\\ell$ results are difficult to compare conclusively owing to the low signal-to-noise ratio ($\\sim 2$--$4$).  The results presented here improve the situation by: (i) expanding the coverage of the CBI mosaics for higher $\\ell$ resolution, (ii) integrating further on the deep fields, and (iii) combining the deep and mosaic data for a single power spectrum estimate over the full range of $\\ell$ covered by the CBI.  The CBI results presented by \\citet{Mason03} and \\citet{Pearson03} were based on data obtained between January and December of 2000.  \\citeauthor{Mason03} analyzed the data resulting from extensive integration on three chosen ``deep fields'' to constrain the small-scale signal; the analysis of \\citeauthor{Pearson03} used data with shallower coverage of a larger area (``mosaics'') to obtain better Fourier-space resolution.  Further observations were conducted during 2001; these were used to extend the sky coverage of the mosaics in order to attain higher resolution in $\\ell$, and to go somewhat deeper on the existing deep fields. This paper presents the power spectrum resulting from the combination of the full CBI primary anisotropy dataset, which comprises data from years 2000 and 2001 on both mosaic {\\it and} deep fields.  Two of the mosaic fields (14\\,h and 20\\,h) include deep pointings; there is also a third deep pointing (08\\,h), and a third mosaic (02\\,h).  The CBI data have been recalibrated to a more accurate power scale derived from the {\\it Wilkinson Microwave Anisotropy Probe} ({\\it WMAP}).  The organization of this paper is as follows.  In \\S~\\ref{sec:obscalib} we discuss the observations and {\\it WMAP}-derived recalibration.  In \\S~\\ref{sec:dataanalys} we present images and power spectra derived from the data and explain the methodology employed in their derivation.  In \\S~\\ref{sec:interp} we use these results to constrain cosmological parameters based on standard models for primary and secondary CMB anisotropies.  We present our conclusions in \\S~\\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  The CBI power spectrum is compared with {\\it WMAP} and ACBAR results in Figure~\\ref{fig:cbi_wmap_acbar_spectrum}.  These results, together with those from a host of other ground- and balloon-based experiments in recent years, are consistent with the key predictions of structure formation and inflationary theories: The universe is close to flat; the initial spectrum of perturbations is nearly scale invariant; oscillations and damping in the power spectrum evince the expected signatures of sub-horizon scale causal processes; initial conditions are Gaussian, and are consistent with adiabatic fluctuations; and the magnitude of fluctuations from the largest scales down to galaxy cluster scales is consistent with what is needed to produce locally observed structures through gravitational collapse.  For discussion of these points see \\citet{bondproc}, \\citet{peiris}, and references therein.  The concordance of observational results with theoretical expectations has permitted cosmological parameters to be determined with precision.  In this work we obtain: $\\Omega_b h^2=0.0225^{+0.0009}_{-0.0009}$, $\\Omega_c h^2=0.111^{+0.010}_{-0.009} $, $\\Omega_{\\Lambda}= 0.74^{+0.05}_{-0.04} $, $\\tau_C= 0.11^{+0.02}_{-0.03} $ , $n_s= 0.95^{+0.02}_{-0.02} $, $t_0=13.7^{+ 0.2}_{- 0.2} $ Gyr, and $\\Omega_m=0.26^{+0.04}_{-0.05} $ from a selection of current primary anisotropy data including CBI, {\\it WMAP}, ACBAR, and Boomerang, and using the flat plus weak-$h$ prior (see Table~\\ref{tab:cmbonly}).  Similar results are obtained when large-scale structure priors are incorporated (Table~\\ref{tab:cmblss}).  As discussed in \\S~\\ref{sec:interp} a flat prior (i.e., assumption that $\\Omega_{\\rm tot} = 1$) is imposed on most of our parameter analysis; while supported by observational data this does impose a strong constraint, and some parameter estimates would be less accurate without it.  A marginal detection of a running scalar spectral index remains, and is consistent with that presented by \\citet{spergel03}.  As discussed in \\S~\\ref{sec:params}, the addition of CMB data from $600 < \\ell < 2000$ significantly improves constraints on $\\Omega_b h^2$, $n_s$, the amplitude of the primary anisotropies, the age of the universe, and $\\tau_C$ relative to what is obtained with only large-scale CMB data (see Figure~\\ref{fig:paramcurves}).  In the absence of a restrictive $\\tau_C$ prior the $\\ell < 600$ data leave significant degeneracies which are broken by the higher-$\\ell$ experiments  (see Figures~\\ref{fig:moneyshot} and \\ref{fig:moneyshot1}).  We note that the improvement between the ``CBI+{\\it WMAP}'' and ``CBI+ALL'' cases comes primarily from the addition of the Boomerang data. Improvements are also seen in analyses which allow a running scalar spectral index (Figure~\\ref{fig:s8alpha}).  The tight constraint on the baryon density, $\\Omega_b h^2 = 0.0225^{+0.0009}_{-0.0009}$ compares favorably with observationally determined BBN values of $\\Omega_b h^2 = 0.0214 \\pm 0.0020 $ \\citep[][]{kirkman}.  We have also obtained an accurate measurement of $n_s$ from the CMB data only, $n_s=0.95 \\pm 0.02$. These results are robust with respect to prior assumptions, such as flatness, imposed on the analysis. By way of comparison the {\\it WMAP}-only values for these parameters are $\\Omega_b h^2=0.0243\\pm0.0016$ and $n_s=1.01 \\pm 0.05$. The breaking of these degeneracies largely relies on the ratio of power levels on small angular scales to those on large angular scales, so the precision of these results has benefited from the accurate cross-calibration with {\\it WMAP}. The CBI data also favor a negative running scalar spectral index $\\alpha_s = -0.087\\pm 0.028$ (CBI+ALL+LSS), consistent with the results from {\\it WMAP} combined with LSS constraints  \\begin{figure*} \\plotone{f14.eps} % \\caption{The CBI+{\\it WMAP}+ACBAR Spectrum + high $\\ell$ points from BIMA. The curves at high $\\ell$ show the levels of SZ power expected in representative models using moving mesh hydrodynamics simulations (dotted) and smooth particle hydrodynamics (dashed) simulations (see text).  The green and pink curves correspond to 30 GHz and 150 GHz, respectively.  In these simulations $\\sigma_8^{\\rm SZ}=0.98$, which also fits well the WMAP and CBI observations at lower $\\ell$ for the case of a running spectral index (see Table 5). The highest-$\\ell$ ACBAR point has been displaced slightly to lower $\\ell$ for clarity. } \\label{fig:cbi_wmap_acbar_spectrum_2} \\end{figure*}  In Figure~\\ref{fig:cbi_wmap_acbar_spectrum_2} we show the same data as plotted in Figure~\\ref{fig:cbi_wmap_acbar_spectrum}, now on a log-log plot and with additional curves which show the expected level of SZE power for the two sets of simulations discussed by \\citet{Bond04}. Note that the fortuitous ``agreement'' between the CBI and ACBAR power levels at $\\ell>2000$ is not expected if the power has a significant component due to the Sunyaev-Zel'dovich Effect because of the different observing frequencies.  Nevertheless, given the uncertainties in these two measurements, it can be seen that the models span a range of power at high $\\ell$ which fits both the CBI and ACBAR observations.  The detection of power at $\\ell > 2000$ is consistent with the results presented by \\citet{Mason03}, although somewhat lower.  We find a bandpower $355^{+137}_{-122} \\, {\\rm \\mu K^2}$ ($68\\%$ confidence, including systematic contributions).  By combining this result with high-$\\ell$ results from BIMA and ACBAR we detect power in excess of that expected from primary anisotropy at $98\\%$ confidence.  This result includes a marginalization over expected primary anisotropy power levels.  Assuming the signal in excess of expected primary anisotropy is due to a secondary SZ foreground we determine $\\sigma_8^{\\rm SZ} = 0.96^{+0.06}_{-0.07} \\, (68\\%)$.  The lower confidence level of the detection of an excess, and also the smaller values of $\\sigma_8^{\\rm SZ}$, are chiefly due to the lower high-$\\ell$ bandpower we obtain and the inclusion of the uncertainty in the primary anisotropy bandpower at $\\ell > 2000$.  The strong dependence of the observable power on $\\sigma_8$ gives rise to firm upper limits on $\\sigma_8$ but a tail to low values (Figure~\\ref{fig:s8sz}). It should be borne in mind that there are systematic uncertainties in the theoretical prediction of the power spectrum due to secondary SZ anisotropies which correspond to a $10\\%$ systematic uncertainty in $\\sigma_8$.  An appreciable fraction of CBI data were rejected by vetoing NVSS sources, and furthermore the uncertainty in the power level of the source population remaining after the NVSS veto is a limiting factor at $\\ell > 2000$.  In late 2004 a sensitive, wideband continuum receiver will be commissioned on the Green Bank Telescope (GBT) to deal with both of these issues.  This will result in a more sensitive determination of the total intensity power spectrum at all $\\ell$ covered by the CBI. Since the end of the observations reported here, the CBI was upgraded and dedicated to full-time polarization observations."
    },
    "0402/astro-ph0402503_arXiv.txt": {
        "abstract": "The acceleration of the expansion of the universe arises from unknown physical processes involving either new fields in high energy physics or modifications of gravitation theory.  It is crucial for our understanding to characterize the properties of the dark energy or gravity through cosmological observations and compare and distinguish between them.  In fact, close consistencies exist between a dark energy equation of state function $w(z)$ and changes to the framework of the Friedmann cosmological equations as well as direct spacetime geometry quantities involving the acceleration, such as ``geometric dark energy'' from the Ricci scalar.  We investigate these interrelationships, including for the case of superacceleration or phantom energy where the fate of the universe may be more gentle than the Big Rip. ",
        "introduction": "\\label{sec.intro}  The acceleration of the expansion of the universe poses a fundamental challenge to the standard models of both particle physics and cosmology.  In both cases addition of an unknown physical component, called dark energy, or modification of gravitation, possibly arising from extra dimensions, is required. Most attention has been paid to dark energy as a high energy scalar field, a physical component contributing a presently dominating energy density, characterized by a time varying equation of state. But acceleration is fundamentally linked to gravitation through the Principle of Equivalence and changes to the framework of the Friedmann cosmological equations governing the universal expansion would play a natural role.  Observations from next generation cosmological probes will map the expansion history $a(t)$ at 1\\% precision, offering the possibility of characterizing the physics responsible for the acceleration. This can be used to test specific models inspired by unified physics involving string theory, supergravity, extra dimensions (e.g.\\ braneworlds), or scalar-tensor gravity, say.  Alternately, one can derive general parametrized constraints on the expansion history and propagate these through into quantities such as an effective dark energy equation of state, extra terms in the Friedmann equations, or spacetime geometry characteristics.  Not only the magnitude of the constraints but the interpretation of them is important.  We investigate to what extent one can use a common parametrization to describe these very different areas of new physics, and conversely how they can be distinguished.  In \\S\\ref{sec.w} we briefly review dark energy as a scalar field component of the universe.  A general modification of the Friedmann equation is analyzed in \\S\\ref{sec.dh}.  We examine in \\S\\ref{sec.rh} the fundamental and general relation between acceleration and spacetime geometry, specifically involving the Ricci scalar, to motivate modifications of gravitation as a possible source of the acceleration -- ``geometric dark energy''. In \\S\\ref{sec.super} we address the issue of superacceleration and whether this leads to a Big Rip.  We conclude in \\S\\ref{sec.concl}, with thoughts on future prospects for understanding how cosmological observations will lead us to specific new physics. ",
        "conclusions": "\\label{sec.concl}  To face the challenge of determining the fundamental physics responsible for the acceleration of the universe, we need to bring to bear next generation observations of the expansion history and possibly its dependent growth history.  The precision and accuracy of these future observations will guide us a long way to identifying new physics.  We see that at the heart of the next step lies a single function -- the effective equation of state $w(z)$.  Mapping this describes the cosmology; models with the same function, or equivalently same expansion history, will agree on the cosmological tests, whether distance-redshift, growth of structure, etc.  Furthermore the simple parametrization in terms of the present value, $w_0$, and a measure of the time variation, $w_a$, proves extraordinarily robust regardless of the exact reason for elaborating on the matter density term in the Friedmann equation.  This is not to say there is no complementarity between cosmological probes; indeed that is a crucial ingredient in constraining the {\\it values} of the equation of state parameters.  And next generation experiments will be superb at achieving this.  The simplicity of a two parameter functional form means we cannot easily appeal to ``naturalness'' to decide which physics model -- dark energy or modified gravity, say -- is a most likely explanation.  Despite the models considered here, though, there is no guarantee that an arbitrary modification $\\delta H^2$ can be fit in terms of $w_0$, $w_a$. Regardless, the function $w(z)$ encodes all the standard, ``smooth'' information regardless of origin.  We have illustrated this for several classes of physics including scalar field dark energy, modifications of general relativity in the Friedmann equation, and direct acceleration through Ricci ``geometric dark energy'', both in general and for specific models.  Explicit examples of the fits were given for probes such as magnitude-redshift, growth factor or gravitational potential, and distance to the CMB last scattering surface.  This held even for models with quite large time variation of the effective equation of state.  One possible breakdown of the simple dark energy mimic ability might occur through the curvature of the gravitational potential decay behavior; the slope is remarkably model independent at low redshifts and asymptotically matter dominated at high redshift, but the localized deviation in between might provide a clue to the accelerating physics. Precision observations of the integrated Sachs-Wolfe effect or the lensing induced CMB bispectrum, yet untested, might be useful probes for this.  We considered the implications of acceleration in general, regardless of origin, through the Ricci scalar curvature. This is pleasingly directly related to the expansion and fate of the universe.  In a conformal horizon history diagram (Fig.\\ \\ref{fig.ahinv}) we illustrate conditions for both acceleration and superacceleration, and briefly discuss the role of superacceleration in particle production that could nullify the Big Rip and indeed possibly provide an attractor for the universe to an apparent cosmological constant state.  The picture of an achievable and wide ranging goal in measuring $w(z)$ is attractive. In our quest for understanding fundamental physics, though, we always want to push deeper.  The virtues of simplicity and broad applicability contest with lack of leverage in separating the root causes.  But it is only in the absence of new dynamics, new equations of motion, that the equation of state $w(z)$ or the expansion history $a(t)$ rules all.  New terms -- interactions or graininess -- lead to complexity but a grip on deeper details of the new physics.  This graininess could come from an observable consequence of dark energy perturbations or a noncanonical sound speed, separating it from a ``smooth'' gravity law (though it is only useful if it occurs within a realm accessible to precision observations).  Conversely, couplings in the gravitational sector, going beyond the Ricci spacetime geometry approach analyzed here, could distinguish a gravitational origin from one of dark energy.  This could arise in scalar-tensor theories, or metric perturbation terms $\\dot h$ in the growth equation, or local curvature dependent effects $\\delta R$, e.g. backreaction from structure formation.  This is rather analogous to the situation in early universe acceleration -- inflation theory.  The incredible simplicity and generic power of it in solving cosmological and high energy physics conundra is immensely attractive, and we shouldn't lose sight of it, just as we shouldn't lose sight of the crucial role of $w(z)$.  But acceleration, then and now, is very much more than just a deSitter state.  We {\\it want} complexity in the form of perturbations, tilt, gravitational waves to learn about the details of the fundamental physics.  For the CMB, measuring $\\delta T/T$, or the power spectrum, is a stunning experimental accomplishment, just as $w(z)$ will be, but we want to explore further through nongaussianities, polarization, etc.  So too we look forward to probing gravity, dark energy, and acceleration."
    },
    "0402/astro-ph0402029_arXiv.txt": {
        "abstract": "{ We present results concerning the occurrence of Seyfert galaxies in a new automatically selected sample of nearby Compact Groups of galaxies (UZC-CGs). Seventeen Seyferts are found, constituting $\\sim$3\\% of the UZC-CG galaxy population. CGs hosting and non-hosting a Seyfert member exhibit no significant differences, except that a relevant number of Sy2 is found in unusual CGs, all presenting large velocity dispersion ($\\sigma$$>$400 km\\,s$^{-1}$), many neighbours and a high number of ellipticals. We also find that the fraction of Seyferts in CGs is 3 times as large as that among UZC-single-galaxies, and results from an excess of Sy2s. CG-Seyferts are not more likely than other CG galaxies to present major interaction patterns, nor to display a bar. Our results indirectly support the minor-merging fueling mechanism. \\keywords {galaxies:clusters:general - galaxies:Seyfert - galaxies:interactions} } \\titlerunning{Seyferts in UZC-CGs} ",
        "introduction": "Because of their high number density of galaxies (comparable to the central density in clusters) and relatively low velocity dispersion ($\\approx$ 200-300 km\\,s$^{-1}$), Compact Groups (CGs) are predicted to constitute the most probable sites for strong galaxy-galaxy interactions and mergers to occur. As a consequence, they are also expected to display a high fraction of AGNs, provided they are bound systems \\citep{Hickson92,Diaferio00} and the interaction-activity paradigm \\citep{Bar92,shlosman90} holds true. The detected fraction of AGNs in CGs might then help to constrain the dynamical status of CGs. A high fraction of AGNs would indicate that CGs are not only physical, but also highly unstable, and would thus support the interaction-activity paradigm, as well as hierarchical scenarios in which large isolated elliptical galaxies \\citep{Zablu98,Borne00} are eventually the end-product of every CG. Conversely, a low fractions of AGNs would be more in accordance with recent results indicating that the occurrence of emission-line galaxies decreases in dense environment \\citep{Balogh03,Gomez03} and with optical and IR observations claiming that spiral galaxies in CGs \\citep{Sulentic93,Verdes98,kelm03}, in groups \\citep{Maia03} and in pairs \\citep{Bergvall03} do not show starburst and/or Seyfert enhancement typically expected in interacting galaxies. In general, a strong correlation appears to hold between AGN and the presence of tidal interaction only for very luminous QSOs \\citep{Bahcall97}, while the excess of companions for Seyferts is still controversial \\citep{Dahari85,Keel85,Rafanelli95,Fuentes88,Mackenty89,Keel96,Derobertis98,Schmitt01}.  Concerning CGs, only the Hickson Compact Groups (HCG, Hickson 1982, 1997) have been extensively studied as for many years it was the only large and uniform sample available. The results from this sample are somewhat conflicting. Several HCGs show evidence of ongoing interaction, but components usually remain distinct, with recognizable morphological types \\citep {Sulentic97}. The fraction of blue ellipticals (which are plausible merger remnants) has turned out to be rather low (4 in 55), predominantly associated with faint members (Zepf {\\it et al.} 1991) and similar to the estimated fraction ($\\approx$7\\%) of currently merging galaxies (Zepf 1993). Hickson {\\it et al.} (1989) found the FIR emission in HCGs to be enhanced  compared to a sample of field galaxies, but  Sulentic \\& de Mello (1993) and Verdes-Montenegro {\\it et al.} (1998) suggest there is no firm evidence for enhancement.  The specific issue of AGNs in HCGs has been addressed by the Kelm {\\it et al.} (1998) finding that only $\\approx$2\\% of the member galaxies display a Seyfert spectrum, and that this fraction is similar to that found in galaxy pairs. They also find that HCG Seyferts are hosted by luminous spirals, as is usually the case \\citep{Heckman78}. However, a relevant population of low-luminosity AGNs (LLAGNs) in HCGs as well as in the SCG sample (Iovino 2002) has been revealed by means of deep resolution spectroscopy (Coziol {\\it et al.} 1998a, 1998b, 2000), with a significant preference for early-type hosts. Coziol {\\it et al.} (2000) state that AGNs (including low luminosity/dwarf sources) are the most frequent (41\\%) activity type encountered in CGs. Shimada {\\it et al.} (2000) confirm a high fraction of LLAGNs in HCGs, but claim that there is no statistically significant difference in the frequency of occurrence of emission line galaxies between the HCGs and the field. Whether a CG environment really triggers an AGN at all, or whether CGs are favourable hosts for low-excitation, low luminosity AGNs only, remains controversial, however.  In this paper we address the issue of Seyfert occurrence in CGs making use of the new large sample of UZC-CGs (Focardi \\& Kelm 2002), selected from a 3-D magnitude limited  catalogue (UZC, Falco et al. 1999). In Sect. 2 the UZC-CG and the Seyfert samples are presented, in \\S 3 UZC-CGs with and without a Seyfert are compared; a similar comparison for HCGs is performed in \\S 4. In \\S 5 the frequency of Seyferts in UZC-CGs and in a single-galaxy sample (selected in UZC) are discussed. In \\S 6 and 7 we address the relative occurrence of Sy1, Sy2 and LINERs, and in \\S 8 the presence of interaction patterns.  A Hubble constant of $H_{o}$\\,=\\,100\\,km\\,s$^{-1}$\\,Mpc$^{-1}$ is used throughout. \\section {The samples} UZC-CGs (Focardi\\,\\&\\,Kelm 2002) have been extracted from the UZC catalog (Falco et al. 1999) using an objective algorithm. UZC lists redshift for nearly 20\\,000 galaxies in the northern sky and is 96\\% complete for m${_B}$\\,$\\le$\\,15.5 galaxies. UZC-CGs are systems of 3 or more galaxies lying inside a 200$h^{-1}$\\,kpc radius area and radial velocity within 1000\\,km\\,s$^{-1}$ from the center. Possible ACO clusters substructures  have been excluded from the UZC-CG sample.  The present analysis is restricted to 192 UZC-CGs (639 galaxies) in the 2500-7500 km\\,s$^{-1}$ radial velocity range. The lower limit in radial velocity  has been set to avoid possible Local Supercluster structures and major contamination of distances by the effect of peculiar motions. The upper limit avoids including a large population of galaxies with uncertain morphological classification.  Seyfert galaxies are identified by cross correlating UZC-CG galaxies with the Veron-Cetty\\&Veron (2001, V\\&V) AGN catalogue. V\\&V is still the largest all-sky available catalogue of bright nearby Seyfert galaxies. It is neither complete nor homogeneous, nevertheless it is not severely affected by survey biases (i.e. there are no large sky regions in which the absence of Seyferts can be attributed to the lack of data) and can be used at least for a preliminary analysis.  The NED database has further been inspected to include any additional AGN and to check V\\&V's activity classification. It has also been used to assign morphological classification, available for 75$\\%$ of the galaxies. \\begin{table*} \\begin{center} \\caption{Seyferts in the UZC-CG Sample. } \\begin{tabular}{||r|l|c|c|l|r|r|r|r|l||} \\hline UZC-CG  & name  &V\\&V & NED  & morphology & cz  & m$_B$ &lum. rank$^1$ &inter.$^2$ &other id. \\cr \\hline &               &     &     &        &      &     &   &   & \\cr 19  & NGC449    & Sy2 & Sy2 & S      & 4750 & 15.2 &3&  & MKN 1 \\cr 23  & NGC513    & Sy2 & Sy2 & Sb/c   & 5840 & 13.4 &1&  & ARK 41\\cr 30  & UGC1479       & Sy2 & Sy2 & Sc & 4927 & 14.8 &2&  &       \\cr 57  & UGC3752       & Sy2 &     & S? & 4705 & 14.5 &2&  &       \\cr 109 & ZW63.06       & Sy2 &     &    & 3988 & 15.5 &3&  &       \\cr 143 & NGC3798       & Sy1 &     & SB0    & 3563 & 13.9 &1&  &       \\cr 144 & NGC3822   & Sy2 & Sy2 & Sb     & 6138 & 13.7 &1& Y & HCG58a \\cr 156 & NGC4074       & Sy2 & Sy2 & S0 pec.& 6802 & 15.4 &6&  & AKN347 \\cr 162 & NGC4169       &     & Sy2  & S0/a   & 3755 & 12.9 &1&    & HCG61a\\cr 186 & IC4218        & Sy1 & Sy1 & Sa     & 5808 & 14.4 &1&  & \\cr 200 & ZW45.099      & Sy2 & Sy2 &        & 6870 & 15.4 &2& & \\cr 207 & NGC5395   &     & Sy2 & Sb     & 3522 & 12.6 &4& Y  & Arp84, IZw77 \\cr 240 & NGC5985       & Sy  & Sy1 & Sb   & 2547 & 12.0 &1&  & \\cr 242 & NGC5990   & Sy2 &     & Sa pec.& 3726 & 13.1 &1& &  \\cr 272 & NGC6967       & Sy2 & Sy2 & SB0    & 3768 & 14.3 &3&  & \\cr 276 & UGC11950  & Sy2 &     & E      & 6224 & 14.3 &1&   & \\cr 279 & NGC7319   & Sy2 & Sy2 & SBpec. & 6652 & 14.8 &2& Y   &HCG92c \\cr &               &     &     &        &      &     &   &   & \\cr \\hline \\end{tabular} \\note{col.8 indicates the luminosity rank of the Sy galaxy within its group} \\note{col.9 indicates the presence of major morphological-peculiarities/interaction-patterns visible on POSS plates.} \\par \\end{center} \\end{table*} In Table 1 basic data for each Sy1 and Sy2 in UZC-CGs are listed: UZC-CG number (column 1), Seyfert name (column 2), V\\&V and NED activity classification (column 3 and 4), morphological classification (column 5), radial velocity (column 6), apparent blue magnitude as in UZC (column 7), luminosity rank (1 = first ranked, 2 = second ranked ...) (column 8), presence of interaction pattern (column 9), other identification (column 10).  To identify LINERs (L) we have also inspected the list by Carrillo et al. (1999). LINERs are named Sy3 in V\\&V, and occasionally AGN in NED. LINERs in UZC-CGs are listed in Table 2. LINERs are found in galaxies of earlier Hubble type than Seyferts and their nuclear continua are usually dominated by old stars \\citep {Heckman80,Ho03}. We have kept Sy1 and Sy2 separate from LINERs because the nature of the central power-source in LINERs remains uncertain; they might constitute a transition class between non thermal objects and starburst \\citep{Carrillo99} or between normal galaxies and Seyferts rather than a true low luminosity extension of the AGN sequence. ",
        "conclusions": "We have investigated the occurrence of Seyfert galaxies in a nearby sample of 192 UZC-CGs. We find 17 Seyferts among the 639 member galaxies, indicating that only a minor fraction ($\\sim$ 3\\%) of UZC-CG galaxies host a Seyfert.  When comparing velocity dispersion, number of large scale neighbours and morphological content of Sy-CGs and non-Sy-CGs, no significant differences are found, although some Seyferts (5 Sy2) appear associated with 'extreme' CGs, presenting large $\\sigma$ ($>$400\\,km\\,s$^{-1}$), many neighbours and an unusually high number of Elliptical members.  This suggests that Sy-CGs and non-Sy-CGs are drawn from the same parent population, and that the presence of the Seyfert is not linked to specific CG properties.  We also find the fraction of Sy in CGs (3\\%) to be significantly higher than the fraction of Sy in a single-galaxy sample (1\\%). Curiously, the enhanced fraction of Seyferts in CGs reflects the behaviour of Sy2 alone: while 14 Sy2 and 3 Sy1 (5:1) are found in UZC-CGs (the ratio is 6:1 in HCGs), in the single-galaxy sample there are 33 Sy2 and 22 Sy1 (3:2), suggesting that the location within a group potential and the triggering of type 2 Seyferts are related.  No excess of interacting galaxies or barred galaxies is found among Sy-CGs, indicating that strong galaxy-galaxy interaction is not a common Seyfert fueling mechanism \\citep{Derobertis98,shlosman90} in CGs. But any minor merging (between a galaxy and a satellite/dwarf companion) fueling mechanism \\citep{taniguchi99} is not ruled out. The large number of isolated Seyferts we find in our analysis, indirectly supports the minor-merging mechanism."
    },
    "0402/astro-ph0402517_arXiv.txt": {
        "abstract": "{ In three previous papers (Pelat 1997, 1998 and Moultaka \\& Pelat 2000), we set out an inverse stellar population synthesis method which uses a database of stellar spectra. Unlike other methods, this one provides a full knowledge of all possible solutions as well as a good estimation of their stability; moreover, it provides the unique approximate solution, when the problem is overdetermined, using a rigorous minimization procedure. In Boisson et al. (2000), this method has been applied to 10 active and 2 normal galaxies.\\\\ In this paper we analyse the results of the method after constraining the solutions. Adding {\\it a priori} physical conditions on the solutions constitutes a good way to regularize the synthesis problem. As an illustration we introduce physical constraints on the relative number of stars taking into account our present knowledge of the initial mass function in galaxies. In order to avoid biases on the solutions due to such constraints, we use constraints involving only inequalities between the number of stars, after dividing the H-R diagram into various groups of stellar masses.\\\\ We discuss the results for a well-known globular cluster of the galaxy M31 and discuss some of the galaxies studied in Boisson et al. (2000). We find that, given the spectral resolution and the spectral domain, the method is very stable according to such constraints (i.e. the constrained solutions are almost the same as the unconstrained one). However, an additional information can be derived about the evolutionary stage of the last burst of star formation, but the precise age of this particular burst seems to be questionable. ",
        "introduction": "\\label{sec:intro}   The search of the stellar populations inside unresolved galaxies has been the aim of several studies since the sixtees. Two different approaches have been adopted for this purpose: the direct approach of which methods are usually called ``the evolutive synthesis methods'' (e.g. Tinsley 1972, Charlot \\& Bruzual 1991, Bruzual \\& Charlot 1993, Leitherer et al. 1999, Fioc \\& Rocca-Volmerange 1997, Vazdekis, 1999, Bruzual \\& Charlot 2003) and the inverse one of which these are called ``the synthesis methods'' (e.g. Faber 1972, O'Connell 1976, Joly 1974, Bica 1988, Schmidt et al. 1989, Silva 1991, Pelat 1997,1998, Moultaka \\& Pelat 2000). \\\\ In the first approach, one decides an {\\it a priori} model for the history of the star formation occuring inside the studied galaxy, and by means of theoretical stellar evolutionary tracks and of a stellar database, derives quantities that are directly compared to the observed ones. From different input models, one retains the model that best fits the observed quantities.\\\\  In the inverse approach, usually, no {\\it a priori} model is necessary to derive the stellar populations and one uses exclusively the observables in order to deduce the stellar spectral types and luminosity classes by means of a minimization procedure. This minimization task is not a very simple one since the absolute minimum is usually difficult to find because the ``objective function'' (which is the function that one has to minimize) can rarely be minimized analytically.\\\\  Whatever the approach is, the problem of stellar population synthesis often suffers from the lack of true solutions or, on the contrary, from their degeneracy (i.e. multiple solutions) and/or finally from their instability (i.e. small errors around the observations can induce discontinuities in the solutions). These three inconveniences have been controlled in the inverse method described in Pelat (1997,1998, hereafter Paper I and II) where all the solutions are well identified and the minimization procedure is rigorously treated, as well as in Moultaka \\& Pelat (2000) where the stability of the various solutions is analysed (hereafter Paper III). \\\\  In this paper, we study the influence upon the different solutions of astrophysical constraints included {\\it a priori} when searching for a solution. The concept of constraining the solutions in the inverse methods, has been adopted by O'connell (1976), Pickles (1985) and Silva (1991) but no estimation of the induced bias has been given by these authors to our knowledge.\\\\  In the next section, we recall briefly the inverse method and its error analysis; in the third section, we describe constrained models. Finally, in section \\ref{sec:results}, we show and discuss constrained versus unconstrained results for the globular cluster G170 located in M31, the LINER NGC4278, the starburst NGC3310 and the Seyfert 2 galaxy NGC2110. In the last section we make a general description of the behaviour of constrained solutions in the 27 central regions of the twelve galaxies studied in Boisson et al. (2000), hereafter Paper IV. ",
        "conclusions": "\\label{sec:concl}  The ideal case for a spectral synthesis giving a synthetic distance equal to zero would be the case where the signal to noise ratio of the galactic and stellar spectra goes to infinity and where the stellar database is complete. In such a case, all stars with spectral types later than the spectral type at the turnoff position would have non zero contributions to luminosity. But in practice all unconstrained solutions show many zero contributions; this is due to the finite signal to noise ratio of our spectra and to the limitation of the stellar database which itself is due to the finite spectral resolution.  Therefore, constraining the stellar population would {\\it {a priori}} reduce the number of zero contributions because of the additive information introduced in this process. But as can be seen in the previous results, no large improvement in eliminating the zero contributions has appeared. This is probably due to the not perfect adequation of the observational data. Actually  as the constraints are expressed by large inequalities (i.e. equalities are allowed) optima are usually located on the border of the domain in which solutions are constrained.\\\\  The stellar synthesis method with constraints presented in this paper has been applied to the 27 regions of galaxies studied in Paper IV. In general, all 27 regions present ``Standard mode'' solutions equal or very similar to the unconstrained solution. Moreover, all zero contributions in the unconstrained solutions remain null in the ``Standard mode'' or have small ill-defined values and all well- and ill- determined contributions remain respectively well- and ill- defined. This result shows that ``Standard mode'' solutions are generally included inside the error bars of the unconstrained solution and when they are not, their synthetic distances are at several $\\sigma$ from that of the unconstrained solution. \\\\  In the ``Decreasing IMF mode', the number of star classes contributing to the synthesis is often larger than that of the unconstrained solution and of the ``Standard mode''. This fact affects especially dwarf stars and is due to the sharper distribution of stars in the mass groups of the H-R diagram in this mode. For the same reason and because the number of constraints is larger, the synthetic distances here are in general larger than those of the previous mode and of the unconstrained solution.\\\\   In both modes (``Standard'' and ``Decreasing IMF'') some solutions satisfy their constraints on the border of the domain. This shows that in such cases constraints are somehow too strong and induce bias. However, these solutions provide some indications, thanks to the error bars, allowing one to find acceptable solutions that satisfy the desired conditions inside the domain of constraints (see previous examples). \\\\  As a matter of fact, the solution of the least square problem is the one that minimizes the synthetic distance; this happens often on the border of the domain of constraints but the goal is not to find the optimal mathematical solution, rather a ``realistic'' or physical one next to the minimum. \\\\  All previous results are very well confirmed in the case of the globular cluster G170. This is a very important point since this object is constituted of a single burst of star formation; consequently, any deviation of the behaviour of the resulting stellar population due to the inclusion of astrophysical constraints can clearly be detected in this object. \\\\  This study has shown that the inverse method described in this paper and in Papers I, II and III is very stable against the inclusion of additional astrophysical constraints, and is, therefore, very reliable.\\\\  However, constraining the solutions and using the information provided by the error analysis allows one to find similar solutions with younger bursts of star formation. Thus, it is crucial to perform tests such as in the previous section and to discuss the results, especially the different possible locations of the turnoffs, i.e. the age of the last burst of star formation.\\\\  \\appendix"
    },
    "0402/astro-ph0402667_arXiv.txt": {
        "abstract": " ",
        "introduction": "In the preceding talk, John Hawley presented an overview of studies of MHD turbulence-driven accretion onto black holes.  After outlining the principal physical mechanisms, he summarized how to simulate numerically such a system. At the close of his talk, he described the main structures that are seen.  In this talk, I will report in greater detail some of the properties of the accretion flows observed in these simulations and indicate some of their implications for observations.  Ultimately, it is the magnetic field that controls how matter accretes, so the first topic in this review will be the distribution of field intensity, the topology of field-lines, and the way both depend on black hole spin (\\S II).  Once this basis has been established, it will be possible to see how the magnetic field controls the accretion rate (\\S III). Accreting matter releases energy via dissipation of fluid motions and magnetic field; the heated gas can then generate photons.  Although these simulations do not treat the thermodynamics of accretion explicitly, they do offer strong hints about how radiation may occur; some of these are discussed in \\S IV.  We conclude this review with a list of some of the more interesting implications of our results (\\S V). ",
        "conclusions": ""
    },
    "0402/astro-ph0402451_arXiv.txt": {
        "abstract": "We present new $JHK$ photometry on the MKO-NIR system and $JHK$ spectroscopy for a large sample of L and T dwarfs. Photometry has been obtained for 71 dwarfs and spectroscopy for 56. The sample comprises newly identified very red objects from the Sloan Digital Sky Survey (SDSS) and known dwarfs from the SDSS and the Two Micron All Sky Survey (2MASS). Spectral classification has been carried out using four previously defined indices (from Geballe et al. 2002, G02) that measure the strengths of the near infrared water and methane bands.  We identify 9 new L8--9.5 dwarfs and 14 new T dwarfs from SDSS, including the latest yet found by SDSS, the T7 dwarf SDSS J175805.46$+$463311.9. We classify 2MASS J04151954$-$0935066 as T9, the latest and coolest dwarf found to date.  We combine the new results with our previously published data to produce a sample of 59 L dwarfs and 42 T dwarfs with imaging data on a single photometric system and with uniform spectroscopic classification. We compare the near-infrared colors and absolute magnitudes of brown dwarfs near the L--T transition with predictions made by models of the distribution and evolution of photospheric condensates.  There is some scatter in the G02 spectral indices for L dwarfs, suggesting that these indices are probing different levels of the atmosphere and are affected by the location of the condensate cloud layer.  The near-infrared colors of the L dwarfs also show scatter within a given spectral type, which is likely due to variations in the altitudes, spatial distributions and thicknesses of the clouds. We have identified a small group of late L dwarfs that are relatively blue for their spectral type and that have enhanced FeH, H$_2$O and K~I absorption, possibly due to an unusually small amount of condensates.  The scatter seen in the $H-K$ color for late T dwarfs can be reproduced by models with a range in surface gravity.  The variation is probably due to the effect on the $K$-band flux of pressure-induced $\\rm H_2$ opacity. The correlation of $H-K$ color with gravity is supported by the observed strengths of the $J$-band K~I doublet. Gravity is closely related to mass for field T dwarfs with ages $>10^8$~yrs and the gravities implied by the $H-K$ colors indicate that the T dwarfs in our sample have masses in the range 15 -- 75 $\\rm M_{Jupiter}$. One of the SDSS dwarfs, SDSS J111010.01$+$011613.1, is possibly a very low mass object, with log $g$ $\\sim$ 4.2 -- 4.5 and mass $\\sim$ 10 -- 15 $\\rm M_{Jupiter}$. ",
        "introduction": "Since the discovery of dwarfs of spectral type later than M as companions to nearby stars (Becklin \\& Zuckerman 1988; Nakajima et al. 1995) major observational and theoretical progress has been made, thanks to sensitive new wide-area surveys at optical (0.4--1.0~$\\mu$m) and infrared (1.0--2.5~$\\mu$m) wavelengths and models of atmospheres of dwarfs with effective temperatures between those of the coolest stars and the giant planets (see the reviews by Chabrier \\& Baraffe 2000; Burrows et al. 2001). Two new spectral classes have been identified later than type M: the L dwarfs, characterized by the disappearance of gas-phase TiO and VO, and the T dwarfs, characterized by methane band absorption in the $H$ and $K$ spectral regions (Mart\\'{\\i}n et al. 1997, 1999b; Kirkpatrick et al. 1999, 2000; Strauss et al. 1999; Leggett et al. 2000, 2002b; Geballe et al. 2002, hereafter G02; Burgasser et al. 2002a; Hawley et al. 2002).  Most L dwarfs and all T dwarfs are brown dwarfs.  These objects are of interest because they occupy the mass range between that of stars and giant planets; because many of them are likely to have the intrinsic properties of giant planets, which at present cannot be directly observed; and because they allow the investigation of the initial mass function to substellar masses.  Field L and T dwarfs have been discovered in large numbers in recent sky surveys: the Deep Near Infrared Survey (DENIS, Epchtein 1997); the Two Micron All Sky Survey (2MASS, Skrutskie et al. 1997; Beichman et al. 1998), and the optical Sloan Digital Sky Survey SDSS (York et al. 2000). Including objects described in the present paper, there are now about 280 L dwarfs and 58 T dwarf systems published (e.g. Delfosse et al. 1997, 1999;  Kirkpatrick et al. 1999, 2000;  Burgasser et al. 2002a, 2003e; G02). This large sample has been used to establish a complete spectral sequence from L0 to T9 (G02; Burgasser et al. 2002a; McLean et al. 2003;  present paper).  Unlike stars, brown dwarfs lack a sustained source of thermonuclear energy, and hence cool continuously, passing through the L and T stages, with their initial spectral types depending on their masses. Their observational properties are thus a function not only of mass and metallicity, but also of age. Not all dwarfs of a given spectral type or effective temperature are identical; they have different gravities and different colors, the latter likely due to differing amounts of particulate matter in the atmosphere.  For example, in mid to late L dwarfs there is a large scatter in the $JHK$ colors and apparently no one-to-one correspondence between effective temperature and spectral type (Leggett et al. 2002a; Golimowski et al. 2004).  Such considerations drive searches for and measurements of additional dwarfs, in order to more fully characterize their atmospheres.  2MASS and SDSS have been highly complementary  in the discovery of L and T dwarfs.  Most of the flux from late-type dwarfs is emitted longward of 1 $\\mu$m, and the $J-H$ and $H-K$ colors of M and L dwarfs become redder with decreasing effective temperature, allowing the identification by 2MASS of large numbers of L dwarfs (Kirkpatrick et al. 1999, 2000).  However, in the transition from L to T, $\\rm CH_4$ absorption appears in the $H$ and $K$ regions (and $\\rm H_2$ absorption predominantly at $K$), strengthening with later spectral type and causing the T dwarfs to become increasingly blue in their $JHK$ colors.  The $JHK$ colors of early T dwarfs are similar to those of the common K and M stars, making their identification in 2MASS very difficult. However in the SDSS filters the dwarfs simply become redder and thus the early T dwarfs have been found primarily in SDSS imaging (Leggett et al. 2000).  Only three objects with spectral types between T0 and T3.5 have been identified from sources other than the SDSS -- Mart\\'{\\i}n et al. (2001) in an optical and near-infrared imaging survey of the $\\sigma$ Orionis cluster tentatively classified one member as T0, Liu et al. (2002) found a distant field T3--T4 dwarf in a deep $Iz$ survey, and McCaughrean et al. (2003) found that $\\epsilon$ Indi B, discovered in a high proper motion optical survey, is a binary consisting of a T1 and T6 pair (see also Scholz et al. 2003, Smith et al. 2003 and Volk et al. 2003).  To date, all T dwarfs later than T7 have been found in the 2MASS database (Burgasser et al. 2002a) but we expect such objects to be found in the SDSS imaging data as sky coverage is increased.  In this paper we present near-infrared photometry and spectra of new and previously reported L and T dwarfs (including 14 new T dwarfs) and compare the colors and spectra with predictions from state of the art model atmospheres with and without clouds.  The new objects observed are described in the next section, and the new $JHK$ photometric and spectroscopic observations are described in \\S 3, where we also derive spectral types.  \\S 4 presents colors, spectral types and absolute magnitudes for the entire body of near-infrared data on L and T dwarfs which we have accumulated to date.  In \\S 5, we compare these data to model atmospheres. The conclusions are given in \\S 6. ",
        "conclusions": "\\subsection{Characteristics of L and T Dwarf Atmospheres}  As effective temperatures of dwarfs cool to those of the late M dwarfs and below, two chemical changes occur in their photospheres that strongly impact their emergent spectral energy distributions.  The first to occur, for late M dwarfs, is the appearance of corundum ($\\rm Al_2O_3)$ grains within the photosphere (Jones \\& Tsuji 1997) and the formation of condensate clouds.  At the even lower effective temperatures of the L dwarfs iron and silicate are the most important condensates\\footnote {These condensates are frequently termed ``dust'', but this can be misleading since in many dwarfs the iron will be in the liquid phase (Lodders 1999). Hence we generally prefer the terms ``condensate'' to refer to grains or drops of the condensed phase, and ``cloud'' as the region in the atmosphere within which the condensed species are found.}.  The effect of the clouds is to weaken or veil the molecular absorption bands and to redden the $JHK$ colors of L dwarfs (see e.g. Ackerman \\& Marley 2001, Allard et al. 2001, Marley et al. 2002, Tsuji \\& Nakajima 2003).  The extent of these effects depends upon the number, size, and vertical distribution of the condensates; for spectral modeling these parameters must either be computed from a model or somehow specified.  Ackerman \\& Marley (2001) developed a one-dimensional model of mixing and sedimentation for this purpose.  In their model upward vertical mixing of gas and condensate replaces condensates that fall through the cloud base, while far above the cloud base sedimentation efficiently cleanses the atmosphere of condensates. Tsuji \\& Nakajima (2003)  model the cloud by specifying the temperature range within which the condensates are found.  They describe their limits to be the points at which the atmosphere is cool enough for condensation but hot enough that the condensates are small enough to remain suspended in the atmosphere and are less prone to sedimentation. The temperature domain in which the cloud is formed depends on the details of the model used, but is around $\\sim$1500--1700~K (Ackerman \\& Marley 2001) or $\\sim$1800--2000~K (Tsuji \\& Nakajima 2003). In the T dwarfs the cloud layer lies near the base or below the wavelength-dependent photosphere and plays a smaller role in determining the observed flux distribution.  The second and later chemical change that occurs in these high-pressure, low-temperature atmospheres is the formation of additional molecular species in the photosphere, most importantly CH$_4$.  CO and H$_2$O are abundant in M dwarf atmospheres, but by mid-L the abundance of CH$_4$ becomes significant at the expense of CO (Noll et al. 2000). At moderate spectral resolution CH$_4$ absorption is not seen in the near-infrared until temperatures drop to those of the late L dwarfs, at which point $K$-band CH$_4$ features are detectable, and at T0 (by definition, G02) CH$_4$ absorptions are seen at both $H$ and $K$.  The increasing CH$_4$ absorption largely accounts for the increasingly blue $JHK$ colors of the T dwarfs with later spectral type, more than compensating for the reddening due to the decreasing effective temperature.  For dwarfs of type T5 and later, pressure-induced H$_2$ becomes a significant opacity source.  This opacity depresses the flux in the $K$ band, and to a lesser extent the $H$ band, and also contributes to the blue near-infrared colors.  For a useful summary of the important molecular species and the wavelength ranges in which they are observed see Figure 15 of Burrows et al. (2001).  \\subsection{Clouds, Molecules and Classification Schemes}  Spectral classification schemes for L and T dwarfs have been developed using both the red and the near-infrared spectral regions.  In the late 1990s Kirkpatrick et al. (1999) and Mart\\'{\\i}n et al. (1999b) developed schemes using the strengths of various absorption features and pseudo-continuum slopes seen in optical spectra to classify the L dwarfs. A few years later, Burgasser et al. (2002a) and G02 presented schemes using the strengths of the near-infrared molecular absorption bands  to classify the T dwarfs, and in the case of G02, L dwarfs also. While the Burgasser et al. (2002a) and G02 schemes for T dwarfs give results in very close agreement, the G02 scheme for L dwarfs can give results that differ by as much as 2.5 subclasses from those given by the Kirkpatrick et al. (1999) scheme, suggesting that there are significant differences in the optical and infrared classification of L dwarfs. Dwarfs in our sample whose optical and infrared spectral classes differ by more than one subclass are identified in Table 9.  Some of the scatter is due to the small differences in the infrared indices from one subtype to the next, combined with measurement errors.  However, the differences in optical and infrared spectral types are not entirely random. Where there are differences, G02 generally assign earlier spectral types to those objects that are redder than average in $J-K$, and later types to those which are bluer (Stephens 2001, Figure 7.5).  Models of L and T dwarf atmospheres give some insight into the variations seen among the spectral indices.  Generally speaking, the spectra of L and T dwarfs are less sensitive to the effects of cloud decks in the 0.7--1.0~$\\mu$m spectral region than at wavelengths longer than 1$\\mu$m. This is because, for effective temperatures corresponding to the earliest L types, optical depth unity in the far-red is reached below the cloud deck, but as the cloud is still fairly optically thin it does not substantially influence the red spectrum.  At lower effective temperatures, the clouds place a ``floor'' on the region from which the emergent flux arises, but because of the large opacity in the far-red due to initially refractory diatomics and water and later to K~I and Na~I resonance line absorption, most of the outgoing red flux arises from above the cloud decks.  Thus for dwarfs with effective temperature cooler than about 1800~K (types $\\sim$L3 and later, Leggett et al. 2002a,b; Golimowski et al. 2004) slight changes in the cloud deck optical depth have little effect on the emergent red spectrum.  In the near-infrared (particularly the $Z$ and the $J$ bands), the windows between the molecular bands of water and other opacity sources allow flux to emerge from very deep in the atmosphere.  In these regions the opacity floor imposed by the clouds substantially alters the depth to which one can see into the atmosphere (see Figure 7 of Ackerman \\& Marley 2001 and Figure 4 of Marley et al. 2002). Thus for dwarfs with $\\rm T_{eff}$ in the range from about 1800 to 1500~K (roughly L3 to L7), slight changes in the cloud profile substantially alter the near-infrared spectrum.  This proposed atmospheric structure implies that spectral typing schemes for mid to late L dwarfs that rely on far-red spectra (e.g. Kirkpatrick et al. 1999) tend to be less sensitive to the vertical distribution of condensates in the atmosphere than schemes that rely upon near-infrared spectra or spectral indices (G02). Stephens (2001, 2003) considered the boundaries of the regions employed in the G02 spectral typing system and found that the flux in the G02 bandpasses usually originates from within the cloud decks.  For the earliest L dwarfs, cloud opacity is not significant, but for the mid L dwarfs, the G02 1.5 $\\mu$m water index can be a more sensitive indicator of {\\it cloud optical depth} than of {\\it effective temperature}. At the same time, the 2.2 $\\mu$m methane index is more sensitive to $\\rm T_{eff}$ than to cloud properties, since this spectral region is more opaque and optical depth unity is reached higher in the atmosphere. However, the classification system of G02 relies heavily on the 1.5 $\\mu$m index, as it is the only infrared index that covers the entire L spectral type range. This index undergoes a much larger change through the L sequence than does the other useful infrared index, CH$_4$--K.  Note that G02 do not claim that defining the spectral type is equivalent to measuring the effective temperature; they use a simple classification scheme based on spectral appearance in which the effects of gravity and clouds are not separated from those of temperature.  This larger sample of dwarfs also suggests some inconsistency in the infrared classification of late L dwarfs.  While the G02 scheme provides excellent internal consistency for T dwarfs, the H$_2$O 1.5 $\\mu$m index tends to give a later spectral type than does the CH$_4$ 2.2 $\\mu$m index for dwarfs in the range L5--L9.5.  This tendency was not apparent in the smaller G02 sample; the results for the present larger sample suggest that an adjustment of the flux ratio definitions as a function of spectral type for H$_2$O 1.5 $\\mu$m and CH$_4$ 2.2 $\\mu$m in the L5--L9.5 range could give better internal consistency.   \\subsection{Observed Color as a Function of Spectral Type}  Figure 3 plots, for the final sample presented in Table 9, various colors against spectral type (determined by the G02 scheme apart from four L dwarfs classified optically, see Table 9);  typical error bars are shown. Only those dwarfs with types determined from their spectra are shown. Clouds strongly affect both colors and spectral types of mid L dwarfs classified from the near-infrared indices, as discussed in \\S 5.1 and \\S 5.2.  This effect can be seen in Figure 3, where the spread in the $Z$ through $K$ colors is greatest from L3 to L7.5, just when the clouds in the detectable atmosphere are expected to be most optically thick.  The overall conclusion is that color cannot be used as an (infrared) spectral type indicator for L dwarfs. For T dwarfs, $z-J$ and $J-H$ appear to be reasonable indicators of type.  Given improved sensitivity, $i-z$ may also be a useful T-type indicator.  Note, however, that $i-z$ is expected to turn blueward at $\\rm T_{eff}~\\lesssim$~600~K as the Na, K, and other alkalis condense into solids and the opacity of the Na and K lines falls (Burrows et al. 2002; Marley et al. 2002). The late T dwarfs show significant scatter in their $H-K$ and $J-K$ colors. This can be understood in terms of the onset of pressure-induced H$_2$ absorption, which is very gravity sensitive (Borysow, Jorgensen \\& Zheng, 1997). As described in \\S 5.6, we can interpret the observed spread in $H-K$ as a range in surface gravity for the field T dwarf population.  The $Z-J$ colors of the mid T dwarfs also show considerable scatter (Figure 3).  This is an intriguing result, because the $Z$ and $J$ bands are the most transparent windows into these atmospheres, with the $Z$ band being the clearer of the two.  As such, these bands are particularly sensitive to any deep variations in opacity between atmospheres, such as might be related to the upper reaches of any remaining deep silicate cloud (see Figure 7 of Ackerman \\& Marley 2001).  The variation may thus be due to differences in the process(es) responsible for the removal of condensates at the L to T transition. Since the $Z$ band is sensitive to the far wings of the optical Na-D and K I resonance lines (Burrows \\& Volobuyev 2003) these variations might arise from differences in the removal of gaseous Na and K, due to gravity or metallicity effects. More detailed modeling of the L to T dwarf transition and the removal of atmospheric condensates is required to account for these observations.  \\subsection{Families of Models for Comparison with the Data}   In the remainder of this paper we compare observed colors with two varieties of models from Marley et al. (2002).  In the first type of model, condensate opacity is ignored, although condensation chemistry is accounted for in chemical equilibrium and molecular opacities. We term these the ``cloud-free'' models.  In the second type the effects of condensate opacity are computed using the Ackerman \\& Marley (2001) cloud model. In these ``cloudy'' models the efficiency of condensate sedimentation is parameterized by $f_{\\rm sed}$\\footnote{ Ackerman \\& Marley (2001) employed the parameter ``$f_{\\rm rain}$\" to describe the efficiency of condensate sedimentation in a brown dwarf atmosphere.  Strictly speaking, ``rain'' is falling water, and this term has now been replaced by $f_{\\rm sed}$.}.  When the sedimentation efficiency is high (large $f_{\\rm sed}$) both the optical depth and vertical extent of the cloud are small.  In the extreme case of no condensate sedimentation $f_{\\rm sed}=0$.  \\subsection{$J-H$ and $H-K$ Colors of L Dwarfs: Unusually Red and Unusually Blue L Dwarfs}  Figure 4 shows $J-H$ plotted against $H-K$ for the L dwarfs in the sample, where spectral subclass ranges are indicated by different symbols. Overlaid are cloudy model sequences with $\\log g=5$ and $f_{\\rm sed}$ values of 3 and 5.  The $f_{\\rm sed}=3$ models match the data well, although a shift in modeled $J-H$ color of about 0.15 mag would encompass many more of the data points. The discrepancy is likely attributable to the modeled TiO bands at $J$ being too deep, causing the $J$-band magnitudes to be too faint and the $J-H$ model color to be too red. Whether this is a shortcoming in the chemical equilibrium calculation or the molecular opacities themselves is as yet unclear.   Log~$g=$4 models make the $H-K$ colors bluer, in better agreement with the data, but this is an unlikely gravity for these field dwarfs. (Burrows et al. 1997 show that if $\\rm T_{eff}$=1500--2200~K and age=1--5 Gyr then log $g\\geq$5.0.)  The detailed distribution of most of the data points in this color-color space is challenging to interpret as both the models and the observations show that $J-H$ and $H-K$ first become redder and then bluer with falling $T_{\\rm eff}$.  Thus the colors double back on themselves.  Despite the scattered distribution, some extreme objects stand out. The L7.5 dwarf 2MASS J2244$+$2043 and the L5.5 dwarf SDSS J0107$+$0041 are quite red in both colors.  This may imply that their condensate cloud decks are more optically thick than average, which could arise from either less efficient sedimentation ($f_{\\rm sed}\\sim 2$) or higher metallicity.  Our sample also includes four late-type L dwarfs --- SDSS J0805$+$4812, SDSS J0931$+$0327, SDSS J1104$+$5548 and SDSS J1331$-$0116 --- that are unusually blue for their spectral types.  As shown in Figure 3, they are bluer than average at $J-H$ by about 0.2~mag and at $H-K$ by about 0.1~mag.  Applying the shift to the models described above of about $-0.15$ in $J-H$, Figure 4 suggests that these dwarfs are better described by the $f_{\\rm sed}=5$ models, i.e. the sedimentation efficiency is high and the cloud optical depth is small. The spectral indices for all four of these objects show a large range --- the $H$-band indices imply a late type of  L9 to T1 while the $K$-band index gives an earlier type of L5.5--7.5 (the $J$-band index only implies a type earlier than T0). Our spectra show that they have enhanced FeH, K~I and $\\rm H_2O$ absorption (although the spectrum of SDSS J1104$+$5548 is noisy).  Figure 5 shows the $J$-band spectra for SDSS J0805$+$4812, SDSS J0931$+$0327 and SDSS J1331$-$0116 bracketed by more typical L5.5 and L9 dwarfs.  Gorlova et al. (2003) and McLean et al. (2003) show that the equivalent widths of the $J$-band FeH and K~I features peak at spectral types around L3 and then become smaller as the Fe condenses into grains and K is lost to KCl. The strengths of these features in SDSS J0805$+$4812, SDSS J0931$+$0327 and SDSS J1331$-$0116 are similar to those of the early L types, while their H$_2$O bands are more typical of the latest L dwarfs, supporting the interpretation that we are looking deep into unusually condensate-free atmospheres.  The possibility of low metallicity should also be considered. Burgasser et al.  (2003a) identified a late L dwarf (2MASS J05325346$+$8246465) whose extremely blue near-infrared colors are similar to those of the mid T types.  This high velocity dwarf appears to be an extremely metal-poor halo subdwarf with strong FeH features as well as H$_2$ absorption which depresses the  $H$ and $K$ band fluxes.  Cruz et al. (2003) identify two early L dwarfs (2MASS J1300425$+$191235 and 2MASS J172139$+$334415) that are bluer than average at $J-K$ by about 0.2~mag.  As condensate clouds are optically thin for early L dwarfs, and these dwarfs have significant proper motion, they may also be part of a low-metallicity population.   \\subsection{$J-H$ and $H-K$ Colors of T Dwarfs: Gravity and Mass Determinations}  The $H-K$ and $J-K$ colors of T5--T9 dwarfs are scattered (Figure 3), even though the $H$- and $K$-band indices of G02 yield consistent classifications. Figure 6 shows $J-H$ against $H-K$ for the T dwarfs in the sample with sequences from the cloud-free and cloudy $f_{\\rm sed}=5$ models by Marley et al. (2002) overlaid.  The synthetic $H-K$ colors for the late T dwarfs, and the range in color over a plausible range of gravities of log~$g=4.5$--5.5, reproduce the observed colors extremely well.  Gl 570 D provides a further test for the model/data correspondence shown in Figure 6.  Geballe et al. (2001) fit models to the absolute luminosity of Gl 570 D and used age constraints to find $\\rm T_{eff}=784$--824 K and a surface gravity in the range 5.00--5.27. This gravity range is consistent with that implied by Figure 6.  The effective temperature implied by Figure 6, however, is high by about $150\\,\\rm K$.  Although Geballe et al. (2001) found that their best fitting models generally reproduced the $JHK$ spectrum of Gl 570D quite well, there were notable discrepancies.  In particular the notorious inadequacy of the $H$-band methane opacity database and the tendency of all clear-atmosphere models to overestimate the water-band depths limit the fidelity of the fit. These deficiencies both result in the best-fitting temperature contours in Figure 6 being somewhat too warm.  Further, the trends shown in Figure 6 may break down for lower temperatures. One interesting challenge is the T9 dwarf 2MASS J0415$-$0935, the latest and coolest T dwarf currently known, with $\\rm T_{eff}\\approx$700~K (Golimowski et al. 2004, Vrba et al. 2004). Figure 6, 8 and 9 show  that instead of being bluer in $J-H$, $H-K$ and $J-K$ than Gl 570D, it is {\\it redder}, by 0.2 magnitudes, in $J-K$. While models by Marley et al. (2002) and Burrows et al. (2003) predict that, indeed, the coolest dwarfs will become redder in $J-K$ with falling $\\rm T_{eff}$ and the onset of water cloud formation, this happens only for models with $\\rm T_{eff}~\\lesssim$~500~K  unless the brown dwarf is older than 7 Gyr and more massive than 40 $\\rm M_{Jupiter}$. The condensation of the alkalis into their solid chloride forms may also lead to redder colors at these kind of temperatures (Lodders 1999, Marley 2000, Burrows et al. 2003).  Despite these discrepancies, the overall trends seen in Figure 6 can be understood in the context of our presently limited understanding of brown dwarf atmospheres.  At the effective temperatures of late T dwarfs, the $K$-band flux is very sensitive to gravity.  This is because the opacity of pressure--induced $\\rm H_2$ absorption is proportional to the square of the local gas number density.  Higher gravity objects of a given $\\rm T_{eff}$ tend to be cooler at a given pressure than lower gravity dwarfs, and thus have denser, more opaque, atmospheres.  Hence high gravity objects tend to be dimmer at $K$ and bluer in $H-K$ than comparable lower gravity objects. 2MASS J0937$+$2931 shows a particularly depressed $K$-band flux, standing out in Figure 6, presumably due to strong $\\rm H_2$ opacity (Burgasser et al. 2002a). Structural models imply an upper limit to $\\log g$ of 5.5 for brown dwarfs (e.g. Burrows 1997) and therefore Figure 6 suggests that 2MASS J0937$+$2931 may be both a high gravity {\\it and} a low metallicity dwarf, as also suggested by Burgasser et al. (2003b).  While $\\rm H_2$ opacity is also enhanced by decreasing metallicity (e.g. Saumon et al. 1994, Borysow et al. 1997), variations in metallicity are less likely than variations in gravity for this sample of local field brown dwarfs.  $H-K$ appears to be an easy to obtain and straightforward indicator of gravity for late T dwarfs. This is significant as, for a field sample with a likely range in age of 1--5~Gyr (Dahn et al. 2002 estimate 2--4~Gyr based on kinematic arguments), gravity corresponds directly to mass. This tight relationship is due to the small dependence of radius on age or mass for brown dwarfs older than about 200~Myr (Burrows et al. 2001). Figure 9 of Burrows et al. (1997) shows that log~$g=4.5$ implies a mass of 15~$\\rm M_{Jupiter}$, log~$g=5.0$ a mass of 35~$\\rm M_{Jupiter}$, and log~$g=5.5$ a mass of 75 $\\rm M_{Jupiter}$. We list the $H-K$ implied surface gravities for the later T dwarfs in Table 11.  Recent investigations of spectroscopic gravity indicators for L and T dwarfs include those by Lucas et al. (2001), Burgasser et al. (2003b), Gorlova et al. (2003), Mart\\'{\\i}n \\& Zapatero Osorio (2003) and McGovern et al. (2004). Figure 2 of Mart\\'{\\i}n \\& Zapatero Osorio shows synthetic spectra for $\\rm T_{eff}=$1000~K from COND models by Allard et al. (2001).  The models imply that the lines of  K~I at 1.243 and 1.254 $\\mu$m become weaker with increasing gravity for T dwarfs\\footnote{Theoretically this is explained by the fact that the column abundance of molecules above a given pressure level $P$ in an atmosphere is proportional to $P/g$, where $g$ is the gravity.  In a higher gravity atmosphere an outside observer must, all else being equal, look to higher pressure to observe the same column of absorber as in a lower gravity atmosphere.  The calculations of Lodders (1999) show that with rising pressure at a fixed temperature chemical equilibrium increasingly favors KCl over K.  Thus in a higher gravity atmosphere the total column of potassium above the floor set by the continuum opacity is less than in a lower gravity model.  The sensitivity to gravity should be greater in later T (lower $\\rm T_{eff}$) atmospheres since the line forming region (roughly 1000 and 1500~K in the 1.15 and $1.25\\,\\rm \\mu m$ regions respectively)  falls closer to the chemical equilibrium boundary than for earlier T dwarfs (see Figure 2 of Lodders 1999). In addition higher pressures produce greater line broadening, thus decreasing the line depth. }. Figure 7 shows $J$-band spectra of three T6($\\pm0.5$) dwarfs (SDSS J1110$+$0116, 2MASS J2339$+$13 and 2MASS J0937$+$2931) and two T8 dwarfs (2MASS J0727$+$1710 and Gl 570D). These dwarfs span a range in $H-K$ color and are identified in Figure 6. It can be seen that as $H-K$ increases from top to bottom in Figure 7, the K~I lines strengthen, supporting the interpretation of increasing $H-K$ as being due to decreasing gravity. The increase in K~I equivalent widths for these dwarfs is confirmed by the higher resolution NIRSPEC data of McLean et al. (2003).  Comparison of our Figure 7 with Figure 2 of Mart\\'{\\i}n \\& Zapatero Osorio (2003) shows that the depths of the K~I lines seen in 2MASS J0937$+$2931 are similar to the model predictions for log~$g \\sim$5.5, and that the lines seen in SDSS J1110$+$0116 are almost as strong as the synthetic spectrum with log~$g \\sim$3.5.  Figure 6 and the models of Marley et al. (2002) imply that SDSS J1110$+$0116 has log~$g$ between 4.0 and 4.5. There is as yet no direct measurement of the effective temperature of this dwarf, but the effective temperatures of other T6 dwarfs are in the range 900--1075~K (Golimowski et al. 2004).  If SDSS J1110$+$0116 has $\\rm T_{eff}\\approx 1000\\,\\rm K$ and $\\log g \\approx 4.2$, the evolutionary models of Burrows et al. (2003, their Figure 1) imply that it is a 10~$\\rm M_{Jupiter}$ brown dwarf with an age of about 1$\\times$10$^8$ years, i.e. similar to that of the Pleiades cluster.  However we argue in \\S  5.7.4 that one might shift the model contours on Figure 6 up and to the right to bring them into better agreement with observed $J-K$ and measured effective temperatures.  In that case SDSS J1110$+$0116 would have a somewhat larger gravity, mass, and age.  We thus adopt a more conservative estimated mass of 10 -- 15$\\,\\rm M_{Jupiter}$ and an age of 1 -- 3 $\\times10^8$ years. We note also that the candidate young-cluster T dwarf S Ori 70 has an unusually red $H-K_s$ color (Zapatero Osorio et al. 2002) apparently consistent with $\\log g < 4$.   \\subsection{Absolute Magnitudes: The L/T Transition}  It has been apparent since the earliest discoveries of L and T dwarfs that L-type objects evolve into T-type objects as they cool.  With decreasing effective temperatures the condensation level for the principle L dwarf condensates (iron and silicates) falls progressively deeper in the atmosphere. Unless upward mixing is very efficient, the clouds will eventually disappear beneath the photosphere, whose location is strongly wavelength-dependent for these objects. Both the general evolutionary cooling trend and the removal of condensates produce lower atmospheric temperatures; under these conditions the equilibrium chemistry rapidly begins to favor CH$_4$ over CO as the dominant C-bearing species.  With less photospheric condensates to veil the molecular bands and the growing importance of CH$_4$ opacity in the $K$-band, the objects turn to the blue in $J-K$.  Marley (2000) employed a simple, one scale-height thick cloud layer to demonstrate that the sinking of a finite cloud deck explains the red to blue transition in $J-K$.  This was subsequently confirmed by models employing more elaborate cloud models (Marley et al. 2002; Tsuji  2002; Allard et al. 2003).  The absolute magnitudes presented here can be used to better understand this behavior.  Figures 8 and 9 show, respectively, absolute $J$ and $K$ magnitude against spectral type and $J-K$ color.  A fifth order polynomial fit is shown to absolute magnitude against spectral type and the coefficients of the fit are given in Table 12.  Known binaries were removed from the sample before fitting the data; the mean scatter around the fit is 0.4 magnitudes for M$_J$ and 0.3 magnitudes for M$_K$.  As noted by Burgasser et al. (2002b) these data suggest that the L to T dwarf transition may be more complex than implied by the picture of a continuously sinking cloud.  The most notable discrepancy between the simple picture and the data shown in these figures is the brightening seen at $J$-band as the objects transition from L to T.  In this section we summarize various suggested mechanisms to explain this behavior and compare their predictions to the photometric data presented here.  \\subsubsection{Thin Cloud Decks}  A finite-thickness cloud deck forming progressively lower in the atmosphere will eventually disappear from sight.  In the opposite extreme, dust that is well mixed through the entire observable atmosphere, as in the DUSTY models of the Lyon group (e.g. Allard et al. 2001), will by definition never disappear.  Models of such objects show that they simply become progressively redder as they cool and thus, due to veiling of the changing molecular bands, never exhibit an L to T transition.  Although models with finite-thickness cloud decks do move from red, L-like colors to blue, T-like colors, they tend to do so relatively slowly.  This is because the cloud always has a finite thickness and thus does not disappear from a given bandpass instantaneously.  During the time the cloud is departing, say from $J$-band visibility, the overall atmosphere is continuing to cool and become fainter.  Thus in a color magnitude diagram (Figures 8 and 9) models with finite-thickness clouds that are opaque enough to reach the colors of the latest L dwarfs ($J-K\\sim2$) tend to leisurely turn to the blue as they cool and so reach the colors of the early T dwarfs at too faint magnitudes.  The $f_{\\rm sed} = 3$ model in Figure 8 is an example.  Tsuji \\& Nakajima (2003) proposed an interesting solution to this problem. They found a family of models with relatively thin cloud decks in which the turnoff from red to blue in $J-K$ was a function of gravity.  In these models low-gravity 10 $\\rm M_{Jupiter}$ objects depart from what might be called the ``L-type cooling sequence'' (the progressive reddening in $J-K$ with later spectral type) and turn from red to blue at a point almost 2 magnitudes brighter than  high-gravity 70 $\\rm M_{Jupiter}$ objects.  A similar, though less extreme, bright turn off is seen in the $f_{\\rm sed}=5$ family of models in our Figures 8 and 9. Tsuji \\& Nakajima then suggest that there is not a single evolutionary path in which objects first fade at $J$ band as they get redder and then brighten as they turn blue.  Rather they propose that the brighter transition T dwarfs are low mass objects that cooled to mid-L type and then turned from red to blue colors around $M_J \\sim 13.3$.  Dimmer transition objects would represent intermediate-mass brown dwarfs that turned off the L cooling sequence at a later L type and redder $J-K$, and and the latest Ls would represent the highest mass objects.  This model makes a number of interesting predictions.  First, the $M_J$ vs. $J-K$ phase space between the L and T dwarfs should eventually be found to be fairly evenly populated both at brighter and fainter magnitudes than is shown by the transition objects detected to date.  Second, bright early T dwarfs, like SDSS J1021$-$0304 (T3) or 2MASS J0559$-$1404 (T4.5) should be fairly low mass objects while the latest and reddest L dwarfs, like 2MASS J1632$+$1904 (L7.5), should be fairly high mass objects.  Emerging gravity indicators should be able to test this hypothesis.   \\subsubsection{Patchy Clouds}  Plotted in the right panels of Figures 8 and 9 are model sequences from Marley et al. (2002). Both cloud-free and cloudy models are shown, the latter with sedimentation parameters $f_{\\rm sed}=3$ and 5.  For the $f_{\\rm sed}=5$ and the no-cloud models, 3 gravities are shown ($\\log g = 4.5$, 5, and 5.5). Only $\\log g = 5$ is shown for $f_{\\rm sed}=3$. Model effective temperatures are given on the right axis of Figure 9.  The general agreement between the observed L colors and the $f_{\\rm sed}=3$ models seen in Figures 4, 8 and 9, suggests that the $JHK$ colors of the L dwarfs can be explained by a uniform global cloud model.  The cloud's vertical extent and optical depth are limited by sedimentation. Models with much less or much more efficient sedimentation would be generally be too red or too blue, respectively, than most of the L dwarf population.  However the existence of a few dwarfs that are redder and a few that are bluer than most (\\S 5.5) implies that about 10\\% of the L dwarf population would require more extreme models.  These variations could arise from either metallicity or sedimentation efficiency differences between objects.  However, the $f_{\\rm sed} = 3$ models turn too slowly to the blue and reach the colors of the bluest T dwarfs at $J$ band magnitudes that are too faint. Faced with this difficulty of finite cloud layers taking too long to disappear in models like those of Tsuji \\& Nakajima (2003), Burgasser et al. (2002b) propose a different mechanism for the L-to-T transition.  Drawing on a suggestion from Ackerman \\& Marley (2001), Burgasser et al. (2002b) propose that the L to T transition region is marked by the appearance of holes in the global cloud deck, not unlike those seen in the ``5-$\\rm \\mu m$ hot spots'' on Jupiter.  Deeply-seated flux, particularly in the clear $Z$ and $J$ windows would then pour out of these holes, pushing the disk-integrated color to the blue.  Indeed Burgasser et al. (2002b) found that a relatively small fraction of holes would appreciably move a late L-type object towards the blue in $J-K$.  They argued that such a mechanism would explain the brightening observed at $Z$ and $J$, and not at other bands, across the transition, and also the observed resurgence in FeH absorption from the latest Ls to the early to mid Ts.  To illustrate the effect such holes might have,  we have joined with a dotted line the magnitude:color values for the cloudy and cloud-free models at $\\rm T_{eff}=1300$ K in both Figures 8 and 9.\\footnote{Note that Burgasser et al. presented a slightly more sophisticated interpolation, computing colors using a weighted sum of flux produced by clear and cloudy temperature profiles.  The straight lines in Figures 8 and 9 are reasonable approximations to their ``clearing models''.} The agreement between the datapoints for T1 to T3 dwarfs and the cloudy to cloud-free interpolations in Figures 8 and 9 is generally good. The parameters  $\\rm T_{eff}=1300$ K and $\\log g =$ 4.5--5.5 appear to bracket the known transition objects in these plots, in fair agreement with Golimowski et al. (2004) who show that there is an apparent plateau at $\\rm T_{eff}\\approx$ 1450~K for types L7 to T4. The cloud clearing model simply posits that at a given effective temperature the global cloud deck begins to break up, perhaps because it has settled sufficiently deeply into the convection zone that it becomes subject to the global circulation pattern. The observed constancy of $\\rm T_{eff}$ across the transition is not required by this hypothesis, although it does raise problems for the Tsuji \\& Nakajima (2003) suggestion of continuous cooling across the transition.  If the clearing does happen over a narrow temperature range, with a large spectroscopic change occurring over a small change in temperature and luminosity as the brown dwarf cools, we would not expect to discover many early T dwarfs. However about one-third of our T dwarf sample is made up of types T0--T3.5, with a possible dearth of T3--4 types (Figure 3). A study of the spectral type distribution in an SDSS magnitude-limited sample will be the subject of a future paper (Collinge et al. 2002).  \\subsubsection{Sudden Downpour}  Rather than relying on spatial inhomogeneities, Figures 8 and 9 suggests a third possibility for the L to T transition, which we term the ``sudden downpour'' model.  The $f_{\\rm sed} = 3$ models do a reasonably good job of reproducing the colors of the latest L dwarfs.  Like the thin Tsuji \\& Nakajima (2003) cloud, models with more efficient sedimentation (larger $f_{\\rm sed}$) turn off the L dwarf cooling track sooner (at brighter magnitudes) than seems to be consistent with the available data.  However, one might argue that L dwarfs first cool at essentially constant $f_{\\rm sed}$, then at around $\\rm T_{eff} = 1300\\,\\rm K$, $f_{\\rm sed}$ begins to gradually increase from $\\sim 3$ to infinity at roughly fixed effective temperature.  This rapid increase in the efficiency of sedimentation would, in essence, produce a torrential rain of condensed iron and silicate grains.  Unlike the Tsuji \\& Nakajima (2003) mechanism this would begin at the late L spectral type for all masses. T1 to T4 dwarfs would represent different stages of this cloud thinning process. Figures 8 and 9 shows where the $f_{\\rm sed}=5$ models would lie, for example.  Once grains are essentially completely removed from the atmosphere the object would continue to evolve and cool.  A useful diagnostic for evaluating the various transition models may be gravity. The sudden downpour mechanism, for example, would predict that the T3.5 dwarf SDSS J1750$+$1759 (see Figure 8) would have $\\log g \\sim 5.4$. Evolution tracks for the patchy cloud model curve downward at blue $J-K$ compared to the straight lines shown in the figure and thus this model would predict a smaller gravity, say $\\log g = 5$.  On the other hand, Tsuji \\& Nakajima (2003) would predict that since this relative bright T dwarf has already made the transition to blue $J-K$, it must be relatively low in mass and have a substantially lower gravity, say $\\log g = 4$.  There are also gravity tests among the late L dwarfs, although these are more subtle since all of the models would predict that the faintest and reddest L dwarfs will have progressively higher masses.  Tsuji \\& Nakajima (2003) predict that the lowest mass dwarfs turn to the blue relatively early.  The downpour model also predicts that lower masses turn blueward earlier than higher masses as can be seen by studying the tracks for the various gravities in the $f_{\\rm sed}=5$ case. However this turn off happens at later spectral types than in the Tsuji \\& Nakajima model and is accompanied by a subsequent brightening in $M_J$.  For example the L7 dwarf labeled in Figure 8 could have $\\log g \\ge 4.8$ under the cloud clearing and downpour models, but Tsuji \\& Nakajima (2003) would predict a higher minimum gravity, likely around 5.3 or so. Certainly there are hints in Figure 8 that there is a width in $M_J$ to the transition and this will facilitate such tests.  There are also other diagnostics to consider.  The patchy-cloud model straightforwardly accounts for the resurgence seen in FeH absorption across the transition (Burgasser et al. 2002b), and it is not clear that other models can account for this.  Clearly more modeling of all mechanisms must be completed to better define gravity and other diagnostics of the transition mechanism.  Unfortunately gravity indicators among the late L and early T field dwarfs are elusive and it may be difficult to use them to definitively test the models. (One might expect that the Li test (Mart\\'{\\i}n, Rebolo \\& Magazzu 1994) could be used to identify high gravity dwarfs since only brown dwarfs more massive than 0.06~M$_{\\odot}$ will have burned Li during their evolution.   As a practical matter, however, the Li line at 0.6708~$\\mu$m in late L and early T dwarfs is detectable only with the largest telescopes due to the lack of continuum flux in this region. Also, in T dwarf atmospheres LiCl and LiOH become the dominant Li bearing molecule (Lodders 1999) and these species are currently undetectable.) Self consistent evolutionary models be developed as well as mechanisms to explain the onset of either patchiness or varying $f_{\\rm sed}$ at a particular $\\rm T_{eff}$.  We plan to more fully explore these issues in a future paper.   \\subsubsection{Additional Considerations}  Discrepant objects to note in Figures 8 and 9 are Kelu-1 (L3), SDSS J0423$-$0414 (T0) and 2MASS J0415$-$0935 (T9).  2MASS J0415$-$0935 is significantly redder in $J-K$ than the models would predict, as discussed in \\S 5.6.  Kelu-1 and SDSS J0423$-$0414 are both superluminous by about 0.75 magnitudes, suggesting that they may be pairs of identical dwarfs in unresolved binary systems. Kelu-1 has been imaged by $HST$ and by Keck, with no evidence of duplicity found (Mart\\'{\\i}n et al. 1999a; Koerner et al. 1999), while SDSS J0423$-$0414 has not been imaged at high resolution to our knowledge.  Cruz et al. (2003) classify this dwarf as L7.5 using red spectra; however, CH$_4$ bands are clearly seen in the G02 spectrum (Figure 3 in G02), and the bolometric correction (i.e. ratio of $K$-band flux to total luminosity) is more compatible with an infrared classification of T than L (Golimowski et al. 2004).  Although it has been suggested that this object consists of a late L and early T close pair (Burgasser et al. 2003b) a discrepancy between the optical and infrared types is not unexpected.  As described in \\S 5.2, different wavelength regions probe different levels of the photosphere. It is likely that the optically derived spectral types are more representative of effective temperature, and as Golimowski et al. (2004) show that there is an apparent plateau at $\\rm T_{eff}\\approx$ 1450~K for types L7 to T4 we would expect optical types to be earlier than infrared types for the L/T transition objects.  The earlier optical classification is in fact observed for six L8 to T0 dwarfs in our sample, as indicated in Table 9.  Although the clear atmosphere models do a good job of reproducing the colors of the later T dwarfs in the $J-H:H-K$ diagram (Figure 6), they tend to predict bluer $J-K$ colors for these objects than is observed (Figures 8 and 9).  This seems to be a generic problem with clear models (see also Tsuji \\& Nakajima 2003).  The $J-K$ result implies that the model tracks shown in Figure 6 should slide up and to the right, consistent with the suggestion in \\S 5.6 that the temperature contours are too warm.  Because of the overall shape of the model contours the quality of the fit in $J-H:H-K$  would remain about the same.  The discussion of the gravity signature seen in $H-K$ is still qualitatively valid, although SDSS J1110$+$0116 would be expected to have a somewhat higher gravity.  Finally, we consider detectability limits for SDSS using Figures 3, 8 and 9. Many SDSS T dwarfs have to be selected as $z$-only objects and they are either not detected or only barely detected at $i$ (see Table~1).  The nominal (5$\\sigma$, better than $1.5''$ seeing) $z$ detection limit is 20.8, and the faintest SDSS T dwarf discovered to date is close to this limit, at $z$ = 20.4.  Since the colors of late (e.g. T8) dwarfs are $z-J\\approx$3.8, the corresponding $J$ limit is $\\approx$16.6; Table~9 therefore indicates that SDSS should be able to detect T9 dwarfs to 10 pc.  The latest SDSS dwarf discovered to date is a T7, but we anticipate that some later types will be found."
    },
    "0402/astro-ph0402498_arXiv.txt": {
        "abstract": "We present deep Ka-band ($\\nu \\approx 33$~GHz) observations of the CMB made with the extended Very Small Array (VSA). This configuration produces a naturally weighted synthesized FWHM beamwidth of $\\sim 11$~arcmin which covers an $\\ell$-range of 300 to 1500. On these scales, foreground extragalactic sources can be a major source of contamination to the CMB anisotropy. This problem has been alleviated by identifying sources at 15~GHz with the Ryle Telescope and then monitoring these sources at 33~GHz using a single baseline interferometer co-located with the VSA. Sources with flux densities $\\gtsim 20$~mJy at 33~GHz are subtracted from the data. In addition, we calculate a statistical correction for the small residual contribution from weaker sources that are below the detection limit of the survey.  The CMB power spectrum corrected for Galactic foregrounds and extragalactic point sources is presented. A total $\\ell$-range of $150-1500$ is achieved by combining the complete extended array data with earlier VSA data in a compact configuration. Our resolution of $\\Delta \\ell \\approx 60$  allows the first 3 acoustic peaks to be clearly delineated. The is achieved by using mosaiced observations in 7 regions covering a total area of 82 sq. degrees. There is good agreement with WMAP data up to $\\ell=700$ where WMAP data run out of resolution. For higher $\\ell$-values out to $\\ell = 1500$, the agreement in power spectrum amplitudes with other experiments is also very good despite differences in frequency and observing technique. ",
        "introduction": "\\label{sec:introduction}  The angular power spectrum of primordial anisotropies in the cosmic microwave background (CMB) has become an important tool in the era of ``precision cosmology''. Since the first statistical detection of CMB fluctuations on large angular scales ($\\ell =2-30$) by the $COBE$-DMR instrument (Smoot et al. 1992), several experiments have detected acoustic peaks in the power spectrum in the $\\ell$-range $100-1000$ (Lee et al. 2001; Netterfield et al. 2002; Halverson et al. 2002; Beno{\\^i}t et al. 2003; Scott et al. 2003) and a fall-off in power at high $\\ell$-values  (Dawson et al. 2002; Grainge et al. 2003; Kuo et al. 2004; Readhead et al. 2004). More recently, the Wilkinson Microwave Anisotropy Probe, henceforth $WMAP$, has provided unprecedented measurements over the $\\ell$-range $2-700$ (Bennett et al. 2003a; Hinshaw et al. 2003a). The WMAP 1-year power spectrum is cosmic variance limited up to $\\ell=350$ and delineates the first 2 peaks at $\\ell \\sim 220$ and $\\ell \\sim 550$ with exceptional signal-to-noise ratios. The new data have provided detailed cosmological information on a wide range of parameters (Spergel et al. 2003) and have raised new questions to be answered. However, the angular resolution of $WMAP$ limits the power spectrum to $\\ell \\ltsim 800$. At high $\\ell$-values, the origin of the excess power initially observed by the Cosmic Background Imager (CBI) at $\\ell = 2000-4000$ (Mason et al. 2003) has also generated much interest. It is clear that the high-$\\ell$ CMB power spectrum is one of the challenges for future CMB experiments including the ESA Planck satellite (Tauber 2001; Lawrence 2003) due for launch in 2007.  The Very Small Array, henceforth the VSA (Watson et al. 2003, hereafter Paper~I), is a purpose-built radio interferometer that has measured the CMB angular power spectrum between $\\ell = 150$ and 900 in a compact array configuration  (Scott et al. 2003, Paper III) and more recently up to $\\ell=1400$ in an extended array (Grainge et al. 2003). This paper describes the complete set of extended array data observed during October 2001 $-$ June 2003. The data cover a larger area of sky than those analysed by Grainge et al. (2003) by adding more pointings to the previous VSA fields and by observing 4 new fields. We use mosaicing techniques to increase the sensitivity, reduce the sample variance and to facilitate finer $\\ell$ resolution $-$ or equivalently reduce correlations between $\\ell$ bins $-$ over the range $\\ell=300-1500$.  The paper is organised as follows. Section~\\ref{sec:extended_array} summarises the telescope parameters and the extended array configuration. In section~\\ref{sec:abs_and_data} we describe the observations, data reduction and calibration of the data, including a range of data checks. In section \\ref{sec:foregrounds} we discuss the various foregrounds, particularly discrete radio sources, and the corrections that were made to the data. The main results, CMB mosaic maps and CMB power spectrum covering $\\ell=150-1500$, are presented in section~\\ref{sec:results}. A morphological analysis of the power spectrum is given in section~\\ref{sec:ps_gaussians}. Section~\\ref{sec:discussion} is a discussion of the results and comparisons with other data, followed by conclusions in section~\\ref{sec:conclusions}. The cosmological interpretation of these data, both on their own and combined with other data, is described in a companion paper (Rebolo et al. 2004). ",
        "conclusions": "\\label{sec:conclusions}  We have presented high sensitivity and foreground subtracted measurements of the CMB power spectrum up to $\\ell =1500$ observed with the VSA. The cosmological interpretation is described in a companion paper (Rebolo et al.~2004). The final extended array data set contains a factor of $\\sim 4$ more than those presented by Grainge et al. (2003). The sky-coverage has increased by a factor of $\\sim 3$ utilising 4 additional 3-field mosaics and increasing the original 3-field mosaics to 7-field mosaics, resulting in a significant increase to the signal-to-noise ratio. Mosaicing techniques allow for smaller correlations between adjacent bins and hence, can give finer resolution in $\\ell$. However, in this paper we have used the same binning scheme as Grainge et al. (2003) for comparison and to reduce the bin-bin correlations ; the $\\ell$ resolution is $\\Delta \\ell \\approx 60$. We adopt a new absolute calibration scheme based on Jupiter temperature measurements with WMAP to give a 3 per cent accurate absolute scale in the power spectrum. This reduces absolute values by 4 per cent in temperature or 8 per cent in the power spectrum compared to previous VSA data.  The VSA power spectrum (Fig.~\\ref{fig:ps}) is in good agreement with data from other CMB experiments. There are now good signal-to-noise detections of the first 3 peaks in the CMB power spectrum and a damping tail at high $\\ell$ which continues to $\\ell \\sim 1500$. The different instruments have different potential systematic problems and they cover a wide range of frequencies. The agreement therefore indicates that none of the experiments are seriously contaminated by foreground emission or other systematic effects (see also Griffiths \\& Lineweaver (2004)). From the analysis of the peak structure (section~\\ref{sec:ps_gaussians}), there is extremely good agreement on the peak structure between experiments, except for a slight discrepancy in the amplitude of the 3rd peak which may have important consequences for the cosmological models. The slight discrepancy ($\\sim 1.6\\sigma$) in the overall scaling between VSA and WMAP may also be an important factor in terms of fitting cosmological models. This is discussed by Rebolo et al. (2004).  In the near future, the VSA will be reconfigured with even larger horns and longer baselines. It will allow the VSA to increase the $\\ell$-range even further, up to a maximum of $\\ell \\sim 3000$, while maintaining good $\\ell$ resolution without loss of overall temperature sensitivity. The assistance of high frequency radio source surveys, to identify and measure the flux densities of sources over large areas of sky, will be a crucial role in keeping the discrete source foregrounds under control. The One Centimetre Receiver Array (OCRA; Browne et al. 2000), a prototype multi-beam receiver system currently begin tested on the Torun 32~m telescope, will be ideal for such purposes. Moreover, OCRA is currently operating at $30$~GHz and hence direct source subtraction, without the need to extrapolate to other frequencies will be possible in the near future."
    },
    "0402/astro-ph0402337_arXiv.txt": {
        "abstract": "We report the discovery of a young massive stellar cluster and infrared nebula in the direction of the CS molecular cloud associated to the IRAS point source 16132-5039. The analysis of the mid-infrared images from the more accurate MSX catalog, reveled that there are two independent components associated with the IRAS source. The integral of the spectral energy distribution for these components, between 8.28 $\\mu$m and 100 $\\mu$m, gave  lower limits for the bolometric luminosity of the embedded objects of $8.7 \\times 10^4 L_\\odot$ and $9 \\times 10^3 L_\\odot$, which corresponds to  ZAMS O8 and B0.5 stars, respectively. The number of Lyman continuum photons expected from the stars that lie along the reddening line for early-type stars is about 1.7 $ \\times$ 10$^{49}$ s$^{-1}$, enough to produce the detected flux densities at 5 GHz. The NIR spectrum of the nebula increases with frequency, implying that free-free emission cannot be the main source of the extended luminosity, from which we conclude that the observed emission must be mainly dust scattered light. A comparison of the cluster described in this paper with the young stellar cluster associated with the IRAS source 16177-5018, which is located at the same distance and direction, shows that the mean visual absorption of the newly discovered cluster is about 10 magnitudes smaller and it contains  less massive stars, suggesting that it was formed from a less massive molecular cloud. ",
        "introduction": "Massive stars are born within dense molecular clouds forming clusters and associations; during their formation and early evolution they are often  visible only at infrared and radio wavelengths. With the advent of the large bidimensional near-infrared array detectors, the morphological and photometric studies of extremely young galactic stellar clusters have benefited. At near-infrared wavelengths (1 to 2.5 $\\mu$m) it is possible to probe deep into the dense dust clouds where star formation is taking place.  The strong ultraviolet (UV) radiation emitted by the massive stars dissociates and ionizes the gas, forming compact HII regions seen at radio wavelengths. A large fraction of the  radiation could also heat the dust, which eventually radiates  in the far-infrared (FIR); for this reason, compact HII regions are among the brightest and most luminous objects in the Galaxy at 100$\\mu$m.  Because massive stars evolve very fast, they are extremely rare and difficult to find, except when related to the emission of the surrounding molecular cloud. In that sense, CS and NH$_3$ lines at radio frequencies, characteristic  of high density gas,  are good  tracers of massive star forming regions.  In this paper we present observations in the near infrared,  of a young stellar cluster of massive stars in the direction of  the IRAS point source I16132-5039. This work is a part of a survey aimed to the identification of stellar populations in the direction of IRAS sources that have colors characteristics of ultracompact HII regions \\citep{wood89} and strong CS (2-1) line emission \\citep{bronf96}. The studied region is located in the  direction of another massive star forming region associated to the IRAS point source 16177-5018 \\citep{rom03}; they are part of the RCW 106 complex, located in the southern Galactic plane at a distance of 3.7 kpc \\citep{cas87}. The near-infrared cluster presented here had recently been located  by Dutra et al. (2003) by visual inspection of the 2MASS images; however their work contains only the possible cluster location, and as Persi, Tapia \\& Roth (2000) already showed  for NGC6334IV,  a high concentration of K band sources can be due to localized low extinction and be mistaken with a stellar cluster, leading to false identifications.   The cluster is associated with the radio source 332.541-0.111,  which presents continuum radio emission at 5 GHz as well as hydrogen recombination lines \\citep{cas87}. The observations and data reduction are described in section 2, the results are presented in section 3 and our main conclusions are sumarized in section 4. ",
        "conclusions": "The combined false-colour infrared image, ($\\it{J}$ blue, $\\it{H}$ green and nb$\\it{K}$ red) of the whole field is displayed in Figure 2, together with amplified individual images of the nebular region at all bands. All of them, but especially  the $\\it{H}$ and nb$\\it{K}$ images, shows the presence of a small  spheroidal nebula with a bright star at its center.  \\subsection{The IRAS source}  In Figure 3 we present a contour map constructed from the LNA nb$K$ image, with a beam size of $2\\times 2$ pixels which shows the region around IRAS16132-5039. The contours start at 2.2$\\times 10^{-4}$ Jy/beam and are spaced by this same value. The IRAS coordinate  has an intrinsic error  delimited by the ellipse plotted in the figure. A more accurate position for the IR source was obtained from the Midcourse Space Experiment - MSX point source catalog $(psc)$ \\footnote{ http://www.ipac.caltech.edu/ipac/msx/msx.html}. The MSX surveyed the entire Galactic plane within $\\mid b \\mid \\leq 5^\\circ$ in four mid-infrared spectral bands centered at 8.28, 12.13, 14.65 and 21.34 $\\mu$m, with image resolution of 19 arcsec and a global absolute astrometric accuracy of about 1.9 arcsec \\citep{price01}. We found one MSX source, with coordinates $\\alpha(\\rm{J2000)=16^{h}17^{m}02.47^{s}}$, $\\delta(\\rm{J2000)=-50^{d}47^{m}03.5^{s}}$ within the IRAS uncertainty ellipse; its coordinates coincides with the star we labeled IRS1.  However, looking at the MSX images, we found another closeby source, although outside the IRAS uncertainty ellipse, with coordinates $\\alpha(\\rm{J2000)=16^{h}16^{m}55.94^{s}}$, $\\delta(\\rm{J2000)= -50^{d}47^{m}07.8^{s}}$. In Figure 4 we present our $H$ band image  overlaid with the 8.28$\\mu$m band MSX contour diagram, with the contours spaced by $ 8\\times 10^{-6}$ W m$^{-2}$ sr$^{-1}$, starting at $2.8\\times 10^{-5}$ W m$^{-2}$ sr$^{-1}$. The same source was found in the contour plots from the other MSX bands. We designated the stronger source as A and the other as B. While the IRS1 object and the infrared nebulae are related to source A , source B  is also associated with a small nebular region that shows a concentration of NIR sources, as can also be seen in Figure 4.  Since only the 8.28 $\\mu$m flux density was given in the MSX $psc$ for both sources, we integrated the flux density of each individual source for the four mid-infrared bands;  the results are presented in table 1, which also shows the values from  MSX and IRAS catalogs. Our integrated flux density for source B in the 8.28 $\\mu$m coincides within 5\\% with the value given in the MSX catalogue, but it is 30\\% larger for the A source. This discrepancy can be understood if we consider that the automatic MSX algorithm subtracted part of the source B flux density as background contribution. The relative contribution of source B decreases with increasing wavelength, explaining also why it was not resolved by the MSX algorithm. It should be noticed that the reported IRAS flux density at 12 $\\mu$m coincides with the sum of our derived flux densities at 12.13 $\\mu$m within 10\\%, while the MSX value was about 50\\% lower.  In Figure 5 we plotted the mid to far-infrared spectral energy distribution of the sources A and B, without any correction for absorption. We assumed that the IRAS flux density in the far infrared is divided between the two sources in the same way as in the mid-infrared (14.65 and 21.33 $\\mu$m): about 90\\%  originating from  source A and 10\\%  from source B. We then integrated the observed flux densities in the mid-far infrared, assuming a distance of 3.7 kpc, obtaining a luminosity $L_{\\rm A}=8.7 \\times 10^4 L_\\odot$ and $L_{\\rm B}=9 \\times 10^3 L_{\\odot}$ for sources A and B, respectively. Assuming that the IR flux density represents a lower limit for bolometric luminosity of the embedded stars, we derived ZAMS  spectral types O8 and B0.5   for the  sources A and B, respectively  (Hanson et al. 1997).    \\subsection{Cluster population}  In order to examine the nature of the stellar population in the direction of the IRAS source, we analyzed the stars in two delimited regions: one that we labeled \"cluster\", which contains the nebula and another that we labeled \"control\", which has a stellar population that we believe is dominated by \"field\" stars, as illustrated in Figure 6. We represented the position of all objects detected in the $H$ band by crosses and we can see that the small region labeled \"cluster\" shows a concentration of sources. In figure 7 we present comparative ($J-H$) versus ($H-K$) diagrams for the stars detected in  the  $J$, $H$ and $K$ images, together with the position of the main sequence, giant branch and reddening vectors for early and late type stars (Koornneef 1983).  We see that the stellar population of the \"cluster\" region is quite different from that of the \"control\" region, with many sources lying on the right side of the reddening vector for early type stars, showing excess emission at 2.2 $\\mu$m. In the \"control\" region the majority of the sources have colors of reddened photospheres, with many objects located along the reddening vector for late type stars. It is well established that very young pre-main sequence objects present large infrared excess due to the presence of warm circumstellar dust (Lada \\& Adams 1992). Our results suggest that the stellar population in the direction of the IRAS source is very young, as can be inferred from their position in the $(J-H)$ versus $(H-K)$ diagram.  We separated the cluster sources from the field stars, by selecting all sources that lie to the right  or on the reddening vector for early type stars in the cluster's color-color diagram. Table 2 shows the coordinates and photometry  of all selected sources ($J$, $H$ and $K$ magnitudes from both LNA and 2MASS surveys).   Further information about the nature of the selected objects in Table 2 can be extracted from $J$ versus $(J-H)$ color-magnitude diagram shown in Figure 8. We used this diagram instead the $K$ versus $(H-K)$ color-magnitude diagram  to minimize the efect of the \"excess\" of emission in the NIR on the derived stellar spectral types. The locus of the main-sequence for class V stars at 3.7 kpc (Caswell \\& Haynes 1987) is also  plotted, with the position of each spectral type earlier than A0 V indicated. The intrinsic colors were taken from Koornneef (1983) while the absolute $J$ magnitudes were calculated from the absolute visual luminosity for ZAMS taken from Hanson et al. (1997). The reddening vector for a ZAMS B0 V star, taken from Rieke $\\&$ Lebofsky (1985), is shown by the dashed line with the positions of  visual extinctions  $A_V = 10$ and 20 magnitudes indicated by filled circles. We also indicated the sources with and without \"excess\" in the color-color diagram by open and filled triangles respectivelly.  For the assumed distance, we estimated the spectral type of the stars that do not present excess emission in the near infrared, by following the de-reddening vector in the color-magnitude diagram, for the others we only gave a rough classification; the results are shown in the last column of Table 2.  The main source of error in the derived spectral types arises from uncertainties in the assumed cluster distance, which was derived from the velocities of the radio hydrogen recombination lines. Since the closest distance was used, the main uncertainties come from the errors in the galactic rotation curve model and related parameters. From the works of Blum et al. (1999, 2000, 2001) and Figuer\\^edo et al. (2002) comparing kinematic with spectroscopic distances we find that they do not differ in more than 1 kpc, in which case the change in the luminosity class would be of two sub-spectral types for early O stars and one for early B stars.  We must notice that source IRS1, which is associated with  MSX source A, has an estimated spectral type of at least O5, reddened by about $A_{V} \\approx$ 14 magnitudes. Besides, there are five objects (IRS3, IRS11, IRS18, IRS21 and IRS33) associated with the mid-infrared source B; IRS3 is probably an  O8 ZAMS star reddened by about $A_{V}$= 7 magnitudes, while the others have estimated spectral types of early-B stars. The corresponding  bolometric luminosities are in agreement with the lower limits derived from the integrated mid-far infrared flux densities, corresponding to O8 and B0.5, for sources A and B respectively, as seen in section 3.1.   A lower limit to the number of Lyman-continuum photons produced in the star forming region, can be calculated taking into account only the stars that do not show \"excess\" in the color-magnitude diagram (IRS1, 6, 8, 9, 10, 13, 14, 16, 20, 23, 26, 27 and 35 in Table 2). It was computed from the relation  given by Hanson et al. (1997), resulting in $1.7 \\times 10^{49}$ photons s$^{-1}$. Is interesting to note that IRS1 is responsible for more than 90\\% of the Lyman continuum photons, being the main ionization source in the whole region.   It is also possible to obtain  the number of ionizing Lyman continuum photons $N_{Ly}$ from the radio continuum flux density given by Caswell \\& Haynes (1987), using the expression derived by Rubin (1968): \\begin{equation} N_{Ly}= \\frac{5.59\\times 10^{48}S(\\nu) D^2 T_e^{-0.45} \\nu^{0.1}}{1+f_i[<He^+/(H^++He^+)]} \\end{equation} where  $\\nu$ is in units of 5 GHz and $f_i$ is the fraction of  helium recombination photons that are energetic enough to ionize hydrogen. Using the values of $T_{e} = 4500$ K, $S(\\nu) = 3.3$ Jy and $D$= 3.7 kpc taken from Caswell \\& Haynes (1987), we find that $f_i\\approx 0$ and $N_{Ly} \\sim 6 \\times 10^{48}$ photons s$^{-1}$, compatible with the lower limit derived from the observed stars.    \\subsection{The Infrared Nebula}   We can see from the detailed  $J,$ $H$ and nb$K$ images in figure 2 that the nebula presents a spheroidal shape, approximately symetric around the IRS1 source. In figure 9 we shown contour maps of the  nebular region at $\\it{J, H}$ and nb$\\it{K}$ bands obtained from our infrared images. The contours were calibrated in flux with the values starting at 0.44 ($J$), 1.6 ($H$) and 2.2$\\times 10^{-4}$ Jy/beam ($K$), with intervals of 0.37, 1.8 and 1.1$\\times 10^{-4}$ Jy/beam respectively. We estimated the total flux density by measuring the area between contours and multiplying by the value of the corresponding flux density per unit area, obtaining $S(J)=0.023$ Jy, $S(H)=0.07$ Jy and $S(K)=0.11$ Jy. We then corrected these results for extinction, using the mean value $<A_{V}>$ = 15.1 magnitudes, derived from the stars in the direction of the nebula that do not present infrared excess (IRS1, IRS6 and IRS10), and the standard extinction law  from Rieke \\& Lebofski (1985). We obtained $S(J)=1.16$ Jy, $S(H)=0.8$ Jy and $S(K)=0.52$ Jy.  In the previous section we showed that the number of ionizing photons available from the detected stars is enough to explain the radio continuum flux density measured by Caswell \\& Haynes (1987). We will investigate now the contribution of free-free emission to the observed nebular IR flux density. Assuming constant density and temperature across the cloud and  local thermodynamic equilibrium, the flux density $S(\\nu)$ due to free-free emission can be written as: \\begin{equation} S(\\nu)=\\tau_\\nu B_\\nu(T)\\Omega \\end{equation} where $\\tau_\\nu$ is the optical depth at frequency $\\nu$,  $B_\\nu(T)$ is the Planck function \\begin{equation} B_\\nu(T)=\\frac{2h\\nu^3}{c^2}\\frac{1}{{\\rm exp}(h\\nu/kT)-1} \\end{equation} and $\\Omega$ is the solid angle of the source given by: \\begin{equation} \\Omega=\\pi{(L/2D)}^2 \\end{equation} where $L$ is the diameter of the ionized cloud and $D$ the distance to the observer. The optical depth at a given frequency $\\nu$ is: \\begin{equation} \\tau_\\nu=\\alpha_{\\rm ff}L \\end{equation} where $\\alpha_{\\rm ff}$ is the free-free absorption coefficient $\\rm (cgs)$ taken from Rybick (1979): \\begin{equation} \\alpha_{\\rm ff}=\\frac{3.7\\times 10^8\\;[1-{\\rm exp}(-h\\nu/kT)]\\;n_e n_i\\;g_{\\rm ff}(\\nu,T)}{\\nu^3\\;Z^{-2}\\; T^{1/2 }} \\end{equation} where $g_{\\rm ff}(\\nu,T)$ is the Gaunt factor  obtained from Karzas \\& Latter (1961).  For two frequencies $\\nu_1$ and $\\nu_2$ the ratio of the corresponding flux densities $S(\\nu_1)$ and $S(\\nu_2)$ may be calculated from: \\begin{equation} \\frac{S(\\nu_1)}{S(\\nu_2)}={exp}[h(\\nu_2-\\nu_1)/kT]\\frac{g_{\\rm ff}(\\nu_1,T)}{g_{\\rm ff}(\\nu_2,T)} \\end{equation}  For a flux density in the 5 GHz continuum of 3.3 Jy and an electron temperature $T_e = 4500$ K, as given by Caswell \\& Haynes (1987) the expected flux densities at the $J, H$ and $K$ bands are 0.05, 0.10 and 0.16 Jy, respectively. We verify that the measured values are much larger than what was expected from the free-free emission, derived from the radio data. In fact, only an absorption as low as 4 magnitudes would explain the inferred spectrum as free-free emission. Since in the direction of the A source, this absorption is incompatible with even the less absorbed star (IRS1), we believe that this is not the main source of extended IR emission.  Accepting the mean value of $<A_{V}>$ = 15.1 magnitudes as the mean value for the absorption in the direction of A source, we found that the corrected flux density  increases with  frequency, suggesting that the observed extended radiation is scattered light from the nearby stars. We adjusted then a black body to the NIR fluxes, obtaining a good fit for T$\\approx$ 16000K,  characteristic of middle-B  stars (Hanson et al. 1997). Lumsden $\\&$ Puxley (1996), analyzing the ultracompact HII region G45.12+0.13, also obtained an extinction corrected flux density that increases with decreasing wavelengths and interpreted it as due to stellar light, scattered by dust through the HII region.     \\subsection{Conclusions}  Near-IR imaging in the direction of the CS molecular cloud associated with the IRAS source 16132-5039, revealed an embedded young massive stellar cluster. We detected 35  member candidates up to our completeness limit, concentrated in an area of about 2 square parsec. All images, but especially  the $\\it{H}$ and nb$\\it{K}$ bands, show the presence of a small  spheroidal nebula with a bright star (IRS1) at its center.  The stars associated with the IRAS point source were identified using more accurate positions from the MSX catalogue. The analysis of the mid-infrared images reveled that there are two  sources associated with IRAS 16132-5039. The strongest coincides with  the position of at least a dozen of OB stars, while the weaker source is associated to less massive objects, with spectral types characteristic of middle-B ZAMS stars. The integral of the spectral energy distribution of the MSX-IRAS sources, between 8.28 $\\mu$m and 100 $\\mu$m, gives lower limits to the bolometric luminosity of the embeded objects of $L=8.7 \\times 10^4 L_\\odot$ and $L=9 \\times 10^3 L_\\odot$, which corresponds to  ZAMS O8 and B0.5 stars, respectively (Hanson et al. 1997). The results are compatible with the spectral types of the objects detected in the NIR, since it is possible that only  part of the energy emitted by the stars is reprocessed by the dusty envelope. In that sense, they can be taken as lower limits to the bolometric luminosity of the embeded stars.  Assuming that the radio emission measured by Caswell \\& Haynes (1987) originates in this region, at a distance of 3.7 kpc, we estimated the number of ionizing Lyman continuum photons  as $N_{Ly} \\sim 6 \\times 10^{48}$ photons s$^{-1}$. On the other hand, the number of Lyman continuum photons expected from the stars that lie along the reddening line for early-type stars is about 1.7 $ \\times$ 10$^{49}$ s$^{-1}$, enough to produce the detected flux densities at 5 GHz. The IRS1 source is enough to account for more than 90\\% of the total number of Lyman continuum photons necessary to ionize the gas.  Analysis of the integrated  flux densities of the NIR  nebula at the $J$, $H$ and nb$K$ bands revealed that they increase with frequency, implying that free-free emission cannot be the main source of the extended luminosity, unless we assume only four magnitudes of visual extinction. Since this value is incompatible with the extinction  derived from the stars that do not shown excess of emission at 2.2 $\\mu m$, we conclude that the observed emission must be mainly dust scattered light.  A comparison of the cluster described in this paper with the young stellar cluster associated with the IRAS source 16177-5018 (Roman-Lopes et al. 2003), which is located at the same distance and direction, shows that the the former contains  less massive stars. Since its mean visual absorption is also about 10 magnitudes smaller than that of IRAS 16177-5018, it is possible that it was formed from a less massive molecular cloud."
    },
    "0402/astro-ph0402101_arXiv.txt": {
        "abstract": "s{Recent observations with {\\it Chandra} and {\\it XMM-Newton}, aided by broad-band spectral coverage from {\\it RXTE}, have revealed skewed relativistic iron emission lines in stellar-mass Galactic black hole systems.  Such systems are excellent laboratories for testing General relativity, and relativistic iron lines provide an important tool for making such tests.  In this contribution to the Proceedings of the 10th Annual Marcel Grossmann Meeting on General Relativity, we briefly review recent developments and present initial results from fits to archival {\\it ASCA} observations of Galactic black holes.  It stands to reason that relativistic effects, if real, should be revealed in many systems (rather than just one or two); the results of our archival work have borne-out this expectation.  The {\\it ASCA} spectra reveal skewed, relativistic lines in XTE J1550$-$564, GRO~J1655$-$40, GRS~1915$+$105, and Cygnus X-1.} ",
        "introduction": " ",
        "conclusions": "In the {\\it Chandra} and {\\it XMM-Newton} era, relativistic Fe~K$\\alpha$ emission lines have been revealed in stellar-mass Galactic black holes.  They have proved to be robust diagnostics of the innermost accretion flow region, which can also constrain the nature of the central black hole.  Archival analyses of the kind briefly detailed here show that relativistic lines are in fact common, and not limited to a few special sources.  In the near future, we look forward to the high effective area, spectral resolution, and broad energy range of {\\it ASTRO-E2}, to further refine studies of this kind.  In the long run, {\\it Constellation-X} will make it possible to study changes in Fe~K$\\alpha$ line profiles over short timescales; such investigations will be especially powerful probes of the corona --- disk interaction in Galactic black hole systems.   \\begin{figure}[htbp] \\epsfxsize=8cm   % \\centerline{\\epsfxsize=4.5in\\epsfbox{lines.eps}} \\caption{Prominent relativistic line profiles in Galactic black holes observed with {\\it ASCA}.  The plots above show the ratio of GIS spectra to a phenomenological continuum model consisting of multicolor disk blackbody, power-law, and smeared edge components.  A narrow Fe XXV/XXVI absorption line and narrow Ca XX absorption line were fit to GRO J1655$-$40 and GRS~1915$+$105.\\label{asca_lines}} \\end{figure} \\eject"
    },
    "0402/astro-ph0402271_arXiv.txt": {
        "abstract": "We present a new set of stellar interior and synthesis models for predicting the integrated emission from stellar populations in star clusters and galaxies of arbitrary age and metallicity. This work differs from existing spectral synthesis codes in a number of important ways, namely (1) the incorporation of new stellar evolutionary tracks, with sufficient resolution in mass to sample rapid stages of stellar evolution; (2) a physically consistent treatment of evolution in the HR diagram, including the approach to the main sequence and the effects of mass loss on the giant and horizontal-branch phases. Unlike several existing models, ours yield consistent ages when used to date a coeval stellar population from a wide range of spectral features and colour indexes. We use {\\em Hipparcos} data to support the validity of our new evolutionary tracks. We rigorously discuss degeneracies in the age-metallicity plane and show that inclusion of spectral features blueward of 4500 \\AA\\, suffices to break any remaining degeneracy and that with moderate $S/N$ spectra (10 per 20\\AA\\, resolution element) age and metallicity {\\em are not} degenerate. We also study sources of systematic errors in deriving the age of a single stellar population and conclude that they are not larger than $10-15$\\%. We illustrate the use of single stellar populations by predicting the colors of primordial proto-galaxies and show that one can first find them and then deduce the form of the IMF for the early generation of stars in the universe. Finally, we provide accurate analytic fitting formulas for ultra fast computation of colors of single stellar populations. ",
        "introduction": "The synthetic stellar spectrum of a galaxy is a well established theoretical tool for investigating the properties of the integrated light from distant galaxies where individual stars cannot be resolved. Since the late 60's several groups have developed different grids of synthetic stellar population models using a variety of stellar interior tracks and observed or theoretical stellar spectra (e.g. \\citet{Tinsley_68,BB77,Renzini_81,Bruzual_83,BO86,Arimoto_Yoshii_87,GR_87,BC93,Worthey_94,BCF94,FiocRocca97,JPMH98,LLG02}). The idea is very simple: stars are born with a given initial mass function (IMF) and they evolve in time according to stellar evolution. At any time one can compute the integrated spectrum by summing up the individual spectra of the stars in the population at that instant. Two major approaches have been used to compute the integrated light of a stellar population: the fuel consumption theorem \\citep{Renzini_81,Renzini_Buzzoni_83} and the isochrone technique, first developed by \\citet{BB77} and later used by \\citet{Charlot_Bruzual_91}. The fuel consumption theorem simply uses the fact that the contribution of stars in any post main--sequence evolutionary stage is proportional to the amount of nuclear fuel they burn at that stage, and approximates the post main--sequence evolution of the stars in the integrated population by the stellar track of the most massive star alive at the main sequence turn--off (MSTO). In contrast, the isochrone technique uses a continuous distribution of stellar masses, and hence tracks, to compute the locus in luminosity and $T_{\\rm eff}$ at a given time for any mass. From this, a smooth isochrone can be computed. The fuel consumption theorem remains an elegant method of studying the fastest stages in the evolution of any population. On the other hand, with the new generation of fast computers, the isochrone technique is clearly the most accurate and precise method of computing the integrated light of any stellar population.  Regardless of the computational technique used to calculate the integrated spectrum of a stellar population, the most important ingredient remains the stellar input: both stellar interior and stellar atmospheric models. Disagreement over stellar interior models (convection, mass loss, opacities), the modelling of post main sequence evolutionary stages and the modelling of stellar atmospheres (opacities, NLTE effects, bolometric correction, mass loss, etc.), combine to make the derived age of even a simple (i.e., no dust, no AGN contamination) stellar population vary by about 10\\% (for a fixed mass and metallicity), depending on which of the currently available synthetic stellar population codes (e.g. \\citet{Charlot_Worthey_Bressan_96,Spinrad+97}) is used to interpret the data (see section 4). A similar problem occurs when trying to date Globular Clusters in the Galaxy, the age of which is currently uncertain by about 10\\% \\citep{Chaboyer95,Jimenez+96,ChaboyerKrauss02}.  We have been motivated to attempt to improve this situation by the fact that the new generation of large 8-10m telescopes can now deliver spectra of galaxies at $z>1$ of sufficient quality to merit accurate age dating. The accurate determination of the ages of high-redshift galaxies can yield important constraints not only on models of galaxy formation, but also on the age of the universe. In particular, in a series of papers \\citep{D+96,Spinrad+97,NDJ01,NDJH03} we have addressed the issue of determining the ages of the reddest known elliptical galaxies at $z \\geq 1$. To aid in the interpretation of our data and others, we have developed a set of simple synthetic stellar population models (SSP) that overcomes some of the problems described above.  Most of the disagreement between the existing synthetic stellar evolution codes stems from the difficulty of modelling accurately the post main-sequence evolution, both because the physics involved in these stages is not completely known (opacities, convection, nuclear rates) and because mass loss strongly controls the final fate of the evolution of the star in these phases. Ideally, one would like to have a robust set of stellar models that are computed self-consistently (i.e. interior, photosphere and chromosphere computed at the same time) and that include the effects of mass loss and dust grain formation. However this is not yet possible, and in any case it is important to realise that mass loss {\\em cannot} be incorporated as a {\\em fixed quantity} for all stars in the population since it varies from star to star.  In the new synthetic stellar population models presented here we have endeavoured to improve the ability of the modelling to accurately reproduce the post-main sequence evolution of real stellar populations by incorporating an algorithm which has been previously applied with success to a number of other stellar evolution problems (e.g. \\citet{Jorgensen_91,JT93,Jimenez+96,Jorgensen_Jimenez_97}). This algorithm accurately simulates the evolution of all post main-sequence evolution stages, and includes a proper modelling of the HB along with an accurate account of the formation of carbon stars on the AGB. Furthermore, the mixing length parameter and the mass loss are properly calibrated using the position of the RGB and the morphology of the HB in real star clusters, respectively.  The other main new feature of our spectral synthesis models is the incorporation of new stellar evolutionary tracks, with sufficient resolution in mass to sample rapid stages of stellar evolution. As a result of these improvements (which we describe in detail in this paper), unlike several existing population codes, our models yield consistent ages when used to date a coeval population from a wide range of spectral features and colour indexes.  Our new SSP code has been applied successfully to a variety of different populations. it has been used to determine the ages of high redshift galaxies \\citep{D+96}, the ages of Low Surface Brightness Galaxies \\citep{PJA97,JPMH98}, the role of star formation and the Tully-Fisher law \\citep{HJ99} and the age of the Galactic disc \\citep{JFK98}. The purpose of this paper is to present the new library of synthetic stellar population spectra and discuss in more detail the physics and assumptions in our SSP modelling procedure.  The paper is organised as follows: in \\S 2 we present the new set of stellar interior tracks and discuss their accuracy when confronted with individual stellar observations. The synthesis models are presented in \\S 3. The degeneracies in the age-metallicity plane are discussed in \\S 4 while systematic errors are considered in \\S 5. The application of synthesis models to primordial proto--galaxies is presented in \\S 6 along with a method to determine the initial mass function of these galaxies. \\S 7 discusses the IR flux density and detectability of primordial proto--galaxies. Our conclusions are presented in \\S 8. An appendix provides fitting formulas for computing broad band colors of SSPs. ",
        "conclusions": "In this paper we have presented new stellar interior tracks and single stellar population models of arbitrary age, metallicity and IMF. Our main findings are:  \\begin{enumerate}  \\item A new set of stellar interior models has been presented. The models cover {\\em all} stages of stellar evolution and we have built isochrones out of them for ages between $10^6 - 1.4 \\times 10^{10}$.  \\item It is possible to construct stellar evolution models that accurately reproduce the properties of individual stars for a wide range of ages and metallicities.  \\item We have presented a new algorithm to compute the evolution of stars in the RGB, HB and AGB. This algorithm makes it possible to explore the effect of variations in some of unknown parameters of stellar evolution like mass loss, mixing length etc.  \\item We have developed a new and fast algorithm to build synthetic stellar evolution spectra and colour--magnitude diagrams of arbitrary metallicity and age.  \\item We have shown that changes in the values of the stellar parameters like mass loss and mixing length can change the predicted colours of a population by as much as 0.4 mag.  \\item We have studied degeneracies in the parameter space (age and metallicity) and shown that these parameters are only degenerated if the wavelength range of the spectrum is very small or only a few spectral features are chosen. Addition of light bluewards of the $4500$ \\AA\\, ($2500 - 4500$ \\AA) significantly reduces this degeneracies and, in fact, lifts them.  \\item It has been shown that systematic errors among different models are at the level of 10-20\\%, despite using completely different stellar input physics. It should be possible to reduce these errors even further.  \\item We have studied the photometric properties of very young proto--galaxies with primordial or very low ($Z=0.01 Z_{\\odot}$) metallicity and no significant effect of dust in their SED. We have named these galaxies ``primordial protogalaxies'', or PPGs. Using the methods of synthetic stellar populations, we predict that PPGs are the bluest stellar systems in the Universe. They can therefore be selected in color--color diagrams obtained with deep broad band IR surveys, and can be detected with the JWST, if they have a SFR of at least $100$~M$_{\\odot}/$yr, over a few million years.  We have discussed the possibility of using the IR colours of PPGs to constrain their stellar IMF, and investigated the possibility that the stellar IMF arising from gas of primordial chemical composition is more ``massive'' than the standard Salpeter IMF. Finally we have argued that broad band photometry can be more efficient than emission line searches, to detect and select PPGs.  \\end{enumerate}  The models are available on the world wide web ({\\tt www.roe.ac.uk/$\\sim$jsd} and {\\tt www.physics.upenn.edu/$\\sim$raulj}). We provide the stellar interior tracks presented in this paper and the single stellar population models and tools to compute photometry and synthetic stellar populations with arbitrary star formation histories."
    },
    "0402/astro-ph0402047_arXiv.txt": {
        "abstract": "We predict the amplitude  of the gravitational redshift of galaxies in galaxy clusters using an N-body simulation of  $\\rm{\\Lambda}$CDM universe. We examine if it might be possible to detect    the   gravitational   effect   on   the    total   redshift observed for galaxies. For  clusters  of mass  $M  \\sim10^{15}\\msun$,  the difference in  gravitational redshift  between the brightest  galaxy and the rest of the cluster members is $\\sim10 \\kms$. The most efficient way to detect gravitational redshifts using information from galaxies only involves using the full gravitational redshift profile of clusters. Massive clusters, while having the largest gravitational redshift suffer from large galaxy peculiar velocities and substructure, which act as a source of noise. This and their low number density make it more reasonable to try averaging over many clusters and groups of relatively low mass. We examine publicly available data for 107 rich clusters from the ESO Nearby Abell Clusters Survey (ENACS), finding no evidence for gravitational redshifts. Test on our simulated clusters show that we need at least $\\sim 2500$ clusters/groups with $M> 5 \\times 10^{13}\\msun$ for a detection of gravitational redshifts at the 2$\\sigma$ level. ",
        "introduction": "General  Relativity  predicts  redshift  of  photons due to a gravitational  field.  When  a  photon with  wavelength  $\\lambda$  is emitted in a gravitational potential  $\\Phi$, it will lose energy when it  climbs up in the  gravitational  field and  will consequently  be redshifted. The redshift observed at infinity is given in the weak field limit by: \\begin{equation} z_{g} = \\frac{\\Delta \\lambda}{\\lambda} \\simeq \\frac{\\Delta \\Phi}{c^{2}} \\label{eqn:zg} \\end{equation} where $\\Delta \\lambda$, $\\Delta \\Phi$ are respectively the difference in wavelength, and difference in potential between where the photon is emitted and where it is observed.  If we consider galaxies as sources of the photons, the gravitational redshift effect is so tiny that we take it for granted that  a measurement of the total galaxy redshift can be assumed to be the sum of Hubble expansion and peculiar velocities. In this paper we examine whether this is always the case, and in particular whether galaxies in galaxy clusters could have measurable values of $z_{g}$.   Since the gravitational  potential depends on the mass distribution around  galaxies,  the gravitational redshift, if observable, should be most evident in dense environments. In an early study by Nottale (1976), the redshift difference between pairs of clusters was compared to the richness difference. A supposed strong effect was found, with the pair member of higher richness having a systematically high redshift. However, when Rood \\& Struble (1982) rexamined this with a larger sample, their result showed no such correlation. Nottale (1990)  discussed that the effect should be looked for in galaxies at the centers of  galaxy clusters, by comparing their redshifts with those of galaxies at the cluster edges. Stiavelli \\& Setti (1993) carried out a related test in individual elliptical galaxies, finding at $99.9\\%$ confidence that elliptical galaxy cores are redshifted with respect to the galaxy outer regions, explaining this as a result of gravitational redshift.  The study of gravitational redshifts in galaxy clusters was taken further by Cappi (1995), who modelled clusters  using different density profiles including a de Vaucouleurs law. It was predicted that the gravitational redshift is non-negligible in very rich clusters. For example, the centers of clusters of masses $10^{16} \\msun$ should be redshifted by as much as $300 \\kms$ with respect to infinity. Broadhurst \\&  Scannapieco (2000) modelled the effect using a Navarro Frenk \\& White (1997) (hereafter NFW) density profile, and suggested that the gravitational redshift of metal lines in the cluster gas could eventually be used to map out the potential directly. As the gravitational redshift  is sensitive to the distribution of mass in the innermost regions of clusters, it could be  used as a  probe to constrain the amount of dark matter there. Gravitational redshifts would provide complimentary information to gravitational lensing (e.g., Sand \\etal 2003), as unlike lensing they do not depend on the mass density projected along the line of sight.  Here we use an $N$-body simulation made publically available by the Virgo Consortium (Frenk et al.  2000) to estimate the magnitude of the effect of gravitational redshifts on galaxy clusters in a $\\rm{\\Lambda CDM}$ universe. We examine possible observational strategies and determine if galaxy gravitational redshifts could be detected with a reasonable   number  of   clusters.  By  using   the  numerical simulation, we will be able to study the effect of substructure in the density and the  potentially complex velocity  field of  realistic clusters. We will see if Nottale's suggestion of measuring the difference in gravitational redshift between the  central galaxy  and galaxies at  the edge of  the cluster is realizable in practice. The potential wells should be deeper for the most massive clusters, which means larger gravitational redshifts. However, massive clusters are rare, so that there may not be enough large mass clusters in reality, something which we will investigate.  The  gravitational redshift  of the galaxy closest to the center of the potential is always positive with respect to the other galaxies, whereas the Hubble velocity component and peculiar velocity components can have either sign. We  extract the gravitational redshift information  by summing the total  redshift from galaxies in many  clusters  so that  the  noise from  Hubble expansion and peculiar velocities is averaged out. Small mass clusters may be better candidates than massive ones in  this sense because they are more abundant in  the Universe.  The  effective  mass range  for  detection is  also discussed in this paper.  The  structure of  the  paper is  as  follows. In  \\S2, the  numerical simulation used as our model universe is described. In \\S3, we discuss the magnitude of the gravitational  redshift and the number clusters needed in order to detect the  effect observationally. In \\S4, we briefly examine constraints on gravitational redshifts from an observed sample of cluster galaxies, take from the ENACS survey. Finally,  in \\S5, we summarize our results and conclude. ",
        "conclusions": "We have explored the gravitational redshifts of galaxies in the potential wells of simulated galaxy clusters. In particular we have concentrated on the possibility of detecting gravitational redshifts from surveys of cluster galaxies. Our main findings are as follows:  \\begin{itemize} \\item For clusters of  $\\sim10^{15}\\msun$, the difference in redshift between the central galaxy  and the  other cluster  members is  $\\sim 10 \\kms$. \\item The gravitational potential of clusters does not exhibit much substructure, even when the clusters are not relaxed. The most important source of noise on a measurement is from galaxy peculiar velocities. \\item The  most sensitive method for detecting gravitational redshifts with galaxy data  is to use the gravitational redshift profile as a function of impact parameter. \\item For a given number of galaxy redshifts measured, the most efficient  strategy would  be  to target  clusters  of $\\sim5  \\times 10^{13}\\msun$, which have lower noise from galaxy peculiar velocities than more massive clusters. \\item For a 2$\\sigma$ detection, we require $\\sim 5000$ clusters with  $M >  5 \\times 10^{13} \\msun$ when the measurement uncertainty is 30 km/s and $\\sim 2500$ clusters when the measurement uncertainty is negligible. \\end{itemize}  Our result for the shape of the gravitational  redshift profile are consistent with that of Cappi(1995) and Broadhurst and Scannapieco (200). However, the masses of clusters in our simulation are generally smaller than those used as examples by these authors. For example Cappi (1995) found that gravitational redshift with respect to infinity of the central regions of a cluster of mass $M  = 5 \\times 10^{15} \\msun$  to be $> 100 \\kms$ and  predicted the  effect  for even  more massive clusters  with $M  = 10^{16}    \\msun$ ($v_{g}  \\sim  300  \\kms$). Unfortunately, even though the gravitational redshift in a cluster with mass $10^{16} \\msun$ is large, there is little chance of finding such a massive cluster  in the Universe.  We can  roughly estimate the number of clusters  in the universe whose mass is $M = 8 \\times 10^{14}  \\msun$ or greater. First, we approximate the mass function of clusters given by (Bahcall \\& Bode, 2003) as a power law: \\begin{equation} \\log n (>M_{1.5,com}=8\\times10^{14} \\msun) \\approx -2 z - 6.75, \\label{eqn:bahcall} \\end{equation} The comoving distance as a function of redshift is given by (e.g., Hogg 1999): \\begin{equation} D = D_{H} \\int^{z}_{0} \\frac{dz'}{E(z')} \\label{eqn:hogg} \\end{equation} where $E(z) \\equiv \\sqrt{\\Omega_{M} (1+z)^{3} + \\Omega_{k} (1+z)^{2} + \\Omega_{\\Lambda}}$  and  $D_{H} =  3000  \\hmpc$. We take $\\Omega_{M} = 0.3$ and $\\Omega_{\\Lambda} = 0.7.$ We then find the total number of  clusters in the universe (up to $z=1$, although raising the $z$ limit makes essentially no difference) : \\begin{equation} N = \\int n dV = 4 \\pi D_{H}^{3} \\times 10^{-6.75} \\int^{1}_{0} \\frac{10^{-2 z}}{E(z)} \\left( \\int^{z}_{0} \\frac{dz'}{E(z')} \\right) ^{2} dz. \\label{eqn:totn} \\end{equation} Integrating Eq.~\\ref{eqn:totn} numerically  we find that there are only $\\sim 600$ clusters ($M >  8 \\times 10^{14} \\msun$) observable in the entire universe. We have seen, however in \\S3, that in any case the larger noise from peculiar velocities in more massive clusters makes a detection of gravitational redshifts problematic anyway. It might  be  better to  use lower  mass clusters in  order to detect  the gravitational redshift, even though effect is small.   As for observed samples of clusters with lower masses, the situation is  more promising. For example, the Sloan Digital Sky Survey has been used to find $\\sim  1,000$ clusters in total with $z < 0.15$ and  $M> 5  \\times 10^{13} \\msun$ so far (Chris Miller, 2003, private communication). For a more on cluster finding in the SDSS, see Nichol (2003) and Annis \\etal (2003). Published samples of SDSS clusters include those of Goto \\etal (2002) and Bahcall \\etal (2003b). Merchan and Zandivarez (2002) have constructed a galaxy group catalog from the 100K data release of the the 2dF  galaxy  redshift survey (refs). They find  $\\sim  1000$ clusters/groups  with masses  $M> 5  \\times 10^{13}  \\msun$. The final 2dF data release contains $\\sim  2 \\times  10^{5}$ galaxies  (Colless et al. 2003), which should enable construction of a group/cluster catalog with $\\sim 2500$ members above this lower mass limit. As we have seen in \\S 3.4, this may be close to the level required in order to make a $2 \\sigma$ detection of the gravitational redshift. A catalog of similar groups selected from the final SDSS survey might allow a 4 $\\sigma$ detection.   Because the effect we are searching for is very small (a redshift on the order of a few km/s), we must be careful to consider possible systematic effects.  The effects of peculiar velocities and cluster substructure have been included in our simulations, and averaging over large numbers of clusters produces reasonable results. One effect we have not included is full selection of clusters in redshift space. In the simulation, we  simply define a cluster in real space with  a linking length of our choice. In  a survey, we select the clusters in redshift space and we  include  some contaminating galaxies  behind and in front of each cluster that are not members. This would act to increase the noise, but not by much, as they will be heavily outnumbered by cluster members.  A real problem would be any effect which gives a systematically positive or negative redshift for the central galaxy with respect to the other galaxies. If we select cluster galaxies using a cone of constant angular size rather than a cylinder of fixed radius, then more contaminating galaxies will come from the larger volume behind the cluster than in front, and the average redshift of the galaxies apart from the central one will change. We have calculated the potential effect of this bias using the cluster-galaxy cross-correlation function data of Croft \\etal (1999) to find the number of contaminating galaxies. We find a relative blueshift of $\\sim 1\\kms$ for the central galaxy. This can be avoided simply by using a cylinder to select cluster members, however.   An additional potential problem which goes in the opposite direction is the fact that with a magnitude limited survey, the contaminating galaxies behind the cluster should be fewer in number because they are at a greater distance. The size of this effect depends on the steepness of the luminosity function of galaxies. It could also however easily be eliminated by applying a local volume limit in the vicinity of each cluster, rejecting galaxies which would be too faint when placed at the far edge of the cluster.  In this paper we have seen that realistically, gravitational redshifts are an extremely small effect, and that looking at the most massive galaxy clusters may not help in detection because the noise is much larger than for small clusters. If a detection is to be made using a galaxy survey, it will be important to carry out cross checks. For example, enough data needs to be a available that clusters can be divided by mass into more than one bin, to make sure that the effect is largest for the high mass clusters (even if the error bar is large). The more tracers of the clusters potential that are available, the better. Broadhurst \\& Scannapieco (2000) have shown that X-ray emission from intracluster gas could be used. Other probes of the potential could include using redshifts of intracluster planetary nebulae (e.g., Feldmeier \\etal 2003).  It may be marginally possible to  detect the gravitational redshift of cluster galaxies with the data available today. However,  more surveys of greater scope have been imagined for the future. Someday this will provide  us enough information to separate the gravitational redshift effect from total redshift.  We thank the anonymous referee for useful suggestions. RACC acknowledges support from the NASA-LTSA program, contract NAG5-11634."
    },
    "0402/astro-ph0402470_arXiv.txt": {
        "abstract": "{ We present moderate resolution (R $\\geq$ 1,800) infrared {\\it K}-band spectra of twelve long-period (P\\raisebox{-.6ex}{orb} $\\geq$ 6 hr) cataclysmic variables. We detect absorption lines from the photospheres of the secondary stars in every system, even though two of them were undergoing outbursts. We have attempted to assign spectral types to each of the secondary stars, and these classifications are generally consistent with previous determinations/estimates. We find evidence for abundance anomalies that include enhancements and/or deficits for all of the species commonly found in {\\it K}-band spectra of G- and K-type dwarfs. There is, however, only one common abundance anomaly: extremely weak CO features. Only two of the twelve objects appeared to have normal levels of CO absorption. We interpret this as evidence for low carbon abundances. In addition, we detect \\raisebox{.6ex}{13}CO absorption in four of the twelve objects. Depleted levels of \\raisebox{.6ex}{12}C and enhanced levels of \\raisebox{.6ex}{13}C indicate that material that has been processed in the CNO-cycle is finding its way into the photospheres of CV secondary stars. In systems with luminous accretion disks, we find that the spectrum of the secondary star is contaminated by a source that flattens (reddens) the continuum. While free-free or classical accretion disk spectra are flatter than the blackbody-like spectra of G and K dwarfs, removal of such contamination from the {\\it K}-band data results in spectra in which the absorption features become too strong to be consistent with those of G and K dwarfs.} ",
        "introduction": "Cataclysmic variables (CVs) are short-period binary systems consisting of a white dwarf primary that is accreting material via Roche-lobe overflow from a low mass, late-type secondary star. The commonly proposed evolutionary history for cataclysmic variables (CVs) establishes that the vast majority of CV secondary stars have undergone very little evolution during their lifetime (see Howell, Nelson, and Rappaport 2001, and references therein). The formation of a CV from a wide binary containing two main sequence stars is envisaged to have three main phases: First, the orbital separation of the wide binary of the pre-CV is rapidly shrunk in a common envelope phase where the secondary star orbits inside the red giant photosphere of the white dwarf progenitor. The second phase is a very long epoch where gravitational radiation, or a magnetically constrained wind from the secondary star (magnetic braking) extracts angular momentum from the binary, resulting in the eventual contact of the photosphere of the secondary star with its Roche lobe. The final phase begins once the secondary star contacts its Roche lobe, mass transfer to the white dwarf is initiated, and all of the phenomena associated with CVs is observed. During the lifetime of the mass transfer phase, the overall mass of the secondary star is gradually reduced. Much of the material accreted by the white dwarf is believed to be lost from the typical CV system through numerous classical novae eruptions. In both the common envelope phases and during classical novae eruptions, material with a peculiar composition can be deposited in the photospheres of CV secondary stars.  Evidence for the existence of peculiar abundance patterns in CVs is growing. For example, using UV spectroscopy, Cheng et al. (1997) found that the carbon abundance was 5$\\times$ solar, nitrogen was 3$\\times$ solar, and silicon was $\\leq$ 0.1$\\times$ solar for the white dwarf in WZ Sge. They suggested that this material was probably transferred from the secondary star. For VW Hyi, Sion et al. (1997) found that the white dwarf appeared to be deficient in carbon (0.5$\\times$ solar), iron (0.5$\\times$ solar), and silicon (0.1$\\times$ solar), but had an excess of nitrogen (5$\\times$ solar), oxygen (2$\\times$ solar), and phosphorous (900$\\times$ solar). Sion et al. (2001) found subsolar abundances for carbon (0.05$\\times$ solar) and silicon (0.1$\\times$ solar) for the white dwarf in RX And. In both CN Ori (Urban et al. 2000) and AH Her (Lyons et al. 2001) subsolar silicon abundances were found. Meanwhile, Sion et al. (1998) and Long \\& Gilliland (1999), estimated that the carbon abundance of the white dwarf in U Gem is about 0.1$\\times$ solar, while the nitrogen abundance is about 4$\\times$ solar. Harrison et al. (2000) found from infrared spectroscopy that the secondary star in U Gem appeared to have extremely weak CO features, suggesting it was deficient in either carbon, or oxygen. If so, the deficit of carbon on the white dwarf for U Gem could be easily explained by the transfer of carbon-poor material from the secondary star. In addition, Harrison et al. found that the secondary star of SS Cyg displayed weaker CO features than it should for its spectral type, along with an apparent magnesium deficit. Mennickent \\& Diaz (2002) report weak CO absorption for the late-type secondary in VY Aqr. It appears that the both the primary and secondary stars in CVs have peculiar compositions.  Recently, G\\\"ansicke et al. (2003) have reported on anomalous N V/C IV line flux ratios in new {\\it HST} STIS ultraviolet spectroscopy for several CVs, and compile a list of ten CVs that all show similar spectra. They conclude that these represent true abundance anomalies, and that nitrogen is strongly enhanced relative to carbon.  What could be the origin of such abundance anomalies? Sion (1999) suggests the abundance anomalies in the white dwarf photospheres may result from the nuclear processing in classical novae explosions. As shown by Jos\\'e, Coc and Hernanz (2001), however, only classical novae with very massive white dwarfs (1.35 M$_{\\sun}$) can produce nuclei such as silicon or phosphorous in thermonuclear burning (and such burning increases, not decreases the abundance of silicon). White dwarf masses in typical CV systems cluster near 0.6 M$_{\\sun}$. The presence of similar abundance patterns in the secondary stars suggest that for most CVs, they are a more likely source for material of peculiar composition. Marks \\& Sarna (1998) performed a detailed theoretical study of the possible effects on the surface abundances of the secondary star due to both evolutionary effects in the secondary itself, as well as that due to sweeping-up of CE material or matter accreted from classical novae ejecta. All such events could place thermonuclear processed material into the photosphere of the secondary star. Considering the first case, Marks \\& Sarna find that the photospheric chemistry of the secondary star could show large abundance variations in carbon, nitrogen and oxygen from evolutionary effects alone. In this scenario, the CNO tricycle is operating in the secondary star either before or during the contact phase. As material is removed from the secondary star, layers where the CNO tricycle was operating are exposed, or mixed to the surface, creating abundance and isotopic variations in CNO species. Especially relevant is their predictions for an overall deficit in carbon, enhancements in nitrogen, and a dramatic change in the ratio of \\raisebox{.6ex}{12}C/\\raisebox{.6ex}{13}C. For this work, however, Marks \\& Sarna only considered initially massive secondary stars (1.0 - 1.5 M$_{\\sun}$), whereas Howell et al. (2001) have shown that massive secondaries are likely to be present in only a small fraction of CVs.  Marks \\& Sarna (1998) also performed a study of the effects on the surface abundances of the secondary star due to sweeping-up of CE material or by accreting classical novae ejecta. They concluded that any material acquired during the CE phase would be thoroughly mixed into the secondary star during the extended period between the CE phase and the time the secondary contacted its Roche lobe. Marks \\& Sarna did find that if the process of accreting novae ejecta was efficient, dramatic abundance and isotopic anomalies could be present in the photospheres of CV secondaries.  G\\\"ansicke et al. (2003) suggest that the anomalous nitrogen to carbon abundances are a natural consequence of a scenario where the initial mass of the current donor star was greater than that of the white dwarf. As described by Schenker et al. (2002), in this situation a short-lived phase of very high, and dynamically unstable mass transfer quickly whittles away the outer layers of the donor star, leading to the production of a CV with a relatively normal mass ratio, but one where the donor star is now the CNO processed core of the massive donor. Schenker et al. propose that the unusual CV system AE Aqr has just completed this phase of evolution. As shown in their Fig. 9, the surface chemical abundance ratios of \\raisebox{.6ex}{12}C/\\raisebox{.6ex}{13}C, and C/N drop by one or two orders of magnitude as the system like AE Aqr evolves to become a CV. This scenario predicts depleted levels of carbon, enhanced levels of nitrogen and \\raisebox{.6ex}{13}C, and that the donor stars will have spectral types that are too late for their orbital periods.  Thus, the detection and measurement of abundance anomalies in a secondary star may provide direct insight into the evolutionary history of a CV. Of course, it remains quite possible that any observed abundance anomalies might arise due to unusual excitation conditions within the non-equilibrium photospheres of irradiated, mass-losing secondary stars. We present new infrared spectra of a dozen long period cataclysmic variables to search for additional abundance anomalies in their secondary stars. We detect the secondary star in every CV, and find evidence for carbon deficits in nearly all of them. In addition, we detect \\raisebox{.6ex}{13}CO for the first time in a CV secondary star. We find that in the case of MU Cen, the strength of the \\raisebox{.6ex}{13}CO feature suggests that the CNO cycle has run to completion (\\raisebox{.6ex}{12}CO/\\raisebox{.6ex}{13}CO = 3.2). We also find a wide range in the strengths of lines from such elements as sodium, calcium, magnesium, silicon, and iron when compared to main sequence stars of the {\\it most appropriate} spectral type. The origin of such a wide range of behavior is not easily identifiable. The detection of enhanced \\raisebox{.6ex}{13}CO certainly suggests that CNO processed material has made it into the atmospheres of a number of CV secondary stars, but whether this is from the accretion of material, or due to the evolutionary history of the secondary star itself, remains unclear. Further high resolution, high S/N spectroscopic observations will be needed to quantify these anomalies. Equally important, however, will be the need for good atmosphere models to help rule out any effects due to peculiar excitation conditions.  In the next section we discuss our observations, followed by a description of the spectra of the objects in section 3, followed by our conclusions in section 4. ",
        "conclusions": "There are two consistent results from our {\\it K}-band survey of long period cataclysmic variables. The first is that just about all of the CV systems studied here exhibit weaker-than-expected CO absorption features for their apparent spectral types. The other common trend is for systems with luminous accretion disks to have redder continuua than expected. There are two plausible explanations for the weakness of the CO features: either carbon and/or oxygen are deficient in the secondary stars of these systems, or weak CO emission is occurring from elsewhere in the system that fills-in the absorption features of the secondary star.  Howell et al. (2004) have recently used Keck to obtain high S/N observations of WZ Sge which reveal both CO and molecular hydrogen (at 2.22 $\\mu$m) in emission. In the case of WZ Sge, the H\\raisebox{-.6ex}{2} emission was about one third the strength of the emission from the first overtone bandhead of CO. There is no evidence for H\\raisebox{-.6ex}{2} emission in any of the objects discussed here. Given this, and the evidence for carbon deficits seen in the UV spectra of some of the white dwarfs in CVs (e.g., U Gem), it seems more likely that the carbon is underabundant in CV secondary stars. Unfortunately, because the CVs in our sample are all long period systems with G and K-dwarf secondaries, the water vapor features seen in the spectra of the M-type stars found in short-period CVs are not present, and it will be difficult to conclusively rule out an oxygen deficit for these objects using {\\it K}-band spectra.  That the {\\it K}-band continua of many long period CVs appear flatter/redder than the spectra of G and K-type secondaries can be explained by having a contaminating source that has a spectrum that is less steep than the F$_{\\lambda} \\propto \\lambda^{-4}$, blackbody-like spectra of G and K-type stars. Both the standard accretion disk spectrum (F$_{\\lambda} \\propto \\lambda^{-7/3}$) and free-free emission (F$_{\\lambda} \\propto \\lambda^{-2}$) are possible sources. To demonstrate what happens when we remove a contaminating source, we present the spectrum of V426 Oph in Fig. 18 from which a flat continuum source (F$_{\\lambda}$ = constant), with 44\\% of its {\\it K}-band flux, has been subtracted. While the resulting spectrum has the identical slope to the K5V template, the absorption lines of sodium and calcium are now too strong for a dwarf star! To achieve the same result with an free-free or accretion disk spectrum requires them to constitute an even larger fraction of the {\\it K}-band flux, with the result of even stronger absorption lines. Only the subtraction of spectra with a positive spectral index can minimize this effect. We are unaware of any physical process that generate such spectra, but observations using SIRTF would help quantify the nature of this emission.  \\subsection{$^{13}$CO and Other Apparent Abundance Anomalies} The goal of our program was to investigate whether the secondary stars in long period CVs show evidence for peculiar abundances. We summarize our results in Table 2, where a ``+\" indicates a possible enhancement, and a ``-\" a deficit (a ``?\" indicates uncertain, while an ``!\" indicates a significant deficit). Ellipsis indicate the spectrum was either too poor, or too contaminated, to have confidence in statements about the abundance of a particular element. A ``0\" indicates that a species seems to be at relatively normal level. While every single object appears to have something peculiar about its spectrum, the {\\it K}-band data alone are not quite sufficient to determine the source of these peculiarities. Only for SS Cyg, RU Peg, and V426 Oph, were the S/N of the spectra sufficiently high to confidently examine them for low-level enhancement/deficits.  Except for the near-universal weakness of CO, there is no apparent pattern in the strength of the lines from any particular element.  This indicates to us that normal cosmic dispersion might be mostly responsible for the observed abundance peculiarities. Given the complex environment in which the spectrum of a CV is emitted, unusual line strengths could also arise due to other effects. One feature, however, does stand out: the detection of \\raisebox{.6ex}{13}CO in four of the systems (MU Cen, EM Cyg, SS Cyg, and AH Her). In Fig. 19, we show close-up views of the spectra of these four objects in the region around the \\raisebox{.6ex}{13}CO\\raisebox{-.6ex}{(2,0)} bandhead. In the normal main-sequence counterparts of the (mostly) K-type stars found in this sample of long-period CVs, any \\raisebox{.6ex}{13}CO absorption is almost undetectable. But we find evidence for fairly strong features in these four CVs. Even in the objects where the S/N of the spectra are quite low (e.g., CH UMa), there is evidence for \\raisebox{.6ex}{13}CO features. {\\it Only for RU Peg and V426 Oph can we rule out enhanced levels of \\raisebox{.6ex}{13}CO.} In the models of Marks \\& Sarna (1998), the presence of \\raisebox{.6ex}{13}C was strong evidence for evolution of the secondary star, resulting from the baring of, or mixing-in material from, layers where the CNO cycle had been operating. In such cases, the isotopic ratios of carbon, nitrogen and oxygen in the photospheres of the secondary star were found to reach unusually large values. In this process, the common isotopes (\\raisebox{.6ex}{12}C,\\raisebox{.6ex}{ 14}N, and \\raisebox{.6ex}{16}O) of these three species become depleted. Given that we simultaneously observe both an apparent \\raisebox{.6ex}{12}CO depletion {\\it and} \\raisebox{.6ex}{13}CO enhancement suggests that material enriched from the CNO cycle is reaching the photospheres of CV secondary stars.  As described earlier, there are at least four presently envisioned paths that can provide CNO isotopic enrichment: 1) the accretion of CNO processed material during the common envelope phase, 2) the accretion of novae ejecta, 3) the possibility that the donor star originally had a higher mass than the white dwarf primary, and that it is now the CNO-enriched core of this massive star, and 4) the secondary stars to have begun to evolve off of the main sequence before becoming a CV. The first two paths are consistent with the current evolutionary paradigm for CVs that requires the secondary stars have undergone little evolution during their lifetimes as short period binaries. The scenario where the secondary stars suffers large CNO enrichment from the accretion of novae ejecta seems difficult to sustain, since the time between classical novae explosions is large ($\\sim$ 10\\raisebox{.6ex}{4} yr) and the amount of material that can be accreted is relatively small ($\\ll$ 10\\raisebox{.6ex}{-4} M$_{\\sun}$). The small amount of ejecta that could be realistically accreted would get mixed-in to the secondary star quickly enough to be become virtually undetectable. A similar scenario is envisaged for the common envelope phase. The time interval between the common envelope phase and the time of first contact of the secondary star with its Roche lobe is believed to be so long ($>$ 10\\raisebox{.6ex}{8} yr), that any accreted material would get thoroughly mixed-into deeper layers within the secondary star.  The apparent detection of enhanced levels of \\raisebox{.6ex}{13}CO indicates that CNO processed material is present in the atmospheres secondary stars of these long period systems, and that they have either undergone some evolution off of the main sequence, or they are the stripped cores of more massive stars.  The main difficulty with either of these two scenarios is that population synthesis models by Howell et al. (2001) find that very few CVs are formed with high mass secondary stars. It is interesting to note that we {\\it do} detect at least one relatively normal CV secondary star (SY Cnc) that has a spectral type similar to the sun. The possibility there are some CV secondary stars with initial masses of $>$ 1 M$_{\\sun}$, does not seem too far-fetched.  \\subsection{The Case for Subgiant Secondary Stars in SS Cyg and RU Peg}  In Harrison et al. (2000), a investigation into the nature of the secondary star in SS Cyg was made by combining a high-precision {\\it HST Fine Guidance Sensor} (FGS) parallax with ground-based photometry. They found that if all of the infrared luminosity of SS Cyg was presumed to be coming from the secondary star, then the secondary star must be a K4 subgiant. Recently, Harrison et al. (2004) have reported an {\\it HST} FGS parallax for RU Peg ($\\pi$ = 3.55 mas). Given that the spectrum of the secondary star in RU Peg appears uncontaminated by accretion disk emission, we thought it interesting to compare it to SS Cyg. In Fig. 20, we present the SEDs of SS Cyg (data from Harrison et al. 2000) and RU Peg (data from Harrison et al. 2004). The SEDs of the two systems are very similar. Their observed K magnitudes differ by $\\Delta$K = 1.12 mag, while their distance moduli differ by 1.16 mag! Thus, these two systems have nearly identical {\\it K}-band luminosities: M\\raisebox{-.6ex}{K} = 3.26 (RU Peg), and M\\raisebox{-.6ex}{K }= 3.30 (SS Cyg). These values should be compared to a K2V (the derived spectral type for RU Peg), which has an absolute magnitude of M\\raisebox{-.6ex}{K }= 4.15, while a K4V (the spectral type for SS Cyg) has M\\raisebox{-.6ex}{K} = 4.48. Thus, if the entire {\\it K}-band luminosities are ascribed to their secondary stars, RU Peg is 0.89 mags above the main sequence, and SS Cyg is 1.18 mags above the main sequence. It is difficult to envision a scenario where accretion disk contamination could supply $\\approx$ 1 mag of luminosity, yet not severely contaminate the secondary star spectrum. It therefore seems quite likely that both RU Peg and SS Cyg have secondary stars that have evolved off of the main sequence."
    },
    "0402/astro-ph0402193_arXiv.txt": {
        "abstract": "Using the NIRSPEC spectrograph at Keck II, we have obtained infrared echelle spectra covering the range $1.5-1.8~\\mu \\rm m$ for the highly reddened bulge globular clusters Terzan~4 and Terzan~5. The clusters are of interest because a blue horizontal branch in Terzan~4 is consistent with very low metallicity, while Terzan~5 has been proposed as possibly the most metal rich known Galactic globular cluster. We report the first detailed abundance analysis for stars in these clusters, and we find [Fe/H]=--1.60 and --0.21~dex for Terzan~4 and Terzan~5, respectively, confirming the presence of a large metallicity spread in the bulge as well as in the halo of the Galaxy. We find $\\rm [\\alpha/Fe]\\approx+0.5$~dex in Terzan~4 and $\\rm [\\alpha/Fe]\\approx +0.3$~dex  in Terzan~5, consistent with what has been found for field stars in the bulge.  The enhanced $\\alpha-$element abundances are also consistent with rapid chemical enrichment, most likely by type~II SNe. Finally, we measure very low  $^{12}$C/$^{13}$C$\\le10$ ratios in all the giant stars, confirming that extra-mixing mechanisms due to {\\it cool bottom processing} occur during evolution along the RGB. ",
        "introduction": "The detailed study of the abundance patterns of iron-peak, CNO and $\\alpha$ elements in field and globular cluster stellar populations is a fundamental step required to understand the star formation and chemical evolution history of our Galaxy. These elements are synthesized in stars of different masses, hence on different timescales and play a crucial role in disentangling evolutionary mixing effects from primordial enrichment (Wheeler, Sneden \\& Truran 1989, McWilliam 1997)  Globular clusters represent the best astrophysical templates for the study of the early Galactic evolution since they are the oldest objects and their age is well constrained, allowing a reliable calibration of a possible age--metallicity relation (McWilliam 1997), a fundamental constraint for any chemical evolution model. In addition, they are equally present in the halo, the thick disk and the bulge of our Galaxy (see e.g. Armandroff 1989, Minniti 1995), therefore they represent unique chemical fossils to trace the early epoch of the Galaxy formation.  Reliable elemental abundances are now available for field stars (see e.g.  Boesgaard et al. 1999; Gratton et al. 2000, Carretta, Gratton \\& Sneden 2000 and reference therein) while are still very incomplete for globular clusters. In halo/disk field stars the average [$\\alpha$/Fe] abundance ratio shows a general enhancement of 0.3--0.5 dex with respect to the Solar value up to [Fe/H]$\\approx$--1 and a linear decreasing trend towards Solar [$\\alpha$/Fe] with further increasing metallicity. The actual position of the knee (i.e. the metallicity at which [$\\alpha$/Fe] begins to decrease) depends on the type~Ia SN timescales and it is also a function of the Star Formation (SF) rate, while the amount of $\\alpha $ enhancement depends on the initial mass function of the progenitors of the type~II SNe (see McWilliam 1997).  While the classic halo globular clusters have a mean [$\\alpha$/Fe]$\\approx$+0.4, too few measurements have been done at the high metallicity end to define the trend there (e.g. Kraft 1994, Carney 1996). Bulge globular clusters are ideal targets to study the behavior of the abundance patterns in the high metallicity domain, but foreground extinction is so great as to largely preclude optical studies of any kind, particularly at high spectral resolution.  The most accurate abundance determinations in the Galactic bulge obtained so far and based on high resolution optical spectroscopy refer to a sample of K giants in Baade's window (McWilliam \\& Rich 1994, hereafter MWR94; Rich \\& McWilliam 2000) and a few giants in two globular clusters, NGC~6553 and NGC~6528 (Barbuy et al. 1999; Cohen et al. 1999; Carretta et al. 2001).  With the availability of NIRSPEC, a high throughput infrared (IR) echelle spectrograph at the Keck Observatory (McLean et al. 1998), came the prospect of measuring the composition of the clusters in the bulge. Recently, we observed two bright giants in Liller~1 and NGC~6553 with NIRSPEC and our abundance analysis has been presented in Origlia, Rich \\& Castro (2002, hereafter ORC02). We find $\\rm [Fe/H]=-0.3\\pm0.2$ and $\\rm [O/Fe]=+0.3\\pm 0.1$ We also measure strong lines for the other $\\alpha-$elements Mg, Ca, and Si, obtaining an overall $\\rm [\\alpha/Fe]=+0.3\\pm0.2$ dex.  In this paper we present the high resolution IR observations and the abundance analysis for a sample of bright giants in two additional globular clusters, namely Terzan~4 and Terzan~5.   These clusters are of special interest because they span the entire abundance range present in the bulge. These clusters suffer from large foreground extinction (E(B--V)$\\ge$2), precluding any optical study at high spectral resolution. Recent HST--NICMOS photometry (Ortolani et al. 2001) suggests old ages for these clusters and confirms previous photometric estimates for their metallicities (Ortolani, Barbuy \\& Bica 1996, 1997): roughly Solar for Terzan~5 and as low as 1/100 Solar for Terzan~4. Ortolani et al. (1997) argue that a blue horizontal branch is present in Terzan~4,  however the small numbers of stars in their color-magnitude diagram (CMD) requires an independent metallicity measurement in order for the case to be more convincing. More recent NICMOS photometry of Terzan~5 (Cohn et al. 2002) confirms the presence of an old turnoff point and a CMD consistent with high metallicity.  The recent photometry confirms that both clusters lie within $\\sim 1$ kpc of the Galactic center.  Our observations and data reduction follow in Sect.~2. Sect.~3 discusses our abundance analysis and in Sect.~4 the resulting metallicities and radial velocities are presented. Our concluding remarks are given in Sect.~5. ",
        "conclusions": "Our high resolution spectroscopy confirms the photometric metallicities of Terzan~4 and Terzan~5 as obtained from the morphology of the RGB (Ortolani et al. 2001).  These two bulge globular clusters were thought to lie at the extremes of the bulge abundance distribution; our detailed abundance analysis indicates that this is indeed the case. Our findings strengthen the analogy between the properties of the bulge globular clusters and the stellar field population in the bulge and remain consistent with a common formation history.  In addition, the detailed study of the abundance patterns indicates oxygen and other $\\alpha-$element enhancement at the level of $\\approx+0.5$~dex in the metal--poor cluster, typical of the halo population, and of $\\approx +0.3$~dex in the metal rich one. The enhanced $\\alpha-$element abundances for Terzan~5 are noteworthy, as enhanced alphas at Solar metallicity is a hallmark of bulge stars, reinforcing the scenario of a rapid enrichment of the bulge, similar to the halo but possibly requiring a higher rate of star formation (see e.g. Matteucci, Romano \\& Molaro 1999; Wyse 2000).  The direct measurement of the $\\rm ^{12}C/^{13}C$ abundance ratio provides major clues to the efficiency of the mixing processes in the stellar interiors during the evolution along the RGB. Both the metal-poor and metal-rich giants in the bulge clusters show very low  $^{12}$C/$^{13}$C$\\le10$ ratios. The classical theory (Iben 1967; Charbonnel 1994 and reference therein) predicts a decrease of $^{12}$C/$^{13}$C$\\approx40$ after the first dredge--up, the exact amount mainly depending on the chemical composition and the extent of the convective zone. However, much lower $^{12}$C/$^{13}$C$\\le$10 values have been measured in several metal-poor halo giants both in the field and in globular clusters (see e.g. Suntzeff \\& Smith 1991, Shetrone 1996, Gratton et al. 2000 and references therein). Also the bright giants of $\\omega$~Centauri show very low $^{12}$C/$^{13}$C$\\le$7 over the whole metallicity range spanned by the multiple stellar population in the cluster (see Brown \\& Wallerstein 1993; Zucker, Wallerstein \\& Brown 1996; Wallerstein \\& Gonzalez 1996; Vanture, Wallerstein \\& Suntzeff 2002; Smith et al. 2002), including the most metal-rich RGB stars (see Origlia et al. 2003). Additional mixing mechanisms due to further {\\it cool bottom processing} (see e.g. Charbonnel 1995; Denissenkov \\& Weiss 1996; Cavallo, Sweigart \\& Bell 1998; Boothroyd \\& Sackmann 1999; Weiss, Denissenkov \\& Charbonnel 2000) are thus a common feature during the evolution along the RGB, regardless the stellar metallicity and environment."
    },
    "0402/astro-ph0402646_arXiv.txt": {
        "abstract": "We discuss the fluxes of high energy neutrinos and gamma-rays expected from AGNs if hadrons can be effectively accelerated to ultra-high energies by their relativistic jets, as currently believed.  Fluxes of multi-TeV neutrinos detectable by $km$-scale detectors like {\\it IceCube} could be expected from powerful blazars where strong accretion-disk radiation is present in the AGN cores. Gamma-ray fluxes of hadronic origin can be important for flares in the compact jets of these sources up to GeV energies, but they will be insignificant for BL Lac objects. Production of UHE neutral beams composed of neutrons and gamma-rays can drive straight collimated jets in the intergalactic medium on multi-kpc scales, which could be resolved by the Chandra X-ray observatory. While we do not expect any significant neutrino flux from these large-scale jets, we predict gamma-rays of synchrotron origin in the energy range from sub-GeV up to TeV energies, which would be detectable by GLAST and ground-based gamma-ray telescopes. ",
        "introduction": "AGN jets, along with GRBs, are powerful accelerators of relativistic particles, and the prime candidate sources of extra-galactic cosmic rays (CRs) with energy spectra extending beyond $10^{20}\\,\\rm eV$.  Detections of non-thermal flares from blazars in the X-ray and particularly in the $\\gamma$-ray domain in the last decade have convincingly demonstrated that the compact inner ({\\it sub-parsec scale}) jets of blazars do accelerate particles to very high energies \\cite{har99,wee00}, presumably due to relativistic shocks.  Although analyses of correlated X-ray and TeV gamma-ray flares in BL Lacs support leptonic models \\cite{mk97,cat97,pia98,kca02}, which imply efficient acceleration of relativistic electrons, the associated acceleration of hadrons is expected with at least the same efficiency as that of the leptons. It might be different in electron-positron pair jets where few hadrons would be present, but comparison of the radio lobe and inner jet powers indicates that jets are composed mainly of electrons and protons \\cite{cf93}, so that a nonthermal hadronic component is very plausible.  Observational evidence for acceleration of hadrons in the jets of AGN requires a target for interactions in the jets. In principle, ultrarelativistic protons accelerated in the inner jets up to energies $\\sim 10^{20}\\,\\rm eV$ could produce synchrotron emission detectable at TeV energies \\cite{aha00,mp01}, but this requires extremely strong magnetic fields $\\sim 20$-100 G in the sub-parsec scale jets.  All other observable consequences of hadron acceleration result from interactions with ambient material or photon fields.  These include neutrinos, detection of which would provide a direct proof for hadron acceleration, and the high energy electromagnetic radiation from the secondary leptons and $\\gamma$-rays. Such evidence could also be provided by production of collimated beams of ultra-high energy (UHE; $\\gg 10^{15}\\,\\rm eV$ for the discussion below) neutrons and gamma-rays formed in the same process of hadronic interactions, which could explain effective transport of energy from AGN cores to multi-kpc scales \\cite{ad03}. ",
        "conclusions": ""
    },
    "0402/astro-ph0402120_arXiv.txt": {
        "abstract": "We present results from our analysis of the Infrared Space Observatory (ISO) data on the L~1641-N outflow region located in the Orion~A molecular cloud.  The data were obtained with the  array camera (ISOCAM) using two broad-band filters, LW2 (6.7~\\micron), \\& LW3 (14~\\micron), and the narrow-band circular variable filter (CVF) which provides low resolution (R=$\\lambda/\\delta\\lambda\\sim40$) spectra in the 5--17~\\micron\\ region. \\par We detect a total of 34 sources in the $7.65^{\\prime}$~x~$8.40^{\\prime}$\\ region covered by the broad-band filters.  Four of these sources have no reported detection in previous studies of the region. The CVF data are available for only the central $3.2^{\\prime}$~x~$3.2^{\\prime}$\\ portion of the total region, providing spectra for the brightest 7 of the sources in that region. \\par We find that the source previously identified as the near-IR counter-part to the IRAS detected point-source (IRAS 05338-0624) is not the brightest source in the wavelength region of the IRAS 12~\\micron\\ filter.  We find instead that a nearby object (within the beam of IRAS and not detected at near-IR wavelengths) outshines all others sources in the area by a factor of $\\sim$2.  We submit that this source is likely to be the IRAS detected point source. \\par A comparison of the near-IR (\\jmh\\ \\vs\\hmk) and mid-IR (\\hmk\\ \\vs $\\isotwothree$) color-color plots shows that while at near-IR wavelengths only four (4) of the sources show evidence for emission above the values predicted from photospheric emission alone (hereafter referred to as excess emission), at least 85\\% of all sources show evidence for excess emission at mid-IR wavelengths.  This result supports similar conclusions from  $L$-band surveys. \\par The CVF spectra suggest a range of evolutionary status in the program stars ranging from embedded YSOs to young disks.  When combined with optical and near-IR age estimates, these results show active current star-formation in the region that has been on-going for at least 2 Myr. ",
        "introduction": "\\label{sec-intro} \\par Intensive observational and theoretical studies of young, optically obscured star-forming regions over the last four decades suggest that stars like our Sun begin their lives in the midst of cold and dense molecular clouds. During the earliest phases of their evolution, these young stars are embedded behind 50-200 magnitudes of visual extinction , and likely develop extended bipolar outflows and jets in the early phases of their evolution towards the main-sequence (for a review, see Lada \\& Lada 2002).   The Orion giant molecular cloud complex is the closest example of such a star-forming region that offers both massive and low-mass stellar population. \\par The L~1641-N region is located in the southern part of the Orion~A molecular cloud. Attention was especially drawn towards L~1641-N by IRAS with the detection of a bright point source (IRAS 05338-0624). Subsequently, \\citet{fukui} reported discovery of two well-separated outflow lobes surrounding the IRAS source.  Further optical and near-infrared imaging studies revealed an enhancement of stellar density near the IRAS source at near-IR wavelengths (Chen \\etal\\ 1996; Hodapp \\etal\\ 1993), located \"among the highest concentrations of HH [Herbig-Haro] objects known anywhere in the sky\"  (Reipurth \\etal 1998). \\par The presence of the HH outflows implies that the L~1641-N region is one of the youngest star-forming environment in the Galaxy.  Dynamic age estimates of the outflow sources presently observed in the Galaxy constrain outflow ages to less than $10^5$~years (Fukui 1989; Fukui \\etal 1993).  Studying these young, optically obscured environments at mid-infrared wavelengths offers several advantages: (i) stars in the earliest phases of evolution are heavily obscured and are frequently undetected at even near-IR wavelengths.  (ii) With imaging data alone, it is difficult to distinguish between actual photospheric emission and reprocessed emission (reflected light).  Thus, near-IR reflection nebulae are often confused with actual sources.  (iii) At mid-IR wavelengths, both the stellar photosphere and the colder star-forming environment are visible. \\par However, the mid-IR region is difficult to observe from the ground because of the presence of high background from the warm Earth atmosphere. The Infrared Space Observatory (ISO, Kessler \\etal\\ 1996) provided the first space-based observatory with the ability to carry out sensitive observations at mid-IR wavelengths.  The ISO observed spectra of the embedded Young Stellar Objects (YSOs) revealed, for the first time, the presence of silicate, CO, water (vapor and ice) absorption features in the photospheric continuum of the stars (see for example, Gibb \\etal~2000).  These features are thought to arise in the colder star-forming environment in which the stars are forming.  Thus, the mid-IR region provides information both on the star itself and about the physical conditions in which they form. \\par In this contribution we provide an analysis of ISO CAMera (ISOCAM, Cesarsky \\etal~1996) mid-infrared broad-band and low-resolution Circular Variable Filter (CVF) observations of the L~1641-N region.  Section~\\ref{sec-obs} describes these observations in detail.  Section~\\ref{sec-phot} describes our source detection, photometry and cross-correlation with other optical and near-IR sources.  We present  and discuss our results in Section~\\ref{sec-results}.  Finally, the conclusions are summarized in Section~\\ref{sec-summary}. ",
        "conclusions": "\\label{sec-results} \\par We detect a total of 34 sources in the region covered by the broad-band filters.  The spatial coverage and sensitivity of the narrow-band filter data allowed us to obtain low-resolution spectra of only 7 of these sources.  All but 4 of the sources are previously reported by the Two Micron All Sky Survey (2MASS) and/or  the near-IR survey of  \\citet{chen}.  Two of the sources (\\#10 and \\#18, in our numbering scheme) are only detected at mid-IR wavelengths: by ISOCAM and by \\citet{chen} in their M-band images.  These results are further discussed below.  \\subsection{The counterpart to IRAS 05338-0624} \\label{sec-irascounterpart} \\par \\citet{chen} suggest source \\#18, in our numbering scheme, as the likely near-IR counterpart to the IRAS detected point source based on their {\\it M}-band imaging survey.  Source \\#18 is within 2 arc-seconds of the IRAS published position and is the brightest source at {\\it M}-band in the \\citet{chen} survey.   \\citet{stanke} also detected a bright (0.4 Jy) point source at 10\\micron\\ coincident with the location of source \\#18.  However, \\citet{stanke} argue that IRAS 05338-0624 cannot be unambiguously associated with source \\#18 because a 2.7-mm dust continuum source (Chen \\etal\\ 1993, 1995) is also coincident with the position uncertainty ellipse of the IRAS source (Figure ~\\ref{fig-m2}, also see Figure 4 of Stanke \\etal.\\ 1997).  In the ISO image, source \\#10 outshines source \\#18 by approximately a factor of two.  Source \\#10 is outside the field-of-view of the \\citet{chen} {\\it M}-band survey and the \\citet{stanke} 10\\micron\\ image, and is $\\sim$1 arc-minute away from the IRAS source position.  There are no other comparably bright sources detected in the ISO images.   The narrow-band CVF spectra (Figure~\\ref{fig-cvfspec}) include the wavelength range of the shortest IRAS filter and are useful for obtaining a more appropriate flux comparison with IRAS measurements.  We used the ISO Spectral Analysis Tool (ISAP) and the IRAS filter transmission curves to simulate IRAS photometry for sources \\#18 and \\#10 using the CVF spectra.  We find source fluxes of 284 and 481 mJy (30\\% errors) for sources \\#18, and \\#10, respectively.  The IRAS published photometry for IRAS 05338-0624 is 481 mJy.  Despite the excellent agreement between the fluxes of IRAS 05338-0624 and source \\#10, we conclude that source \\#18 remains the most likely counter part to IRAS 05338-0624 because it is the brightest source that is consistent with the position uncertainty ellipse of the IRAS source.   We rule out source \\#10 because for it to be the IRAS counterpoint implies a positional error of $\\sim$1 arc-minute in the IRAS position.  We consider this to be unlikely given that IRAS positions have been shown to be reliable to better than 7 arc-seconds (1-sigma) based on statistical comparison between IRAS and the Smithsonian Astrophysical Observatory (SAO) catalog point sources (cf. Section VII of the IRAS Explanatory Supplement, Beichman \\etal 1988).  The 2.7-mm dust continuum source lacks a bright mid-IR counter part in the ISO image implying either that it is not a strong emitter at those wavelengths or that it has dimmed significantly in the decade between IRAS and ISO measurements.  However, we also consider it unlikely that IRAS would not have spatially resolved and detected source \\#10, which is the  brightest source in the ISO image and about 1 IRAS beam-width away from IRAS 05338-0624 (source \\#18).  If we rule out  positional error in IRAS coordinates, then this implies that source \\#10 must have brightened by about at least 200 mJy  (using point source detection limit of 0.3 Jy for IRAS) between the ISO and IRAS epochs.  \\par We also note that the proposed IRAS counterpart, source \\#18, shows a broader PSF than other point sources in the area.  This is likely due to the presence of extended emission or emission knots which are unresolved by ISOCAM and seen in reflected light in the K-band as noted by \\citet{chen}.  \\subsection{The nature of point-like sources} \\par Figure~\\ref{fig-nirclrs} shows the \\jmh\\ \\vs \\hmk, color-color relationship for the 30 sources for which near-IR photometry is available.  In constructing this relationship, we used the photometry from 2MASS database.   The solid lines in the figure show the colors of ordinary dwarf and giant stars from the compilation of \\citet{bessell}.  The dotted lines in Figure~\\ref{fig-nirclrs}  show the effect of extinction on the colors of the reddest giant and dwarf stars.  The distribution of the L1641-N stars on Figure~\\ref{fig-nirclrs} is consistent with those from deeply embedded young stellar clusters. And, four stars, \\#6, \\#11, \\#17, and \\#21 show evidence for the so-called \"excess' emission -- colors that are redder than predicted by simple foreground extinction.  This \"excess\" emission is though to arise from warm circumstellar material (see \\eg\\ Lada \\& Lada 2003).  The emission itself is from warm dust in the circumstellar material. \\par We investigated a similar color-color relationship by using the ISO photometry for which all 4 photometry measurements are available (24 stars).  The results are shown in Figure~\\ref{fig-midclrs}: \\jmk\\ \\vs $\\isotwothree$.  As in Figure~\\ref{fig-nirclrs}, the dotted line shown in Figure~\\ref{fig-midclrs} shows the effect of foreground extinction, but in this case, for a star with intrinsic $\\isotwothree$\\ color of 0 mag.  At mid-IR wavelengths only a handful of estimates are available for the intrinsic colors of ordinary or standard stars.  These estimates show intrinsic values that are either 0 or close to 0 mag.  Thus, we adopted 0 mag as the typical $\\isotwothree$\\ color.  While 4 stars show evidence for excess emission in Figure~\\ref{fig-nirclrs}, all but four, sources \\#5, \\#31, \\#22, and \\#34 show evidence for excess emission at mid-IR wavelengths. This can be explained simply as warm dust emits more at longer than near-IR wavelengths; thus it is easier to detect excess emission at mid-IR wavelengths than near-IR color-color plot.  A similar argument was presented for the Chamaeleon I population by \\citet{nordh} and \\citet{persi}.  Adding source \\#10 and \\#18 in this sample, we conclude that at least 85\\% of all sources detected in the ISO LW2 filter show evidence of excess emission at mid-IR wavelengths. \\par Given the presence of excess emission in most of the observed sources, we conclude that most sources are most likely physically associated.  A conclusion also reached by \\citet{chen} in their analysis, who labeled it as a \"stellar density enhancement\".  The ISO narrow-band CVF spectra for the seven sources detected in the CVF filter shed further light on the nature and evolutionary status of the stars in this grouping.  These ISO spectra are shown in Figure~\\ref{fig-cvfspec}.  Each panel of Figure~\\ref{fig-cvfspec} shows the spectrum from the source identified in the left-hand-side of the panel.  Some of the prominent features associated with young stars are also labeled. \\par Five of the 7 detected sources, \\#11, \\#13,  \\#10,  \\#18, and \\#19, show the broad feature at 9.5 \\micron, which is characteristic of silicate absorption.  Sources \\#10, and \\#19 further show features seen in the very youngest proto-stars such as W 33A (Gibb \\etal\\  2000).  These features are labeled in Figure~\\ref{fig-cvfspec} and have been attributed to ices (CO$_2$, H$_2$O, etc.) in the environment of the young embedded stars (Gibb \\etal\\ 2000).  Source \\#14 shows marginal detections at wavelengths consistent with H$_2$\\  emission.  Sources \\#11 and \\#15, which show evidence for excess emission at near-IR wavelengths (see Figure~\\ref{fig-nirclrs}), also show evidence for an emission feature at wavelengths consistent with the PAH emission at 11.2~\\micron.  We rule out silicate emission because its peak is near 10~\\micron.   PAH emission has been observed in some T-Tauri stars and roughly 50\\% of the Herbig Ae/Be stars in the sample of \\citet{meeus}.   PAH emission can signify the presence of tiny grains and (likely) UV radiation to excite the PAH.  This is consistent with the observed location of sources \\#11 and \\#15 on Figure~\\ref{fig-nirclrs}.  Source \\#19, the faintest object detected in the CVF spectra, has relatively poor signal-to-noise ratio and most narrow features seen in its spectra are likely instrumental artifacts.  The only significant detection in the spectra of source \\#19 is the broad silicate feature. \\par \\citet{evans} provide a  crude evolutionary sequence for the youngest stars based on spectral features in the mid-infrared wavelengths (see their Figure 8).  We use this sequence to interpret our CVF spectra.  Sources \\#10, \\#18, and possibly \\#19 are, thus, in the earliest phase of evolution, labeled as 'embedded YSOs' by \\citet{evans}.  The ice features in the spectra of sources \\#10,  \\#18 are seen in the earliest phases when the proto-star is deeply embedded in a circumstellar envelope.  As the circumstellar environment matures, the ices give way to molecular and atomic features during the 'Emerging YSO' phase.  Source \\#14 in our survey is consistent with such an object.  Sources \\#11, \\#15 show PAH emission but otherwise a featureless spectrum.  These are consistent with the 'Young disk' phase of \\citet{evans}.  The evolutionary status of source \\#13 is unclear.  The presence of silicate absorption is indicative of the embedded YSO phase but the lack of features from other ices is puzzling.  Perhaps this is an intermediate object between the Embedded and Emerging YSO phase. \\par In our interpretation of the spectra the ages of stars are of the order $\\sim$few$\\times10^4$~years for the embedded YSOs to $\\sim$few$\\times10^5$~years for stars with 'Young disk'.  These ages are consistent with \\citet{fukui} based on their dynamical estimates using the outflows.  However,  \\citet{chen} estimate the cluster population to be 1-2 Myr based on 3 stars detected by optical and near-IR surveys.  This inconsistency is likely just the difference between the optically detected and the mid-IR detected population of L 1641-N.  The combination of the two estimates establish this region as one of the youngest regions in Orion with active star-formation and age spread of order 2~Myr among the stars.    Our analysis of the ISOCAM broad-band and CVF data produced the following main results: \\par We detect a total of 34 sources in the region covered by the broad-band filters.  The CVF spectra are available for  7 of these sources.  Two sources (\\#10 and \\#18) are only detected at mid-IR wavelengths. \\par We find that the source previously identified as the counter-part to the IRAS detected point-source in the region is not the brightest source at the IRAS 12~\\micron\\ filter wavelengths.  We propose instead that the IRAS source is likely to be combination of sources with \\#10 (in our listing) being the dominant one.  This source outshines all others by roughly a factor of 2. \\par However, source \\#18, previously associated with the IRAS source, shows a larger PSF than other point sources in the area.  This is likely due to the presence of extended emission or emission knots which are unresolved by ISOCAM.  We conclude that source \\#18 remains the most plausible source of the outflows as suggested by \\citet{chen}. \\par A comparison of the near-IR (\\jmh\\ \\vs\\ \\hmk) and mid-IR (\\hmk\\ \\vs\\ $\\isotwothree$) color-color plots shows that while only two of the sources show evidence for ``excess'' emission, at least 85\\% of all sources show evidence for excess emission at mid-IR wavelengths.  For two of these sources, CVF spectra are available and show emission features from PAH.  The presence of excess emission in most sources supports the assertion by \\citet{chen} that these objects are physically associated, or at least co-located. \\par The narrow-band CVF spectra of the 7 sources detected in the CVF show a range of evolutionary status for the stars, ranging from the earliest 'Embedded YSOs' to the 'Young disks'.  When combined with optical and near-IR age estimates, these results show active current star-formation in the region that has been on-going for 2~Myr. \\par The observed absorption features in the spectra of the sources deserve further study in the form of detailed  modeling.  Future work will include full radiative transfer modeling to interpret the physical conditions responsible for the observed spectral features.  We plan to provide these results in a future contribution."
    },
    "0402/astro-ph0402066_arXiv.txt": {
        "abstract": "Evidence of a mis-aligned secondary bar, within the primary bar of the Large Magellanic Cloud (LMC) is presented.  The density distribution and the de-reddened mean magnitudes ($I_0$) of the red clump stars in the bar obtained from the OGLE II data are used for this study. The bar region which predominantly showed wavy pattern in the line of sight in \\citet{a03} was located. These points in the X-Z plane delineate an S-shaped pattern, clearly indicating a mis-aligned bar. This feature is statistically significant and does not depend on the considered value of $I_0$ for the LMC center. The rest of the bar region were not found to show the warp or the wavy pattern. The secondary bar is found to be considerably elongated in the Z-direction, with an inclination of 66$^o$.5 $\\pm$ 0$^o$.9, whereas the undisturbed part of the primary bar is found to have an inclination of 15$^o$.1 $\\pm$ 2$^o$.7, such that the eastern sides are closer to us with respect to the western sides of both the bars. The PA$_{maj}$  of the secondary bar is found to be 108$^o$.4 $\\pm$ 7$^o$.3. The streaming motions found in the H I velocity map close to the LMC center could be caused by the secondary bar. The recent star formation and the gas distribution in LMC could be driven by the mis-aligned secondary bar. ",
        "introduction": "The off-centered stellar bar is one of the most striking features of the Large Magellanic Cloud (LMC). On the other hand, this is one of the least studied and understood feature of the LMC. The near-IR star count maps presented by \\citet{v01} found the bar to be a smooth structure, even though a {\\it peak} in the ellipticity and change in position angle (PA) were found within the central 2$^o$. Recently \\citet{a03} studied the relative distance within the LMC bar using the de-reddened mean magnitudes of red-clump stars and found that the bar is warped and also found structures in the bar. In an attempt to find out the possible reason for these structures, we came across evidence of a possible existence of a secondary bar within the LMC bar. Bars are a common phenomenon in late-type spirals and Magellanic irregulars \\citep{df73}. Recent studies find that the secondary bars within large-scale bars are also common, occurring in about a third of barred galaxies \\citep{J97, l02, es02, E03}.   The evidence of a possible existence of the secondary bar has been found in the literature. Some of the significant references are discussed below. In the R-band isophotes of \\citet{d57}, the first contour shows the bar, two contours immediately next to this suggests a turn in the top-left and bottom-right corners of the bar. This is the first evidence for the twist of the isophotes within the central region of the LMC bar. The isophotes are based on R-band photometry and hence contribution from the old stars dominate. Such an isophotal twist would  manifest as change in the PA of the major axis in the central region, change in ellipticity and a possible counter-rotation in the inner regions of the LMC. The evidence of the change in PA and ellipticity near the LMC center can be found in the literature and some are indicated below. Figure 3 in \\citet{v01} shows the change in PA and ellipticity, figure 6 in \\citet{vahs02} shows change in PA for the carbon stars, figure 5 in \\citet{k98} shows the change in PA for H I. The presence of the above two features found in the central region of NGC 2950 is taken as the photometric signature of a misaligned secondary bar \\citep{c03}. Recent investigations of double-barred galaxies \\citep{E03, J97} indicate that, in general, the radial plot of the ellipticity reflects double peak corresponding to both the bars in the galaxy. In the case of LMC, figure 3 in \\citet{v01}, indicates the first peak, $\\epsilon_{max} \\sim 0.7$ at r$\\sim$1.$^o$0. This is well within the primary bar, which extends to more than 2$^o$ radius. The second peak appears after a radial distance of 2$^o$.0, though it is not very prominent. The corresponding $\\epsilon_{max} \\sim 0.57$. Evidence of negative rotational velocity with respect to the center of LMC is noticed in a number of cases. The study of CH stars by \\citet{HC88} found that some stars have negative galactocentric velocity. Similar cases are also found in the case of planetary nebulae and old star clusters. \\citet{HC88} state that these stars may be related in someway to the bar of the LMC. Recent studies of carbon star kinematics by \\citet{vahs02} find that within the central 1$^o$, the mean rotational velocity is $\\sim$ $-$28 Km/s. Thus all the above observations point to the possible existence of a misaligned secondary bar. All these are features observed in the projected two dimensions and no information is available on its possible appearance in the line of sight.  In the present study, we explore the presence of the secondary bar within the LMC bar in the projected two dimensional X-Y plane as well as in the X-Z plane. We use the red clump stars in the OGLE II catalogue as the probe for this study. The density of red clump stars in the bar region is used to study the projected pattern. The relative distance estimates in the LMC bar based on the de-reddened mean magnitudes of red clump stars, presented in \\citet{a03} are used to study the pattern in the line of sight. ",
        "conclusions": "The possible existence of a secondary bar within the primary bar of LMC is explored here. The photometric signatures of the secondary bar is found in plentiful in the literature, like the twist of isophotes, change in the PA of the major axis and ellipticity {\\it peak} in the central regions. On the other hand, these signatures were never connected with the possible existence of a secondary bar. The motivation to look for a secondary bar came from the perturbations that were noticed in the primary bar by \\citet{a03}. The radial profile of the maximum density on the east side, resembles the brightness profile of the double barred galaxies along the major axis. The secondary bar is seen only on the east side. This indicates that the secondary bar is not symmetric with respect to the optical center. The secondary bar has disturbed only a part of the primary bar, hence we are also able to estimate the parameters of the undisturbed primary bar. The undisturbed primary bar does not show any east west asymmetry. Though the ellipticity of the bars could not be estimated here, the ellipticity estimations in the literature shows that the ellipticity of the secondary bar is $\\epsilon^s_{max} \\sim 0.7$, whereas that of the primary is $\\epsilon^p_{max} \\sim 0.57$. The catalog of double-barred galaxies presented by \\citet{E03} indicates that the average value of the ellipticity of the secondary bar in 49 galaxies is 0.3, whereas the average value for the primary bar is 0.47. The above values were found to be similar to the average ellipticity of the bars presented in \\citet{J97} for 13 galaxies. Both the data also indicate that $\\sim$ 85\\% galaxies show higher ellipticity value for the primary bar when compared to that of the secondary bar. On the other hand, LMC shows higher ellipticity for the secondary bar, which is seen in $\\sim$ 15\\% of the double-barred galaxies. It can be seen that the difference between the PAs of the primary and the secondary bar is very small. The $\\Delta PA = 8^o.0 \\pm 23^o$.0, which is very small, or close to zero within errors. LMC belongs to the group of 6\\% of double-barred galaxies, which show very small value for $\\Delta PA$ \\citep{E03}. The bars are not aligned in the Z-direction, as indicated by the inclination values of 66$^o$.5 $\\pm$ 0$^o$.9 and 15$^o$.1 $\\pm$ 2$^o$.7 for the primary and the secondary bars respectively. The main signature which reveals the central structure as a secondary bar is the mis-alignment in the Z-direction, more than the ellipticity and isophotal signatures. This is the reason why the central structure is claimed to be a mis-aligned secondary bar. The spiral like patterns on the ends of the secondary bar could suggest a possible counter-clockwise rotation in the X-Z plane. The presence of a mis-aligned secondary bar could give rise to kinematic signatures near the central regions. The negative rotational velocity noticed in the central regions could be due to the secondary bar. The mis-alignment could also produce non-circular motions. As the feature in the X-Z plane does not show any ring, either the secondary bar has a slow pattern speed or it is recently formed. More studies are required to understand this newly found feature in the LMC.  The signatures of the secondary bar could be traced in H I observations. The velocity field of the H I, as shown in figure 4, of \\citet{k98}, indicates a steep velocity gradient just to the north of the center of the bar. This is considered as the dynamical evidence of large scale streaming motions. Such steep velocity gradient was also noticed by \\citet{lr92}. It is quite possible that the secondary bar is responsible for the streaming motions. The elongation of the secondary bar in the Z-direction could give rise to a steep velocity gradient. The H I observations by \\citet{k98} show a two armed spiral pattern in their figure 2. Similar spiral arm features were also observed by \\citet{ss03}, figure 2, and they remark that the two arms are connected by H I but not in a structure that looks like the bar as the position of the optical bar is different. The secondary bar could be the feature which is connecting the two spiral arms. The H I observation of \\citet{r84} find clear indications of non-circular velocity near the center. In figure 5 of \\citet{r84} and figure 8 of \\citet{lr92}, the velocity map of H I clearly indicates a lower velocity to the south of the thickly populated iso-velocity contours near the LMC center, and a higher velocity to the north. This suggests a rotation in the sense that the southern part is moving towards us and the northern part is moving away with respect to the center. This corresponds to counter-clockwise rotation in the X-Z plane. If the secondary bar has a counter-clockwise rotation, then this is in good agreement with the rotation seen in the H I velocity maps. Thus the secondary bar of the LMC could be the missing link between the stellar and the gas distribution in the LMC. The recent star formation and the gas distribution in the LMC could be driven by this mis-aligned secondary bar.  I thank T.P.Prabhu and Daniela Vergani for helpful discussions."
    },
    "0402/astro-ph0402585_arXiv.txt": {
        "abstract": "An abundance analysis for the carbon-enhanced, extremely iron-poor ([Fe/H]$\\sim -3.5$) star CS~29498--043 has been obtained using new high-resolution, high signal-to-noise spectra from the Subaru Telescope. The [\\ion{O}{1}] forbidden line at 6300~{\\AA} and the \\ion{O}{1} triplet feature at 7771-7776~{\\AA} are both clearly detected. The overabundance of oxygen is significant ([O/Fe]$>2$). In addition, Na, Co, and Ni abundances have been newly measured. The abundance pattern from C to Ni of this object is quite similar to that of CS~22949--037, another extremely metal-poor star with large excesses of C, N, O, and the $\\alpha$-elements. The abundance patterns of these two stars suggest the existence of supernovae progenitors that ejected relatively little material from their iron cores during the very early era of nucleosynthesis in the Galaxy. The metallicity in these objects, when one includes the elements C, N, and O in the tally of total metals, is not as low as in the most metal-poor stars, suggesting the existence of quite different formation processes for these iron-deficient objects than pertain to the bulk of other metal-deficient stars. ",
        "introduction": "\\label{sec:intro}  The surface chemical composition of extremely metal-poor stars is believed to reflect the yields of heavy elements produced by the first generations of massive stars in our Galaxy. Reported large variations of the elemental abundance patterns of some species (in particular those associated with carbon and the neutron-capture elements) found in very metal-poor stars (e.g., McWilliam et al. 1995; Ryan, Norris, \\& Beers 1996; Cayrel et al. 2004) suggest a diversity of the nucleosynthesis and explosion mechanisms of the supernovae which were presumably responsible.  Our previous study of the extremely metal-poor star CS~29498--043 ([Fe/H]$\\sim-3.7$)\\footnote{[A/B] = $\\log(N_{\\rm A}/N_{\\rm B})- \\log(N_{\\rm A}/N_{\\rm B})_{\\odot}$, and $\\log \\epsilon_{\\rm A} = \\log(N_{\\rm A}/N_{\\rm H})+12$ for elements A and B.} revealed remarkable overabundances of C, N, Mg, and Si in this object \\citep{aoki02a,aoki02b}. The chemical nature of this object, which is similar to that of another extremely metal-poor star, CS~22949--037 (McWilliam et al. 1995, Norris, Ryan \\& Beers 2001; Depagne et al. 2002), suggests that these objects formed from the yields of supernovae that ejected relatively little material from the regions surrounding their iron cores at the time of their explosion \\citep[e.g., ][]{tsujimoto03,umeda03a}. A remaining important question is the origin of the large overabundances of C and N, because these two elements are expected to be significantly affected during the evolution of low-mass stars. In order to distinguish the contributions of low- and high-mass stars, determinations of the O abundance, as well as those of the $\\alpha$ elements, are quite important. The O abundance is also vital to estimate the metallicity of this object. In this paper, we report a new analysis of the chemical composition of CS~29498--043, including the O abundance, derived from high-resolution, high signal-to-noise spectra obtained with the Subaru Telescope. ",
        "conclusions": "Recent studies of O abundances in very metal-poor stars, within the framework of one-dimensional model atmospheres, have revealed the presence of a possible increase of O/Fe ratios with decreasing metallicity in the range of [Fe/H]$<-2$.  Figure \\ref{fig:ofe} shows [O/Fe] as a function of [Fe/H] for our object and for other stars described in the literature \\citep{israelian01, nissen02, cayrel03, bessell04}. The filled symbols show the [O/Fe] values determined from the [\\ion{O}{1}] line, while open symbols indicate those obtained from either the triplet lines or OH molecular lines. A discrepancy between the results from the different indicators is seen, as discussed in subsection 3.1. However, the O overabundance of CS~29498--043, as well as that of CS~22949--037 \\citep{depagne02}, is unusually high, compared with the other very metal-poor stars ([O/Fe]$\\lesssim +1.0$). The [O/Fe] of the most iron-deficient star presently known, HE~0107--5240 ([Fe/H] = -5.3; Christlieb et al. 2002, 2004) is also as high as found in these two stars \\citep{bessell04}. We note that, even though three of the four objects with [Fe/H]$<-3.5$ have [O/Fe]$\\gtrsim +2.0$, there are stars in this metallicity range in which oxygen lines are {\\it not} detected, and hence whose oxygen abundances are not as high as in these three stars. Though it remains possible that [O/Fe] values continuously increase with decreasing [Fe/H] for stars with [Fe/H]$<-3$, the unusually high oxygen abundances in the above three stars are more likely to arise from a different nucleosynthesis history from that pertaining to other metal-deficient stars.  Figure \\ref{fig:res} shows the relative elemental abundances ([X/Fe]), as a function of atomic number, for CS~29498--043 and CS~22949--037. The similarity of the abundance patterns between these two stars is remarkable. The present analysis for the new spectrum of CS~29498--043 confirms the similarity for O and Na, which show large overabundances in both stars. The [O/Fe] value of CS~29498--043 is higher, by 0.36~dex, than that of CS~22949--037, if the same solar O abundance is adopted. However, we do yet not insist on a real difference of the O abundances between the two stars, taking account of the uncertainty in the O abundance determination (as seen, e.g., in the discrepancy in the results between the [\\ion{O}{1}] line(s) and the O triplet). Recall that for the present analysis, the O abundances in both objects were determined from the [\\ion{O}{1}] 6300~{\\AA} line.  The solid line in Figure \\ref{fig:res} shows the abundance ratios predicted by the supernova model of \\citet{umeda03a} for a 30~$M_{\\odot}$ star, assuming substantial mixing and fall-back (Umeda \\& Nomoto 2004). The large overabundances of C, O, Mg, and Si relative to iron-peak elements are well reproduced by this model. The large overabundance of N in these stars could be associated with the operation of the CN-cycle during their evolution on the red-giant branch \\citep{depagne02}. Note as well that the Co/Fe and Ni/Fe ratios of CS~29498--043 might be slightly lower than those of CS~22949--037. The upper-limit on the Zn abundance of CS~29498--043 is also lower than the Zn abundance of the other object. In the Umeda \\& Nomoto (2004) models, lower Co and Zn abundance ratios indicate a smaller explosion energy and/or more effective mixing. More accurate estimates of the abundance ratios of these elements in the present objects, as well as for other similar stars that might be found in the future, are desirable to better constrain the parameters of the present explosion models.  Finally, we would like to consider the metallicity of these objects and their extremely low iron abundances. We here define the metallicity by the {\\it total abundance} of C, N, O, Mg, Si, and Fe, i.e., [A/H]=[(C+N+O+Mg+Si+Fe)/H]. The metallicities of CS~29498--043 and CS~22949--037, using this definition, are [A/H]$=-1.26$ and $-2.00$, respectively. Oxygen is the most important contributor to the metallicity in these stars (60--70\\% of the metallic species). When we adopt the above definition of metallicity, these three objects might be regarded as {\\it extremely iron-deficient stars} rather than {\\it extremely metal-poor stars}.  The metallicity of the most iron-poor ([Fe/H]$=-5.3$) star HE~0107--5240 \\citep{christlieb02, christlieb04} is [A/H]$=-1.77$, using this definition, adopting [O/Fe]$=+2.4$ \\citep{bessell04}, and assuming [Si/Fe]=0.0. In this case, however, the dominant contributor to the metallicity is carbon (about 95\\%). The abundance pattern of this object can also be explained by supernova models by \\citet{umeda03a}, assuming a larger mixing region and a smaller matter ejection factor. Although this object shows no clear overabundance of $\\alpha$-elements, the origin of its peculiar abundance pattern may be related to those of CS~29498--043 and CS~22049--037.  The extremely iron-deficient ([Fe/H]$<-3.5$) stars found in previous studies (Norris et al. 2001, Carretta et al. 2002, Franc\\c{o}is et al. 2003), in general, {\\it do not exhibit} large excesses of lighter elements, except for CS~22949--037 (Norris et al. 2001). These stars can be regarded as extreme cases of metal-deficient stars with higher iron abundance ($-3.5<$[Fe/H]$\\lesssim-2.5$), which would form from interstellar matter polluted by first-generation supernovae. Their extremely low metallicity suggests high explosion energies of the supernovae that provided metals, and induced the formation of second-generation stars \\citep[e.g., ][]{cioffi88}.  Hence, the formation mechanism of the above three iron-deficient stars with excesses of C and O may be quite different from the other iron-deficient stars.  Even though their total atmospheric metallicities are rather high ([A/H]$\\gtrsim -2$), these three stars are also expected to be produced from material polluted by first-generation supernovae, because of their peculiar abundance patterns.  However, the extremely low iron abundances of these objects might be regarded as resulting from the particular yields of their progenitors, rather than from mixing processes within the interstellar matter from which they were formed. The relatively high total metallicity of these objects may indicate low explosion energies of the supernovae which induced the formation of these objects.  If the species C and O were present in the interstellar matter of the birth clouds of these objects they would supply important cooling sources, in particular for HE~0107--5240, hence their overabundances may help understand the possibility of low-mass star formation in the early Galaxy, as suggested by \\citet{umeda03a} (see also Bromm \\& Loeb 2003). This may also explain the existence of the plethora of C-rich objects amongst extremely iron-deficient stars (Norris, Ryan, \\& Beers 1997; Beers 1999; Rossi et al. 1999). Further abundance studies of stars with extremely low iron abundances are indispensable to establish the proper classification of these stars, and to understand the formation mechanism(s) of early-generation stars in our Galaxy."
    },
    "0402/astro-ph0402250_arXiv.txt": {
        "abstract": "\\noindent Several interacting systems exhibit at the tip of their long tidal tails massive condensations of atomic hydrogen, which may be the progenitors of Tidal Dwarf Galaxies. Because, quite often, these tails are observed edge-on, projection effects have been claimed to account for the large HI column densities measured there. Here we show that determining the velocity field all along the tidal features, one may disentangle projection effects along the line of view from real bound structures. Due to its large field of view, high spectral and 2D spatial resolutions, Fabry-Perot observations of the ionized gas are well adapted to detect a kinematical signature of either streaming motions along a bent tidal tail or of infalling/rotating material associated with a forming TDG. Spectroscopic observations also allow to measure the dynamical masses of the TDGs that are already relaxed and check their dark matter content. ",
        "introduction": "The most impressive and surely most studied interacting systems, such as the Antennae galaxies, exhibit long optical tidal tails that may extend up to 100 kpc. HI observations of such colliding galaxies have shown that the stellar tails have a gaseous counterpart that is usually even more prominent and contain a large fraction of the total atomic gas present in the system. In a few interacting galaxies, the HI tidal tails exhibit some gas concentrations that have apparent masses of up to few $10^{9}$ solar masses. Those are the progenitors of the Tidal dwarf galaxies (TDGs), at least the more massive ones. They are typically found at the tip of the optical tidal tails at distances between 30 and 100 kpc from the merging disks. They might be as massive as the Magellanic Clouds, and on top of being rich in atomic hydrogen gas (Duc et al. 2000), they contain high quantities of molecular gas (Braine et al. 2001, and his contribution in this volume) and form stars with a rate as high as in blue compact dwarf galaxies (Duc \\& Mirabel 1998). The amount of stellar and gaseous material in several TDGs suggests that they are gravitationally bound, although we do not have yet conclusive observational evidence of such a self-gravitating TDG, kinematically independent from its host tidal tail. Moreover if dark matter is made of collisionless material distributed in a large halo, it should not be ejected with tidal stellar and gaseous debris pulled out from colliding disks (Barnes \\& Hernquist, 1992). In that case, TDGs should not contain a significant amount of DM whereas ordinary dwarf galaxies seem to possess a lot of it. This can be checked by determining the dynamical masses of TDGs.  However, the existence of the most massive Tidal Dwarf Galaxies or even of their gaseous progenitors as independent entities has been challenged (see Hibbard et al., in this volume). Indeed, an apparent accumulation of tidal material could in reality be the result of a projection effect. In the 3D space, tidal tails are curved. Seen edge-on, they appear as linear structures and may present at their tip fake mass concentrations due to the presence of projected material along the line of sight. Using numerical simulations, we have shown that such projection effects have a kinematical signature (Bournaud et al. 2004). Indeed, as shown in Fig.~1, the large-scale velocity gradient corresponding to the streaming motions along the expanding tails changes its sign before their apparent extremity whenever a projection effect exists.   \\begin{figure}[ht!] \\centering \\includegraphics[width=7cm]{amram1.fig1.eps} \\caption{Numerical simulations  of two interacting galaxies. The kinematics of one of the tidal tails is analyzed as if it were observed edge-on, assuming that the line-of-sight is aligned with its extremity. This projection effect may result in an apparent accumulation of matter at the tip of the observed tail. The position-velocity diagram derived from this simulation shows a velocity gradient, mainly related to streaming motions along the tail. The sign of this gradient changes before the extremity of the tidal tail, a result which is obtained whenever a part of the tail is aligned with the line-of-sight. Conversely, a velocity gradient that has a constant sign rules out the presence of a projection effect. In addition to the large-scale motions, each TDG or gravitational clump can have an inner velocity gradient, related to its own dynamics.} \\label{simu} \\end{figure}   Such a study requires observations with an instrument enabling the access to both the large-scale and small-scale dynamical structures. Synthesis arrays of radiotelescopes allow to probe the largest ones but do not have enough sensitivity to reach the $1\\arcsec$ spatial resolution required to sample correctly the nearby TDGs progenitors. Because of their low surface brightness, a direct study of the stellar kinematics of tidal features is at the limits of today's telescopes and detectors. The internal dynamics can be more easily approached in the ionized gas component through spectroscopic observations of the emission lines. Slit '1D' spectroscopy already provides some information (see Weilbacher et al. 2003) but is largely insufficient given the complex morphology of colliding galaxies. The Fabry-Perot technique appears as an ideal tool as it combines an integral-field capability, a high spatial and spectral resolution and a large field of view. We present here Fabry-Perot observations of several interacting systems where TDG candidates had been previously identified. Observations were carried out at the European Southern Observatory 3.6m telescope and at the Canada-France-Hawaii 3.6m telescope. The pixel size on the sky varies from 0.86 to 0.91 arcsec; the FOV from 170''$\\times$170'' to 440''$\\times $440'' and the velocity sampling from 10 to 16 km.s$^{-1}$. The data reduction procedure has been extensively described (Amram et al. 1998 and references therein). ",
        "conclusions": "\\noindent We have shown that the  kinematics of tidal tails may be used to check whether the claimed massive TDG candidates observed at their tip  are real bound entities or the result of projection effects. With their high spatial resolution, three-dimensional  Fabry-Perot  observations in the H$\\alpha$ line turn out to be well adapted to the problem. Analyzing position-velocity diagrams derived from FP data obtained at the CFHT and at the ESO 3.6m, we were able to  rule our projection effects as the main contributor to the condensations seen in the northern tail of Arp~242 (NGC 4676) and in the southern tail of Arp~105.  We cannot  exclude that they play a role in the merger IC~1182. The existence of massive entities observed near the tip of several tidal tails is actually supported by the numerical simulations of interacting galaxies by Bournaud, Duc \\& Masset (2003) in which tidal objects with comparable masses are formed.  We have also studied the inner kinematics of a few TDG candidates. Several velocity gradients were found. The kinematics of the northern mass concentration of NGC~5291 is consistent with that expected for a  self-gravitating rotating  object. Its inferred dynamical and luminous mass compare, but higher resolution observations are required to prove that the object does not contain a certain amount of  dark matter."
    },
    "0402/astro-ph0402299_arXiv.txt": {
        "abstract": "We discuss the spectrum of scalar density perturbations from warm inflation when the friction coefficient $\\Gamma$ in the inflaton equation is dependent on the inflaton field. The spectral index of scalar fluctuations depends on a new slow-roll parameter constructed from $\\Gamma$. A numerical integration of the perturbation equations is performed for a model of warm inflation and gives a good fit to the WMAP data for reasonable values of the model's parameters. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0402/astro-ph0402244_arXiv.txt": {
        "abstract": "Recent interferometric observations of the brightest and angularly largest classical Cepheid, $\\ell$~Carinae, with ESO's VLT Interferometer (VLTI) have resolved with high precision the variation of its angular diameter with phase. We compare the measured angular diameter curve to the one we derive by an application of the Baade-Wesselink type infrared surface brightness technique, and find a near-perfect agreement between the two curves. The mean angular diameters of $\\ell$~Car from the two techniques agree very well within their total error bars (1.5\\,\\%), as do the derived distances (4\\,\\%). This result is an indication that the calibration of the surface brightness relations used in the distance determination of far away Cepheids is not affected by large biases. ",
        "introduction": "Cepheid variables are fundamental objects for the calibration of the extragalactic distance scale. Distances of Cepheids can be derived in at least two different ways: by using their observed mean magnitudes and periods together with a period-luminosity relation, or by applying a Baade-Wesselink (hereafter BW) type technique to determine their distances and mean diameters from their observed variations in magnitude, color and radial velocity.  This latter technique has been dramatically improved by the introduction of the near-infrared surface brightness method (hereafter IRSB) by \\citet{Welch94}, and later by \\citet{FG97} who calibrated the relation between the $V$-band surface brightness and near-infrared colors of Cepheids. For this purpose, they used the observed interferometric angular diameters of a number of giants and supergiants bracketting the Cepheid color range. This method has been applied to a large number of Galactic Cepheid variables, for instance by \\citet{GFG97}, \\citet{GFG98}, and \\citet{Storm04}.  Applying the surface-brightness relation derived from stable stars to Cepheids implicitly assumes that the relation also applies to pulsating stars. The validity of this assumption can now be addressed by comparing direct interferometric measurements of the angular diameter variation of a Cepheid to the one derived from the IRSB technique. It has recently been shown by \\citet{Kervella04}, hereafter K04, that the VLT Interferometer on Paranal is now in a condition to not only determine accurate {\\it mean} angular diameters of nearby Cepheid variables, but to follow their angular diameter {\\it variations} with high precision. Using the Palomar Testbed Interferometer, \\citet{lane00,lane02} resolved the pulsation of the Cepheids $\\zeta$~Gem and $\\eta$~Aql as early as 2000, but the comparison we present in this letter is the first where error bars on the derived distance and linear diameter are directly comparable at a few percent level between the interferometric and IRSB techniques.  The star that we will discuss in this letter, $\\ell$~Car, is the brightest Cepheid in the sky. Its long period of about 35.5 days implies a large mean diameter, which together with its relatively short distance makes it an ideal target for resolving its angular diameter variations with high accuracy. In this paper, we will compare the interferometrically determined angular diameter curve of $\\ell$~Car with that determined from the IRSB technique, and we will demonstrate that the two sets of angular diameters are in excellent agreement. Based on the available high-precision angular diameter and radial velocity curves for this star, we will also derive a revised value of its distance and mean radius.  Several authors \\citep{Sasselov94, Marengo03, Marengo04} have pointed out potential sources of systematic uncertainties in the determination of Cepheid distances using the interferometric BW method. In particular, imperfections in the numerical modeling of Cepheid atmospheres could lead to biased estimates of the limb darkening and projection factor. We will discuss the magnitude of these uncertainties in the case of $\\ell$~Car. ",
        "conclusions": "The main point of our paper is to show that with a consistent treatment of the data, the internal accuracy of both methods (IRSB or interferometry) is extremely good: the angular diameter variation observed using the VLTI agrees very well with that derived from the $F_V(V-K)$ version of the IRSB technique as calibrated by FG97. For all the interferometric measurements, the corresponding IRSB angular diameter at the same phase lies within the combined 1$\\sigma$ error bars of the two measurements (Fig.~\\ref{fig.LDfit}). Even more importantly, the mean angular diameter of the Cepheid as derived from both independent sets of angular diameter determination are in excellent agreement, within a few percents.  Unfortunately, this is not equivalent to say that the Cepheid distance scale is calibrated to a 1\\% accuracy. We have drawn attention to remaining sources of systematic errors that can affect Cepheid radii and distances up to several percents. As an illustration of these sources, K04 obtain a distance $d= 603^{+24}_{-19}$\\,pc for $\\ell$~Car, while we obtain $d = 566^{+24}_{-19}$\\,pc from the same interferometric data.  We have already shown that most of the 6\\,\\% difference (equivalent to $1.3\\,\\sigma$) can be explained by the use of different radial velocity data and projection factor. Another thing to consider is the phase interval used. K04 used measurements over the whole pulsation cycle whereas in the IRSB technique, one avoids the phase interval 0.8--1 (Fig.~\\ref{fig.svb}). During that phase interval, that corresponds to the rebound of the atmosphere around the minimum radius, energetic shock waves are created. As discussed by \\citet{Sabbey95}, they produce asymmetric line profiles in the Cepheid spectrum. Recent modeling using a self-consistent dynamical approach also show that the $\\tau=1$ photosphere may not be comoving with the atmosphere of the Cepheid during its pulsation, at the 1\\,\\% level \\citep{Nardetto04}. Such an effect would impact the $p$-factor, modify the shape of the radial velocity curve, and thus bias the amplitude of the radius variation, possibly up to a level of a few percents. As the BW method (either its classical or interferometric versions) relies linearly on this amplitude, a bias at this level presently cannot be excluded.  The interferometric BW method is currently limited to distances of 1--2 \\,kpc, due to the limited length of the available baselines. The IRSB technique on the other hand can reach extra-galactic Cepheids as already demonstrated by \\citet[for the LMC]{gieren00} and by \\citet[for the SMC]{Storm04}. Using high precision interferometric measurements of $\\ell$~Car and other Cepheids, it will be possible to calibrate the IRSB method down to the level of a few percents. From the present comparison, we already see that this fundamental calibration will be very similar to the calibration found by FG97 and \\citet{Nordgren2002}."
    },
    "0402/astro-ph0402591_arXiv.txt": {
        "abstract": "{  How do I find the optimal photometric system for a survey?  Designing a photometric system to best fulfil a set of scientific goals is a complex task, demanding a compromise between often conflicting scientific requirements, and being subject to various instrumental constraints.  A specific example is the determination of stellar astrophysical parameters (APs) -- effective temperature, surface gravity, metallicity etc.\\ -- across a wide range of stellar types.  I present a novel approach to this problem which makes minimal assumptions about the required filter system.  By considering a filter system as a set of free parameters (central wavelengths, profile widths etc.), it may be designed by optimizing some figure-of-merit (FoM) with respect to these parameters.  In the example considered, the FoM is a measure of how well the filter system can `separate' stars with different APs.  This separation is vectorial in nature, in the sense that the local directions of AP variance are preferably mutually orthogonal to avoid AP degeneracy.  The optimization is carried out with an evolutionary algorithm, a population-based approach which uses principles of evolutionary biology to efficiently search the parameter space. This model, HFD (Heuristic Filter Design), is applied to the design of photometric systems for the Gaia space astrometry mission. The optimized systems show a number of interesting features, not least the persistence of broad, overlapping filters. These HFD systems perform as least as well as other proposed systems for Gaia -- as measured by this FoM -- although inadequacies in all of these systems at removing degeneracies remain. Ideas for improving the model are discussed.  The principles underlying HFD are quite generic and may be applied to filter design for numerous other projects, such as the search for specific types of objects or photometric redshift determination. ",
        "introduction": "Surveys of stellar populations are directed at improving our understanding of their formation and evolution. One of the most important ingredients of such surveys is stellar photometry and/or spectroscopy as a means to determine fundamental stellar parameters. These are, in the first instance, atmospheric parameters -- effective temperature, surface gravity and chemical abundances -- from which we may derive stellar masses, radii and ages.  A fundamental question facing the designers of such surveys is what spectra and/or photometric systems are optimal for this purpose. Given the constraints of telescope size and survey duration, the designer must generally trade off spectral resolution with signal-to-noise ratio (SNR), limiting magnitude, number of sources and sky coverage. Some aspects of the design may be obvious from the scientific goals, but the optimal settings of many others will remain uncertain.  Deep surveys of large numbers of objects will often be forced to employ photometry rather than spectroscopy due to confusion and SNR considerations. Given well defined scientific goals, the designer must decide how many filters to use, with what kind of profiles, where to locate them in the spectrum and how much integration time to assign to each. This is usually achieved via a manual inspection of typical target spectra.   But if the survey is intended to establish multiple astrophysical parameters (APs) across a large and varied population of objects, then this method of filter design is unlikely to be very efficient or even successful. Manually placing numerous filters and adjusting their profiles to simultaneously satisfy many different -- and often conflicting -- requirements is likely to be extremely difficult, especially given the vast number of permutations of filter parameters possible.   Even if a reasonable filter system could be constructed in this way, we would not know whether a better filter system exists subject to the same constraints. Is there not a more systematic approach to constructing filter systems?  The approach developed in this article is to use a representative grid of (synthetic) spectra with known APs to construct a filter system in a heuristic fashion.  The grid represents the scientific goals of the survey.  Let us assume that for a given filter system we can calculate a figure-of-merit which is a measure of how accurately the filter system can determine the APs of the grid spectra.  If we consider the filter system as a set of free parameters (central wavelengths, widths, profile shapes etc.), then we may construct a filter system by optimizing the figure-of-merit with respect to these filter parameters. This approach has the advantage that it can exploit the extensive literature on optimization techniques.  The specific technique used here is a type of evolutionary algorithm, a population-based technique designed to perform a stochastic yet directed search of the parameter space, adopting features of biological evolution (section~\\ref{EAs}).  The underlying principle of my approach is to make few prior assumptions about the required filter system and to let the optimization proceed freely within the constraints laid down by the scientific goals and other instrumental considerations.  The model itself, HFD (Heuristic Filter Design), will be described in detail in section \\ref{hfd_model}, but it is worth highlighting now that a crucial aspect is to establish a suitable figure-of-merit. The most obvious would be some average (over the grid) of the precision with which APs are determined. This could be achieved by any one of several regression methods -- e.g.\\ nearest neighbours or neural networks -- used to approximate the mapping between the data space and the AP space, although doing this well for the multiparameter stellar problem is far from trivial (Bailer-Jones \\cite{bj02}, \\cite{bj03}). Furthermore, because HFD works through many (ca.\\ 10$^5$) candidate filter systems, fitting a high-dimensional regression model in each case would be unbearably time consuming.  It turns out that an explicit determination of the performance of the filter system in these terms is not actually necessary.  A suitable figure-of-merit can be constructed when we consider what a filter system does.  Its primary function is to define metrics (e.g.\\ colours) which cluster together similar objects and which separate out dissimilar objects.  A simple example is star--quasar separation. If the filter system is designed to determine a continuous AP it should separate objects in proportion to their differences in this AP. In doing this it defines a local vector in the data space along which the AP varies monotonically. (Only once such a separation has been achieved can this vector be calibrated in terms of the AP.)  When determining multiple parameters (e.g.\\ \\teff\\ and extinction), it is furthermore essential that the local vectors for each parameter are near orthogonal, otherwise a local AP degeneracy exists.  Thus `separation' of sources in HFD must be understood in this more general, vectorial sense.  Section~\\ref{fitness} describes how a figure-of-merit is constructed to respect these requirements.  HFD is used in section \\ref{application} to design filter systems for the Gaia Galactic Survey Mission and their performance is compared to other proposed systems. Gaia is a high precision astrometric and photometric mission of the European Space Agency to be launched in 2010. Operating on the principles of Hipparcos, but exceeding its capabilities by orders of magnitudes, Gaia will determine positions, proper motions and parallaxes for the $10^9$ stars in the sky brighter than V=20 (ESA \\cite{esa00}; Perryman et al.\\ \\cite{perryman01}).  Its primary objective is to study the structure, formation and evolution of our Galaxy. To achieve this, the kinematical information must be complemented with multi-band photometry to determine physical stellar parameters.  HFD is used to design appropriate UV/optical/NIR (i.e.\\ CCD) photometric systems for this survey.  Section \\ref{discussion} then gives a critical discussion of the HFD approach, its features and limitations and discusses how the approach could be extended and approved.  Section \\ref{conclusions} summarises the main results and conclusions of this work. ",
        "conclusions": "I have introduced a novel approach to the design of photometric systems via optimization of a figure-of-merit of filter system performance.  In the present incarnation, this figure-of-merit (or fitness) measures the ability of a filter system to determine multiple stellar astrophysical parameters (APs), by calculating the separation in the data (filter) space between stars with different APs.  The better that sources can be separated (in signal-to-noise units) according to their AP differences, the better the filter system.  This separation is vectorial in nature, meaning that the figure-of-merit is also proportional to the angle between the vectors which define the directions of local variance of each AP: In the ideal filter system these vectors would be mutually orthogonal at all points in the AP space, thereby removing any degeneracy between APs. The fitness is calculated via an instrument model for a grid of spectra, which sample stellar parameters the photometric system must determine.  The optimization is performed with an evolutionary algorithm. In this approach, a population of filter systems is evolved according to the principle of natural selection, such that the fitter filter systems are more likely to survive and to produce more `offspring'. Reproduction takes place by combining or mutating selected parents, resulting in changes of the filter parameters (central wavelength, profile width, integration time), thus providing a stochastic yet directed search of the filter parameter space.  This model, HFD (Heuristic Filter Design), has been applied to design CCD photometric systems for the Gaia Galactic Survey Mission.  The systems were optimized to separate the four APs effective temperature, \\teff, metallicity, \\feh, surface gravity, \\logg, and interstellar extinction toward the star, \\extinct.  Recurrent characteristics of the resulting filter systems are broad overlapping filters, although filters with a half-width above 1500\\,\\AA\\ were consistently disfavoured.  The preferred broadness is not surprising when one realises that each of the APs has a coherent effect on the data over a wide wavelength range.  Narrower filters were found not to improve significantly the orthogonality (vector separation). This tendency toward broader filters than have hitherto been adopted for the Gaia filter systems -- and for stellar parametrization in general -- is one of the main results of this application of HFD. Likewise is the related tendency toward overlapping filters.  This may be indicative of a more efficient use of a multi-dimensional data space than non-overlapping systems.  In terms of the scalar separation of sources, the HFD filter systems perform well at the Gaia target magnitude of G=15, although at the limiting magnitude of G=20 the separation for \\feh\\ and \\logg\\ is unsatisfactory. More significantly, the vector separation is inadequate in parts of the AP grid, and between \\teff\\ and \\extinct\\ in particular considerable degeneracy remains. Yet other systems proposed for Gaia show similar difficulties, and overall HFD performs at least as well as or better than these.  It remains to be seen whether these are intrinsic limitations of broad and medium band photometry for these instrument models of whether improvements to the fitness function alter this.  Either way, this systematic approach to filter system design embodied in HFD shows considerable promise.  A number of improvements to HFD to address some deficiencies were suggested, including the use of more efficient search operators, the use of secondary grids or generalization to nonlinear separation, and the incorporation of multiobjective optimization methods. The latter allows the different objectives of the filter system to be optimized separately, thus avoiding having the problem of weighting and combining heterogeneous objectives.  Specifically with regard to Gaia, HFD may be developed in a number of ways. The most significant is perhaps the inclusion of parallax information: the parallaxes from Gaia will permit an accurate determination of the luminosity and (via \\teff) the radius of many stars, reducing the need to determine \\logg. By including the parallax error model in HFD, filter systems better matched to the available astrometry can be designed.   Beyond this application, HFD represents a generic approach to formalizing filter design by casting it as an optimization problem with few prior assumptions. The key steps are the parametrization of the filter system, the construction of a figure-of-merit, and the design of appropriate genetic operators to search the parameter space. Evolutionary algorithms are particularly appropriate for this problem because the fitness landscape in which the optimization is performed will is frequently complex and noisy.  While these steps are nontrivial, HFD provides a general framework for applying this approach to many other problems. These include the identification of particular types of objects, such as ultra cool dwarfs or metal poor stars, star/quasar separation, the spectral classification of galaxies, and photometric redshift determination.  This is applicable not only to future large scale surveys, but also for more modest surveys on existing ground-based facilities."
    },
    "0402/astro-ph0402072_arXiv.txt": {
        "abstract": "We have made {\\em ROSAT} PSPC and HRI X-ray observations to study the intracluster gas surrounding the powerful radio source Hercules A. The cluster is luminous in X-rays ($L_{\\rm bol} = 4.8\\times 10^{37}$ W), although apparently poor in optical galaxies, and the host of the radio source is the central dominating galaxy of the cluster.  The azimuthally-averaged X-ray surface brightness profile is well fitted by a modified King ($\\beta$) model, with core radius $r_c = 121 \\pm 10$ kpc and $\\beta = 0.74 \\pm 0.03$, but the cluster is elongated parallel to the radio source, especially on the scale of the radio lobes, and fits to individual quadrants give a core radius 50 per cent larger along the radio axis. Part of this elongation appears to be associated with enhanced X-ray emission superimposed on the outer radio lobes, which extend to just over $2r_c$.  There are no obvious depressions in the X-ray emission coincident with the radio lobes, as expected if the relativistic plasma displaces the ICM. However, we show that these depressions may be quite weak, essentially because the main part of the lobes are outside the cluster core. From the surface brightness profile for the PSPC data the X-ray emission extends out to $\\sim$ 2.2 Mpc radius.  In the absence of the powerful jets (which must be a transient phenomenon on cosmological timescales), we would expect a cooling flow at the centre of the cluster; but currently it must be substantially disturbed by the expansion of the radio lobes.  The PSPC spectrum reveals a cool component of the ICM with $0.5 \\lta kT \\lta 1$ keV in addition to the $\\approx 4$ keV component detected by {\\em ASCA} and {\\em BeppoSAX}. The central cooling time could be as low as 2 Gyr if the cool component is centrally concentrated, otherwise it is around 6 Gyr. Cooling is significant on a Hubble time to a radius of about 90 kpc. The modelled central electron density of  $n_0 = 1.0 \\times 10^{4}$ m$^{-3}$ is typical for modest cooling flows.  Finally, we have detected faint X-ray emission from a compact central source, with size $< 15$ kpc and luminosity $\\approx 2 \\times 10^{36}$~W. ",
        "introduction": "Hercules A (3C\\,348) is the fourth brightest DRAGN\\footnote{% Double Radiosource Associated with Galactic Nucleus; see \\citet{Leahy1993} or Leahy, Bridle \\& Strom ({\\tt http://www.jb.man.ac.uk/atlas/}) for a full definition.} in the sky at low frequencies. At a low redshift of $z=0.154$, its power at 178 MHz is $P_{\\rm 178\\,MHz} = 1.9\\times10^{27}$ W Hz$^{-1}$ sr$^{-1}$, with $H_0 = 65$ kms$^{-1}$ Mpc$^{-1}$ and $q_0 = 0$ (used throughout this paper).  Hercules A is identified with a very elongated cD galaxy \\citep[e.g.][]{Sadun.etal1993,Baum.etal1996}, with absolute magnitude $-23.75$ in the R-band \\citep{Owen.etal1989}. It appears to lie at the center of a poor, faint cluster, although measurements of the cluster richness may be uncertain \\citep*{Owen.etal1989,Yates.etal1989,Barthel.etal1996}.  Radio galaxies with the radio and optical luminosity of Her A nearly all show hotspot-dominated Fanaroff-Riley II structure \\citep{LO96}, but Her A is is an exception \\citep{Dreher.etal1984}. Its structure is dominated by its twin jets, which are quite different in appearance. The eastern one is the brightest (highest flux density) radio jet in the sky, and contributes a substantial fraction of the radio luminosity of Her A; a weaker jet to the west leads to a striking series of shells which dominate the western lobe. There are no compact hotspots.  The structure is formally FR class I, but does not resemble typical FR\\,I objects in detail.  X-ray emission from the Hercules A cluster was first detected by the {\\em Einstein Observatory} \\citep{Feigelson.etal1983,Dreher.etal1984}. They estimated the 0.2--4 keV luminosity as 2 $\\times 10^{37}$~W. This is typical of a richness 0 to 1 Abell cluster \\citep{Abell1958}. In their deprojection analysis of the {\\em Einstein} data on clusters of galaxies, \\citet*{White.etal1997} list Hercules A as a possible large cooling flow, although their best-fit model had a zero inflow rate due to their low spatial resolution.  \\citet{Barthel.etal1996} have shown that Hercules A and other radio galaxies embedded in dense cluster gas are anomalously radio-loud (for their far-infrared luminosity), as expected because their confinement is more effective, reducing adiabatic expansion.  In this first paper of a series of three, we report observations of Hercules A in X-rays made with the {\\em ROSAT} PSPC and HRI detectors. \\citet[hereafter Paper II]{Gizanic} present new VLA observations of this powerful DRAGN, and we discuss the results derived from the spectral index and projected magnetic field/fractional polarization images all at 1$''\\!$.4 resolution. In Paper III (Gizani, Leahy \\& Garrington in preparation) we combine the results of Papers I and II to study the magnetic field of the cluster gas surrounding the DRAGN. Preliminary reports have already been published by Gizani \\& Leahy (1996, 1999).\\nocite{Gizania,Gizanib} Our observations have also been discussed, in less detail, by \\citet*{Siebert.99}.  Sections 2 and 3 describe the observations and data reduction for the PSPC and HRI respectively. In Section 4 we analyse the radial surface brightness profile to infer the density distribution in the Her A cluster. In section 5 we compare the radio and X-ray structures to search for detailed correspondences, and review observations of related objects to put our results into context. Section 6 gives our conclusions. ",
        "conclusions": "Our {\\em ROSAT} PSPC and HRI X-ray observations of the intracluster gas in the Hercules A cluster have revealed an extended X-ray emission, extending to radius 2.2 Mpc (PSPC data) and a weak central peak (HRI data).  Based on discrepancies between single-temperature fits to the PSPC, {\\em ASCA} and {\\em BeppoSAX} spectra, we argue that the intracluster gas contains at least two phases, dominated by a $\\approx 4$ keV component (seen by {\\em ASCA} and {\\em BeppoSAX}), but with around 15 per cent of the emission in the {\\em ROSAT} band from material at 0.5--1 keV.  This model significantly improves the PSPC fit compared to one with a single temperature. The 0.1--2.4 keV (absorbed) flux density is $3.0\\times 10^{-15}$ W m$^{-2}$, and the bolometric luminosity is $4.8 \\times 10^{37}$ W. There is no evidence for absorption above the Galactic value.  We have detected X-ray emission coming from a compact source at the centre of the cluster, with a 0.1--2.4 keV luminosity of $2 \\times 10^{36}$~W. Our data does not distinguish between an AGN and galaxy-scale thermal emission (perhaps a central cooling-flow spike). The central electron density (excluding the central source) is $n_{0} \\approx 1.0 \\times 10^{4}$ m$^{-3}$.  The central cooling time is in the range 2--6 Gyr, depending on whether or not the cool gas phase is concentrated in the centre. The cooling radius is 90 kpc. The cluster core should therefore be a cooling flow if we ignore the effect of the DRAGN. Indeed, most of the time, the DRAGN will be absent, as its life-time is certainly short compared to the age of the Universe. But at present the DRAGN, which is several times larger than the cooling radius, is almost certainly depositing more energy into this region of the ICM than is being lost radiatively.  Our best combined (PSPC and HRI) fit to the surface brightness profile, with a model consisting of a PSF, a $\\beta$ model and a background, gives a $\\beta$ parameter of $0.74 \\pm 0.03$, typical for clusters of galaxies. The radio lobes are largely positioned beyond the X-ray core radius, allowing for projection, so they are expanding essentially into a power-law atmosphere with density falling as $r^{-3\\beta} \\sim r^{-2.22}$, quite close to the $r^{-2}$ profile needed to give the lobes a self-similar structure \\citep{Falle1991}.  The thermal pressure at the deprojected distance of the radio lobes is an order of magnitude larger than the minimum pressure of the lobes. Thus the minimum energy in the lobes is a severe underestimate of the actual energy content.  This seems to be typical for DRAGNs, of both FR types, although only in a few cases can it be established as clearly as in Her~A.  Two features of the X-ray emission may result from the interaction of the cluster gas with the radio lobes. The region around the core of the cluster is clearly elongated along the radio axis, as revealed by our on- and off-lobe profiles. In addition there are a number of apparently discrete X-ray enhancements projected on or around the radio lobes; however some of these may be artefacts or due to poor photon statistics in our ROSAT data. Unlike some other well-studied DRAGNs, Her~A as yet does not show X-ray holes coincident with the radio lobes. But we have shown that the such holes could be too shallow to detect in our data, due to projection effects: the main lobes being displaced along the line of sight away from the X-ray bright region around the cluster core. Deep observations with {\\em Chandra} should reveal some sign of these holes, if present, and also clarify the nature of the `enhancements'."
    },
    "0402/astro-ph0402558_arXiv.txt": {
        "abstract": "{ Measurements of average pulse profiles made with a single linear polarization over the range 41--112 MHz are presented for PSR B0950+08. We show that the observed variable structure of the pulse profiles is a result of Faraday sinusoidal modulation of the pulse intensity with frequency. The rotation measure corresponding to this effect, $RM \\approx 4\\ {\\mathrm{rad\\, m^{-2}}}$, is about 3 times greater than the published value of $RM = 1.35\\ {\\mathrm{rad\\, m^{-2}}}$ (Taylor et al. \\cite{taylor}). ",
        "introduction": "PSR B0950+08 is well-studied over a wide frequency range from 24 to 10500 MHz. It exhibits a single pulse profile at frequencies above 400 MHz and a double pulse profile at low frequencies in the range 24--112 MHz (Hankins et al. \\cite{hankins}; Kuzmin et al. \\cite{kuzmin}). There is an interpulse occurring approximately $152{\\degr}$ ahead of the main pulse and a bridge of emission between the interpulse and main pulse (Lyne \\& Rickett \\cite{lyne1}).  The study of average pulse profiles from a large number of pulsars has shown that a pulse profile obtained by averaging several hundred individual pulses is very stable for most pulsars (Helfand et al. \\cite{helfand}; Rathnasree \\& Rankin \\cite{rathnas}). Profile changes can be caused by such phenomena as mode changing and nulling which occur in a small fraction of pulsars. The narrowband profile changes which are observed for the pulsar B0950+08 in the range 41--112 MHz and discussed in this paper are consistent with the effects of Faraday rotation on a linearly polarized pulse emission received by a linearly polarized antenna. Pulsars with considerable linear polarization observed with a single linear polarization show frequency-dependent profiles because of Faraday rotation of the plane of polarization in the interstellar medium. Pulse profiles obtained for different epochs show time-dependent profile shapes caused by a variation of the electron density or magnetic field in the propagation path, mainly due to varying ionospheric contribution to the rotation measure. At the Pushchino Radio Astronomy Observatory (PRAO), the effect of Faraday modulation of the pulse intensity with frequency at the output of the multi-channel receiver is used for the measurements of linear polarization  characteristics of pulsars and the estimation of their rotation measure (Vitkevich \\& Shitov \\cite{vitkev}; Shitov \\cite{shitov1}; Suleymanova \\cite{suleym}).  Changes of the average pulse profile at meter wavelengths for the pulsar B0950+08 were first noticed by Smirnova \\& Shabanova (\\cite{smirnova}). They found that the pulse profile is variable in time and shows narrowband changes with frequency. The explanation of the observed phenomenon by Faraday rotation was problematic because the predicted effect of the interstellar Faraday rotation for the tabulated value of $RM=1.35\\ {\\mathrm{rad\\, m^{-2}}}$ (Taylor et al. \\cite{taylor}) was negligible across the receiver bandpass at 102 MHz. The authors supposed that the narrowband changes of the pulse profile may be due either to the pulsar's intrinsic narrowband emission, manifesting only at low frequencies, or to scintillations of spatially separate sources. However, the low time resolution of their observations (about 10$\\degr$ of longitude) made a detailed analysis of this phenomenon impracticable.  The main purpose of this paper is to explain the observed narrowband profile changes of PSR B0950+08. The techniques used for observations over the frequency range 41--112 MHz  are described in \\S2. In \\S3 a large set of average profiles is investigated with respect to the shape changes with time. In \\S4 we study the narrowband changes of the pulse profile across the 2.56-MHz bandpass and establish a relation between the shape changes with time and the shape changes with frequency. \\S5 discusses the frequency dependence of the pulse profile after removing the effect caused by interstellar scintillation. In \\S6 we compute Faraday rotation effects using a numerical polarization model of the pulse profile. \\S7 describes the narrowband changes of the average profile at the lower frequencies of 88, 62 and 41 MHz. \\S8 presents the results of the timing data analysis for the pulsar. In \\S9 we conclude that the profile changes are due to the Faraday rotation effect. ",
        "conclusions": "We have found that the average pulse profile for PSR B0950+08 observed with a single linear polarization is frequency variable all over the range 41--112 MHz. The narrowband structure of the pulse shape is associated with sinusoidal modulation of the pulsar emission. This process is frequency dependent. According to our measurements, the interval of modulation at different observing frequencies is about 3--9 MHz at 111.87 MHz, 3--4 MHz at 88.57 MHz, 1--1.4 MHz at 62.15 MHz and 0.19--0.21 MHz at 41.07 MHz. Hence, the dependence of the modulation interval on frequency is well described by a power low with index close to 3 in the range  41--112 MHz. A similar frequency structure should be observed in the presence of the Faraday rotation effect. The presence of Faraday rotation is also confirmed by a numerical simulation of the polarization effects for the pulse profiles at 112 MHz.  The frequency interval of Faraday rotation of the plane of polarization of linear polarized emission ${\\Delta}F_{\\mathrm{\\pi}}$ at the observing frequency ${\\nu}$ is determined by rotation in the position angle ${\\theta}$ through  ${\\pi}$ and is calculated from  \\begin{equation} {\\Delta}F_{\\mathrm{\\pi}} = \\frac{{\\pi}\\, {\\nu}^{3}}{2\\, RM\\, c^{2}} = 17.48\\, {\\nu}^{3}\\, / \\, RM \\label{fpi} \\, , \\end{equation}  where ${\\Delta}F_{\\mathrm{\\pi}}$ is in MHz, ${\\nu}$ is in hundreds of MHz, the rotation measure $RM$ is in ${\\mathrm{rad\\, m^{-2}}}$ and $c$ is the light velocity.  Using relation~(\\ref{fpi}), we can estimate the rotation measure $RM$ corresponding to our measurements. For the mean values of ${\\Delta}F_{\\mathrm{\\pi}}$ equal to 6, 3, 1.2 and 0.21 MHz at the frequencies of 111.87, 88.57, 62.15 and 41.07 MHz respectively, the rotation measure $RM$ will be approximately 4 ${\\mathrm{rad\\, m^{-2}}}$. Taking into account variability of ${\\Delta}F_{\\mathrm{\\pi}}$ from one observation to another, we derive the values of $RM$ in the range 3--6 ${\\mathrm{rad\\, m^{-2}}}$. The derived values of $RM$ are very large compared to the rotation measure $RM = 1.35\\ {\\mathrm{rad\\, m^{-2}}}$ measured for PSR B0950+08 at frequencies above 400 MHz (Hamilton \\& Lyne \\cite{hamilton}; Taylor et al. \\cite{taylor}).  Thus, we have shown that the observed narrowband profile changes of the pulsar B0950+08 are a result of the Faraday rotation effect and could be explained as an artifact of observations of linearly polarized pulsar emission with a linearly polarized antenna. We have obtained a new value of the rotation measure corresponding to this effect, $RM \\approx 4\\ {\\mathrm{rad\\, m^{-2}}}$. This leads us to conclude that the published value of $RM$ is incorrect and is actually 3 times greater."
    },
    "0402/astro-ph0402302_arXiv.txt": {
        "abstract": "Weak gravitational lensing can be used to directly measure the mass along a line-of-sight without any dependence on the dynamical state of the mass, and thus can be used to measure the masses of clusters even if they are not relaxed.  One common technique used to measure cluster masses is fitting azimuthally-averaged gravitational shear profiles with a spherical mass model.  In this paper we quantify how asphericity and projected substructure in clusters can affect the virial mass and concentration measured with this technique by simulating weak lensing observations on 30 independent lines-of-sights through each of four high-resolution N-body cluster simulations.  We find that the variations in the measured virial mass and concentration are of a size similar to the error expected in ideal weak lensing observations and are correlated, but that the virial mass and concentration of the mean shear profile agree well with that measured in three dimensional models of the clusters.  The dominant effect causing the variations is the proximity of the line-of-sight to the major axis of the 3-D cluster mass distribution, with projected substructure only causing minor perturbations in the measured concentration.  Finally we find that the best-fit ``universal'' CDM models used to fit the shear profiles over-predict the surface density of the clusters due to the cluster mass density falling off faster than the $r^{-3}$ model assumption. ",
        "introduction": "Weak gravitational lensing, in which mass in a field is measured by the distortion induced in the shapes of background galaxies, has proven to be a powerful tool in the study of clusters \\citep[see reviews by][]{BA01.1,ME99.1}. With the advent of wide-field, multi-chip CCD cameras, weak lensing shear profiles for clusters have been measured to beyond the virial radius \\citep{CL01.1,CL02.1,DA02.1} with a high signal-to-noise. For many of these clusters, however, there is a disagreement in the cluster mass as measured by weak lensing and by strong lensing, X-ray observations, and cluster galaxy velocity dispersions.  One possible origin for the differences in the mass estimates is the error introduced by fitting spherically symmetric mass models to aspherical structures.  \\citet{PI03.1} have calculated that imposing a spherical model on a smooth tri-axial cluster can change the ratio of the measured X-ray mass to weak lensing mass by up to $30\\%$.  \\citet{KI01.1} investigated the effect of small-scale substructure seen in N-body simulations of clusters and concluded that the departures from a smooth mass model caused by these substructures do not greatly effect the weak lensing mass measurements.  Another possible origin for the mass estimate differences is the projection of mass structures along the line-of-sight onto the cluster mass in the weak lensing measurements.  Foreground and background structures, for which there exists no positional correlation with the cluster, do not produce a bias in the weak lensing measurements \\citep{HO03.1}.  Filamentary structures extending from the cluster along the line-of-sight can potentially cause an overestimate of the cluster mass from weak lensing, with estimates of the additional mass measured in N-body simulations ranging from a few percent \\citep{CE97.1, RE99.1} to over $50\\%$ \\citep{ME01.1}.  However, these results are obtained by comparing the total mass projected in a cylinder to that contained in a sphere in the N-body simulation, and not by fitting the shear produced by the projected mass with a projected mass model, as is most commonly done for weak lensing mass determinations of clusters.  In this paper we use four high-resolution N-body simulations of massive clusters to study the effects of cluster asphericity, secondary halos, and filamentary structures on the mass profiles measured by weak lensing. In Section 2 we present the simulations and methods used to project the 3-dimensional simulations to 2-dimensional mass maps.  We discuss the weak lensing techniques and results in Section 3, and present our conclusions in Section 4.  Throughout this paper we assume the cosmology of the simulations ($\\Omega_\\mathrm{m} = 0.3$, $\\Omega_{\\Lambda} = 0.7$, $H_0 = 70$ km/s/Mpc, spectral shape $\\Gamma=0.21$, and spectral normalization $\\sigma_8=0.9$). ",
        "conclusions": "While the best-fit NFW models did, on average, provide a good estimate of the cluster virial mass and mass profile within $r_{200}$, they overestimated the surface mass density in the projections for radii as small as half the virial radius.  This is due to the mass density at large radii falling faster than the $\\propto r^{-3}$ predicted by the NFW profile, and while the mass at radii larger than the virial radius are not considered in 3-D models, they constitute a significant fraction of the surface density of the clusters at radii smaller than the virial radius. The greatest discrepancy is seen in the merging cluster system g72, which has unusually low values of the concentration, and therefore a greater overestimate of the mass at large radius.  The cluster for which the best-fit of the mean shear field does not have a similar $r_{200}$ to the 3-D fit is g51, in which the 2-D fit has a $\\sim 2\\%$ higher value of $r_{200}$ than the 3-D fit.  This cluster has a high ellipticity without any significant substructure at large radius and is in virial equilibrium.  Under the tests of \\citet{JI00.1}, this cluster is one in which the NFW profile should provide a measurement of the total mass of the system.  However, as a result of the high ellipticity, combined with the under-density of mass at large radius in these simulations, the mass density decreases with increasing radius much more rapidly along the minor and intermediate axes of the cluster than is assumed by the NFW model.  As such, the sphere enclosed by the 3-D $r_{200}$ includes a large volume with densities far below those predicted by the NFW model.  Therefore, the $r_{200}$ measured in the 3-D mass distribution underestimates the amount of mass in the cluster, and mean value of the 2-D profiles is a more accurate measurement of the cluster virial mass.  The apparent paradox of the best fit NFW profiles to the shear profiles providing the correct virial masses of the clusters while over-predicting the surface densities, from which the shear profiles are calculated, is a result of the mass sheet degeneracy.  As can be seen in Eq.~\\ref{eq5}, $\\kappa$ profiles related by \\begin{equation} \\kappa ^{\\prime}(r) = \\kappa (r) + (1 - \\kappa (r))\\times \\lambda, \\label{eq9} \\end{equation} for any constant $\\lambda$, produce the same shear profile.  This is due to the shear profiles measuring the change in mass with radius, and therefore weak lensing only being able to measure the mass relative to the density at the outer radius of the measured shear region.  The imposition of a chosen mass model breaks the mass sheet degeneracy, provided the model has a bijective relation between the shape of the surface density profile and the total mass at a given radius.  There is nothing, however, which prevents an incorrect model from being assumed, and therefore measuring a mass profile which differs from the true profile by some value of $\\lambda$ via Eq.~\\ref{eq9}. In this case, while the NFW model does provide a relation between the surface density slope, which is measured by the shear profile, and the total mass, which is not, at the outer edge of the shear profile, the mass density assumed at large radii is incorrect, and the best fit models over predict the total surface density within the fitting region.  Our result that the mass measured by weak lensing observations is affected mostly by the alignment of the line-of-sight to the major-axis of the cluster, and therefore the small-scale substructure is of minimal importance, is in good agreement with the results of \\citet{KI01.1}.  This suggests that for the purposes of modeling weak lensing observations, clusters can be adequately described by a smooth, tri-axial mass distribution.  Massive sub-halos projected onto the cluster core can, however, perturb the measured concentration parameter for the cluster.  The levels of the perturbations of the concentration were typically $\\sim 10-20\\%$, except in the case of the major merger cluster g72, in which case the perturbations were on the order of $50-100\\%$.  Infalling haloes which are outside the virialized region of the cluster can also increase the measured $r_{200}$, as occurred for one projection of g51, but such projections should be detected in redshift surveys \\citep[e.g.][]{CZ02.1}.  \\begin{figure} \\centering \\resizebox{\\hsize}{!}{\\includegraphics{fig10.ps}} \\caption{Plotted above are the amounts of excess surface mass within $r < 3 \\mathrm{Mpc}$ compared to an annular region of $3 \\mathrm{Mpc} < r < 4 \\mathrm{Mpc}$ in projections of a 128 Mpc box with the inner 13 Mpc removed.} \\label{fig10} \\end{figure}  Our results that we find, on average, the correct $r_{200}$, and therefore the correct virial mass, for the clusters is in stark contrast to the results of \\citet{CE97.1} and \\citet{ME01.1}, who found that weak lensing measurements would be consistently higher than the virial masses of the clusters.  The two major differences between our study and theirs are the technique used to measure mass via weak lensing and the size of the projected line-of-sight through the clusters.  Both \\citet{CE97.1} and \\citet{ME01.1} simulated weak lensing observations by calculating the aperture densitometry statistic \\begin{equation} \\zeta = \\bar{\\kappa}(r<r_{\\mathrm{ap}}) - \\bar{\\kappa}(r_{\\mathrm{ap}}<r<r_{\\mathrm{max}}) \\end{equation} which is the mean $\\kappa$ within some radius $r_{\\mathrm{ap}}$ minus the mean $\\kappa$ in an annular region from $r_{\\mathrm{ap}}$ to $r_{\\mathrm{max}}$ \\citep{FA94.1}, which can be measured by convolving the shear profile with a specific kernel.  Rather than calculating the shear produced by the clusters, both papers estimated $\\zeta$ directly from the projected images by measuring the $\\bar{\\kappa}$ within a given radius. However, instead of subtracting off the mean $\\kappa$ within an annular region immediately surrounding the radius used to measure $\\bar{\\kappa}$, they subtracted off the mean $\\kappa$ measured across all of the simulations. Given that the surface density is still decreasing with increasing radius at $r = r_{200}$, the mean $\\kappa$ in an annular region immediately outside the $r = r_{200}$ aperture is higher than the mean surface density in the rest of the simulated regions.  As a result, both papers overestimated $\\zeta$ for their clusters, and therefore overestimated the cluster mass which would have been measured by this weak lensing technique.  Both papers then compared the $\\zeta$ values, converted into a surface mass, with cluster models.  With this method, the chosen cluster model must have the correct mass profile at all radii, as one must integrate along the line of sight to measure the surface mass density, in order make an accurate comparison. The 5-10$\\%$ overestimate of the cluster mass from this technique in \\citet{CE97.1} is comparable in size, although opposite in sign, to the $\\sim 5\\%$ difference between surface mass at $r_{200}$ of the the NFW profiles used here and the true surface mass of the simulation. The larger overestimate of cluster mass via weak lensing from \\citet{ME01.1} is a result of their choice of a cluster model, namely a spherical model which has a mean density of $200\\times \\rho_{\\mathrm{crit}}$ inside a sphere with radius $r_{200}$, and a zero density outside the sphere.  The excess mass which is detected is therefore likely to be that associated with the cluster outside of $r_{200}$.  Indeed, as can be seen in Fig.~\\ref{fig9}, if \\citet{ME01.1} had integrated an NFW profile to the edge of their simulated regions instead of only out to $r_{200}$ to measure the expected surface mass of the clusters, they would have found that the technique was underestimating, instead of overestimating, the cluster mass.  While we project the cluster in a 13 Mpc box, the projected regions of \\citet{CE97.1} (64 Mpc box) and \\citet{ME01.1} (128 Mpc sphere) are much larger.  If there is a structure along the line of sight, such as a filament, then the surface mass of the field will increase, and such structures might be more common in the regions around massive clusters.  Indeed, this is one of the reasons quoted in \\citet{ME01.1} for why they find a greater over-estimate of the mass via weak lensing than was found in \\citet{CE97.1}.  In order to test if the 13 Mpc box size causes us to underestimate the mass which would be measured with our weak lensing technique, we selected three clusters from the original simulation the re-simulations were based on.  For each cluster, we rotated the particles to obtain 10 independent lines-of-sight, cut out a 128 Mpc box around the cluster, and projected along the three sides to give 30 independent projections per cluster.  We then cut out a 13 Mpc box around each cluster, performed the same projections, and subtracted them from the 128 Mpc box projections to obtain the projection of only those structures which would be located outside of the high-resolution region used in the NFW fits. For each projection we then calculated the $\\zeta$ statistic with $r_{\\mathrm{ap}} = 3$ Mpc, the outer radius of our shear fitting region, and $r_{\\mathrm{max}} = 4$ Mpc.  These values of $\\zeta$, converted into a surface mass by multiplying by $\\pi r_{\\mathrm{ap}}^2 \\Sigma _{\\mathrm{crit}}$, are shown as a histogram in Fig.~\\ref{fig10}.  The mean of the $\\zeta$ distribution is consistent with zero, although with a skew that results in a broader tail to the large positive surface masses.  The distribution has a rms of $1.2\\times 10^{14} \\mathrm{M}_\\odot$, which is $\\sim 5\\%$ of the surface mass at this radius for the simulated clusters.  The two projections with high values for the increased mass within $r_{\\mathrm{ap}}$ are both the result of a second, smaller cluster along the line-of-sight.  Thus, expanding the projected region from our original 13 Mpc to 128 Mpc would have only resulted in a small scatter being added to the measured cluster masses, with the exception of the occasional projection which would have a significantly higher mass (up to $25\\%$ additional mass, although these simulations were chosen to not have similar mass neighbors so exclude the possibility of two equal mass clusters being projected along the line of sight).  These projections with higher projected mass, however, are caused primarily by secondary clusters along the line-of-sight, and should be visible in the form of a concentration of galaxies at redshifts slightly different from the main cluster."
    },
    "0402/astro-ph0402628_arXiv.txt": {
        "abstract": "The study of the hard ($E\\gtrsim 10$~keV) energy spectra of X--ray binary pulsars can give a wealth of information on the physical processes that occur close to the neutron star surface. Extreme matter regimes are probed, and precious information on how matter and radiation behave and interact in critical conditions can be obtained. We will give an overview on the most recent results obtained by RXTE and BeppoSAX on this class of objects, in order to pass the baton onto just launched experiments, like INTEGRAL, or soon to be launched, like AGILE and ASTRO-E2. ",
        "introduction": "The discovery of X--ray emission from celestial objects further extended the knowledge of our Universe, which resulted very different from the quiet and calm Universe the astronomers of the last centuries described. Indeed, it was just at the beginning of the 1960's that, with the advent of stratospheric balloons and rockets, it was possible to launch outside our atmosphere some Geiger counters (Giacconi et~al.\\ 1962). In this way, the first discrete X--ray source in our Galaxy was discovered: Sco X--1. In 1966 the first optical counterpart to a galactic X--ray source, Sco X--1, was identified with an old 12$^{\\rm th}$--13$^{\\rm th}$ magnitude star (Sandage et~al.\\ 1966). In the following years a theoretical model was developed according to which galactic X--ray sources are close, interacting binary systems composed by a ``normal'' star and a compact object (Shklovskii 1967). In the same years it was understood that spherical accretion could become non symmetric, leading to the formation of an accretion disk around the compact object, if the accreting matter possesses enough angular momentum (Prendergast \\& Burbidge 1968).  But the decisive step toward the comprehension of this class of objects was achieved with the discovery of pulsed emission from some X--ray sources (Giacconi et~al.\\ 1971; Tananbaum et~al.\\ 1972).  Indeed, the variability on short time scale --- for example the first discovered X--ray binary pulsar, Cen X--3, spins at about 4.8~s --- implies a small emitting region. Furthermore, because the system is not destroyed by the centrifugal force it is necessary that at the surface of the emitting object the gravitational force is greater than the centrifugal one. This implies $\\Omega_p \\gtrsim \\sqrt{G\\langle\\rho\\rangle}$, where $\\Omega_p$ is the spin frequency, $G$ the gravitational constant and $\\langle\\rho\\rangle$ the object mean density. The observed values of $\\Omega_p$ imply $\\langle\\rho\\rangle \\gtrsim 10^6$ g/cm$^3$, and therefore the compact nature of the object responsible of the pulsed X--ray emission was established.  The only compact object able to explain all the phenomena observed in X--ray pulsars, as pulse period range and surface magnetic field strength, is a neutron star (see, e.g., Shapiro \\& Teukolsky 1983).  The standard model explains the X--ray emission as due to the conversion of the kinetic energy of the in-falling matter (coming from the intense stellar wind of an early optical star, in this case we speak about wind-fed binaries, or from an accretion disk due to Roche-lobe overflow, and this is the case of disk-fed binaries) into radiation, because of the interactions with the strong magnetic field of the neutron star, of the order of $10^{11}$--$10^{13}$~G\\footnote{Obtained from conservation of magnetic flux during the process of collapse from a ``normal'' star ($B\\sim 10-100$~G, $R\\sim 10^6$~Km) to a neutron star ($R\\sim 10$~Km) and lately confirmed by the observation of cyclotron resonance features (see below).}. The dipolar magnetic field of the compact object drives the accreted matter onto the magnetic polar caps, and if the magnetic field axis is not aligned with the spin axis then the compact object acts as a ``lighthouse'', giving rise to pulsed emission when the beam (or the beams, according to the geometry) crosses our line of sight.  Subsequent observations clearly demonstrated the binary nature of these objects by observations of X--ray eclipses and Doppler delays in the pulse arrival times (Schreirer et~al.\\ 1972). From these measurements it was possible to ``solve'' the binary systems, obtaining masses in agreement with that expected for a neutron star. But for some X--ray pulsars, with pulse periods in the 5--12~s range, every attempt to find signatures of binary motion was unsuccessful. It was later recognized that they form a class of their own, the so-called anomalous X--ray pulsars, in which their X--ray emission is due to the conversion of the magnetic field energy into radiation (Thompson \\& Duncan 1995; 1996).  We will not discuss here on this quite interesting class (see, e.g., Mereghetti et~al.\\ 2002 for a review). ",
        "conclusions": "The study of X--ray binary pulsars, especially in the hard X--ray band, received a new momentum from the results by BeppoSAX and RXTE. Besides the direct measurement of the neutron star magnetic field strength, the observation of CRFs can give hints on the physical processes occurring in extreme condition of temperature, density, gravity and magnetic field.  It will be up to new missions (INTEGRAL, AGILE, ASTRO-E2) to find the answers to the issues that were raised by their predecessors."
    },
    "0402/astro-ph0402134_arXiv.txt": {
        "abstract": "I partially review our current knowledge on microquasars, making special emphasis on radio interferometric observations, and I comment on emerging trends in this field of astrophysics. ",
        "introduction": "An X-ray binary is a binary system containing a compact object, either a neutron star or a stellar-mass black hole accreting matter from the companion star. The accreted matter carries angular momentum on its way to the compact object, around which, in most cases, forms an accretion disk responsible for the X-ray emission. The most recent catalogs contain a total of 280 X-ray binaries (Liu, van Paradijs, \\& van den Heuvel 2000; Liu, van Paradijs, \\& van den Heuvel 2001). In 131 of such systems the companion star has a spectral type O or B, and they are classified as High Mass X-ray Binaries (HMXBs), where mass transfer takes place via a decretion disk (in Be stars) or via a strong stellar wind or Roche-lobe overflow (in OB supergiants). In 149 X-ray binaries the optical companion has a spectral type later than B, and they are called Low Mass X-ray Binaries (LMXBs), where mass transfer occurs via Roche-lobe overflow.  Among them, this author has found 43 Radio Emitting X-ray Binaries (REXBs), of which 8 are HMXBs that usually show persistent radio emission and 35 are transient LMXBs. Since the detected radio emission (in nearly all cases) displays nonthermal spectra, shows some degree of polarization, and implies high brightness temperatures, it is interpreted as produced by the synchrotron radiation mechanism, which takes place when we have charged particles accelerated in the presence of magnetic fields.  A microquasar is simply a REXB displaying relativistic radio jets that can be imaged at a variety of angular scales using different interferometers. Due to the lack of space, I refer the reader to Mirabel \\& Rodr\\'{\\i}guez (1999) and to Fender (2004) for extended reviews on the topic (including discussion on emission mechanisms, relativistic effects, observed radio/X-ray correlations, energetics, jet formation, etc.). The name microquasar was given not only because of the observed morphological similarities between these sources and the distant quasars but also because of physical similarities, since when the compact object is a black hole, some parameters appear to scale with the mass of the central object. In this context, the temperatures of the inner parts of the accretion disks are of the order of $\\sim$$10^7$~K in the case of microquasars containing stellar-mass black holes and $\\sim$$10^5$~K in the case of quasars containing supermassive black holes ($10^6$--$10^9$~$M_{\\sun}$). This explains why in microquasars the accretion luminosity is radiated in X-rays, while it is done in the optical/UV domain in the case of quasars, and why we had to wait until the era of high-energy astrophysics to discover these sources. On the other hand, the characteristic jet sizes seem to be proportional to the mass of the black hole, since radio jets in microquasars have typical sizes of the order of light years, while radio jets in quasars reach distances up to several million light years in giant radio galaxies. Last, but not least, the timescales are also directly scaled with the mass of the black hole following $\\tau\\simeq R_{\\rm Sch}/c=2GM_{\\rm X}/c^3\\propto M$, being $R_{\\rm Sch}$ the Schwarzschild radius. Therefore, phenomena that take place in timescales of years in quasars can be studied in minutes in microquasars. In this sense, one can say that microquasars mimic, on smaller scales, many of the phenomena seen in AGNs and quasars, but allow a better and faster progress to understand the accretion/ejection processes that take place near compact objects. However, it should be noted that the angular scales in units of $R_{\\rm Sch}$ are much higher for nearby quasars than for microquasars, allowing a {\\it closer} look to the jet formation region.  The number of currently known microquasars is around 16 (Rib\\'o 2002), among the 43 catalogued REXBs. It is interesting to note that some authors (e.g., Fender 2004) have proposed that all REXBs are microquasars (i.e., radio emission is always in the form of relativistic jets), and would be detected as such provided that there is enough sensitivity and/or resolution in the radio observations (some of them conducted in the past). In fact, we have detected radio jets in all REXBs for which detailed observations have been possible, with the only exception of CI~Cam. Moreover, it has been suggested that all X-ray binaries, except X-ray pulsars that disrupt the inner part of the accretion disk where the jets are supposed to be launched, have relativistic jets (Fender 2004). In this context, all of them should be considered microquasars. Although this is a reasonable hypothesis, only deep radio observations of the faint X-ray binary population will allow to confirm or reject it. The planned improvement of existing radio interferometers will definitely help to solve this important issue. ",
        "conclusions": "Microquasars allow to gain insight into jet physics (including formation, matter content, propagation and interaction with the ISM), and on the origin and evolution of black holes. They could also be sources of high-energy $\\gamma$-rays and TeV neutrinos, and eventually be ULXs. There are still lots of open questions that hopefully we will be able to answer with the advent of new instrumentation."
    },
    "0402/astro-ph0402378_arXiv.txt": {
        "abstract": "We present a statistical analysis of the X-ray bursts observed from the 2002 June 18 outburst of the Anomalous X-ray Pulsar (AXP) \\tfn, observed with the Proportional Counter Array aboard the \\textit{Rossi X-ray Timing Explorer}.  We show that the properties of these bursts are similar to those of Soft Gamma-Repeaters (SGRs).  The similarities we find are: the burst durations follow a log-normal distribution which peaks at 99~ms, the differential burst fluence distribution is well described by a power law of index $-1.7$, the burst fluences are positively correlated with the burst durations, the distribution of waiting times is well described by a log-normal distribution of mean 47~s, and the bursts are generally asymmetric with faster rise than fall times.  However, we find several quantitative differences between the AXP and SGR bursts.  Specifically, there is a correlation of burst phase with pulsed intensity, the AXP bursts we observed exhibit a wider range of durations,  the correlation between burst fluence and duration is flatter than for SGRs, the observed AXP bursts are on average less energetic than observed SGR bursts, and the more energetic AXP bursts have the hardest spectra -- the opposite of what is seen for SGRs. We conclude that the bursts are sufficiently similar that AXPs and SGRs can be considered united as a source class yet there are some interesting differences that may help determine what physically  differentiates the two closely related manifestations of neutron stars. ",
        "introduction": "Soft gamma repeaters (SGRs) are an exotic class of Galactic sources that are now commonly accepted as being magnetars -- isolated, young neutron stars that are powered by the decay of an ultra-high magnetic field.  The evidence for high surface fields ($\\sim 10^{14} - 10^{15}$~G) comes from several independent lines of reasoning \\citep{dt92a,pac92,td95,td96a}.  These include: the high dipolar magnetic fields implied by the spin properties of SGRs seen in quiescence  under the assumption of magnetic dipole braking \\citep{kds+98,ksh+99}; the requirement of a magnetar-strength field to confine the energy released in the tails of hyper-Eddington outbursts seen from two SGRs \\citep{mgi+79,hcm+99}; the requirement of a high field to allow the decay rate necessary to power the burst and persistent emission \\citep{td96a, gr92}; and the magnetic suppression of the Thomson cross-section, which allows hyper-Eddington bursts to be observed \\citep{pac92}. For a review of SGRs, see Kouveliotou~et~al.~[these proceedings].    Anomalous X-ray pulsars (AXPs), another exotic class of Galactic neutron stars, have also been suggested to be magnetars \\citep{td96a}. This is because of their anomalously bright X-ray emission which can be explained neither by conventional binary accretion models nor rotation power \\citep{ms95}. Also, their spin parameters, as for SGRs, imply large magnetic fields under standard assumptions of magnetic braking.  They also have similar, though on average softer, X-ray spectra compared with those of SGRs in quiescence. However, unlike SGRs, in the $>20$~yr since the discovery of the first AXP \\citep{fg81}, none was seen to exhibit SGR-like bursts.  For this reason, alternative models involving unconventional accretion scenarios have been proposed to explain AXP emission \\citep{vtv95,chn00,alp01}. See Kaspi~et~al.~[these proceedings] for a review of AXPs.  \\begin{figure} \\includegraphics[width=.36\\textwidth]{gavriilf_f1.eps} \\caption{Distribution of the pulse phases of \\tfn\\ which correspond to the times of the burst peaks (solid points). The solid curve is the folded 2--60~keV light curve of the 2002 June 18 observation with the bursts omitted. \\label{fig:burst phases}} \\end{figure}  The magnetar model for AXPs was recently given a boost when SGR-like bursts were detected from two AXPs.  \\citet{gkw02} reported on the discovery of two X-ray bursts in observations obtained in the direction of AXP \\tfe.  The temporal and spectral properties of those bursts were similar only to those seen only in SGRs. However, the AXP could not be definitely identified as the burster. On 2002 June 18, a major outburst was detected unambiguously from AXP \\tfn, involving over 80 bursts as well as significant spectral and timing changes in the persistent emission \\citep{kgw+03}.  Those bursts demonstrated that AXPs are capable of exhibiting behavior observed, until now, uniquely in SGRs, therefore implying a clear connection between the two source classes.  Such a connection was predicted only by the magnetar model \\citep{td96a}. ",
        "conclusions": "The bursts we have observed for \\tfn\\ are clearly similar to those seen uniquely in SGRs.  As concluded by \\citet{gkw02} and \\citet{kgw+03}, AXPs and SGRs clearly share a common nature, as has been predicted by the magnetar model. The bursts seen in the 2002 June 18 outburst of \\tfn\\ are qualitatively similar to those seen in SGRs, and in many ways quantitatively similar. Specifically: \\begin{itemize} \\item the mean burst durations are similar \\item the differential burst fluence spectrum is well described by a power law of index $-1.7$, similar to those seen in SGRs (and earthquakes and solar flares) \\item burst fluences are positively correlated with burst durations \\item the distribution of and mean waiting times are similar \\item the burst morphologies are generally asymmetric, with rise times usually shorter than burst durations \\end{itemize} However, there are some interesting quantitative differences between the properties of the AXP and SGR bursts.  These may help shed light on the physical difference(s) between these classes.  The differences can be summarized as: \\begin{itemize} \\item there is a significant correlation of burst phase with pulsed intensity, unlike in SGRs \\item the AXP bursts have a wider range of burst duration (though this may be partly due to different analyses procedures) \\item the correlation of burst fluence with duration is flatter for AXPs than it is for SGRs (although when selection effects are considered, this correlation should really be seen as an upper envelope for AXPs and SGRs) \\item the fluences for the AXP bursts are generally smaller than are in observed SGR bursts \\item the more energetic AXP bursts have the hardest spectra, whereas for SGR bursts, they have the softest spectra \\item under reasonable assumptions, SGRs undergo outbursts much more frequently than do AXPs \\end{itemize}  Given the rarity of AXP bursts coupled with the unique information that detection of such bursts provides, observing more outbursts is obviously desirable.  Continued monitoring is thus clearly warranted, and \\xte\\ with its large area and flexible scheduling is the obvious instrument of choice.    \\begin{theacknowledgments} We are grateful to C.~Kouveliotou, M.~Lyutikov, S.~Ransom, M.S.E.~Roberts, D.~Smith, and C.~Thompson for useful discussions. This work was supported in part by NSERC, NATEQ, CIAR and NASA. This research has made use of data obtained through the High Energy Astrophysics Science Archive Research Center Online Service, provided by the NASA/Goddard Space Flight Center. \\end{theacknowledgments}"
    },
    "0402/astro-ph0402652_arXiv.txt": {
        "abstract": "We derive the fraction of blue galaxies in a sample of clusters at $z < 0.11$ and the general field at the same redshift. The value of the blue fraction is observed to depend on the luminosity limit adopted, cluster-centric radius and, more generally, local galaxy density, but it does not depend on cluster properties. Changes in the blue fraction are due to variations in the relative proportions of red and blue galaxies but the star formation rate for these two galaxy groups remains unchanged. Our results are most consistent with a model where the star formation rate declines rapidly and the blue galaxies tend to be dwarfs and do not favour mechanisms where the Butcher-Oemler effect is caused by processes specific to the cluster environment. ",
        "introduction": "The observation by \\cite{bo84} that clusters of galaxies contain a larger fraction of blue galaxies at progressively higher redshift (the Butcher-Oemler effect) contributed to the establishment of the current view of clusters as sites of active galaxy evolution, in which environmental influences dramatically alter the morphologies and star formation histories of their members.  However, as the number of clusters observed has increased, it has become clear that the Butcher-Oemler effect is not solely or simply an evolutionary trend. The large scatter observed in the blue fractions for clusters in narrow redshift ranges \\citep{sma98,mar00,got03} implies the existence of environmental effects, which may compete with, and possibly mimic, evolutionary trends, if the cluster samples evolve with redshift.  Studies of the original Butcher-Oemler sample lend support to this scenario: \\cite{nkb88} measured the surface densities and velocity dispersions of seven Butcher-Oemler clusters and found that the high and low redshift sample differ. \\cite{ae99} show that the X-ray luminosity of  high redshift Butcher-Oemler clusters is higher than that of low redshift clusters, although the X-ray luminosity  function of clusters does not evolve to $z \\sim 0.8$ \\citep{hol02}, arguing that the high redshift sample does not represent the progenitors of modern-day clusters. Similarly, both \\cite{sma98} and \\cite{fai02} find a low blue fraction in their samples of X-ray selected clusters.  The blue fraction has also been observed to depend on a number of other factors, such as:  the luminosity limit used and the cluster centric distance \\citep{ell01,got03}, richness \\citep{mar01}, cluster concentration \\citep{bo84} and, possibly, the presence of substructure \\citep{mru00}. These findings point to the necessity of understanding environmental effects on the blue fraction in order to disentangle evolutionary trends from selection biases in studies of clusters at high redshift.  The aim of the present paper is to study how the fraction of blue galaxies varies as a function of cluster properties and in the field at the same redshift in the local universe, in order to analyse the effects of environment in isolation from the evolutionary trends that give rise to the Butcher-Oemler effect. \\cite{lew02} and \\cite{bal03} have recently discussed how the star formation rate, as measured from the H$\\alpha$ equivalent width, varies with environment within the 2dFGRS sample. However, colours provide a measure of the average star formation histories over longer timescales than H$\\alpha$ and may therefore better reflect the influence of environment on galaxy properties, especially if the mechanisms responsible for the Butcher-Oemler effect operate on longer time frames, as is the case for harassment \\citep{moo96}, for instance.  The structure of this paper is as follows: in the next section we present our analysis; selection of clusters, cluster members and a comparison field sample, definition of the radial and luminosity limits used, $k$-corrections, colour-magnitude relations, density measurements and a discussion of how the blue fraction and its error are measured. The main results are detailed in \\S 3. below and we discuss these in \\S 4.  Throughout this paper we adopt a cosmology with $\\Omega_M=0.3$, $\\Omega_{ \\Lambda}=0.7$. It is normal to define $h \\equiv H_0 / 100\\, {\\rm km\\ s^{-1}\\ Mpc^{-1}}$: here we suppress the $h$ scaling, so that the full explicit meaning of Mpc in length units is $h^{-1}$ Mpc; for absolute magnitudes $M$ stands for $M+5 \\log_{10} h$. ",
        "conclusions": "We have determined the fraction of blue galaxies in a sample of nearby clusters and in the general field at the same redshift. Although the blue fraction varies considerably from cluster to cluster, we find that the variation is a real property of the sample and not due to statistical noise.  The mean value of the blue fraction is $0.11 \\pm 0.01$ for galaxies within $r_{30}$ and $0.13 \\pm 0.01$ for galaxies within $r_{200}/2$. These are higher than the value of $0.03 \\pm 0.09$ from \\cite{mar00}, $\\sim 0.07$ from \\cite{mar01} and $\\sim 0.07$ from \\cite{pim02} at $z < 0.11$, although they are within the errors. One possibility for this discrepancy lies in the possibility of blue interlopers from the field population being included in our redshift sample \\citep{dia01}. We have calculated the level of contamination by integrating the field luminosity function of \\cite{mad02} over the appropriate luminosity range and within the pseudo-volume defined by the aperture used and the range in velocities spanned by cluster galaxies. The blue fraction for this sample was calculated from our density-dependent blue fraction. The average level of contamination is $0.08 \\pm 0.06$ but this depends critically on assumptions concerning the mean density of galaxies in cluster outskirts (which influence both the normalization and the fraction of blue interlopers as per Fig.~5). In addition, there are contributions to the error from both Poisson and clustering statistics. The number of blue galaxies is, in any case, small and therefore the level of contamination is uncertain and the error in its estimate quite large. We therefore ignore the issue of possible contamination of our spectroscopic sample of cluster members.   \\begin{itemize}  \\item A first conclusion to be drawn from the derived blue fractions is that few clusters have $f_b > 40\\%$ and none has $f_b > 60\\%$ .  Since this sample of clusters is, at least to first order, complete and volume limited, it represents a fair sampling of the range of blue fractions encountered in cluster environments. By contrast, samples of clusters at high redshift contain at least a few objects with blue fractions in excess of 40\\% (Fairley et al. 2002, La Barbera et al. 2003 and references therein). This would suggest that the observed evolution is real: contamination in our sample (if real) should only make the local blue  fraction lower, while high redshift clusters are drawn from the high richness envelope that contains few local clusters with high $f_b$. However, it is possible that optically selected clusters at high redshift tend to contain higher fractions of blue galaxies because these make them more conspicuous in blue plates.\\\\  \\item The blue fraction appears to depend on the luminosity limit and cluster centric radius used (Fig.~3), as previously found by \\cite{mar00} and \\cite{ell01}.  This behaviour is similar to that observed for dwarf galaxies in Coma and Abell 2218: \\cite{pra03} have shown that dwarfs have a steeper luminosity function at larger cluster centric radii and are preferentially found at large distances from the cluster core, with these trends being stronger for the faintest dwarfs. Similar trends may have been observed for galaxies in the Coma cluster \\citep{bej02} and appear to persist in the $z\\sim 0.4$ CNOC sample. To the extent that present-day clusters and their populations represent proxies for high redshift objects, these observations imply that a proportion of the blue galaxies are intrinsically low luminosity objects.\\\\  \\item There is no dependence of the blue fraction on cluster properties (Fig.~4). The lack of correlation of the blue fraction with cluster properties suggests that the star formation rate, or its decrease, is not related to the large scale structure: therefore, mechanisms which involve cluster-wide effects, such as ram stripping by the cluster gas, tides, which depend on the cluster mass, or harassment, whose efficiency is proportional to velocity dispersion, are disfavoured by the present study, as well as other explanations that depend on processes specific to the cluster environment. However, as \\cite{bal03} show, the star formation rate is affected by local processes and close interactions are a viable mechanism for the origin of the blue fraction.  The above is apparently in contrast  with some previous studies: \\cite{mru00} claim that clusters containing large amounts of substructure have higher blue fractions and interpret this in terms of a shock model during collisions of subclusters. However, their result is based on three clusters and may not be valid for the general population. \\cite{mar01} and \\cite{got03} suggest that the blue fraction depends on richness. One possibility is that, since they use a fixed 0.7 Mpc aperture, and given the radial dependence observed in Fig.~3, that a spurious richness dependence may be induced by aperture effects. We have derived blue fractions for a 0.7 Mpc aperture in our clusters and tested for correlation with richness but failed to detect a significant signal.  Similarly, we failed to detect any correlation with richness if we use the same luminosity range as \\cite{mar01}. However, we see in Fig.~4 that rich clusters tend not to have large blue fractions and we also observe that the blue fraction in the 0.7 Mpc aperture shows large scatter. If the sample used by \\cite{mar01} and \\cite{got03} is biased towards richer clusters at high redshift, the combination of these two effects may produce a spurious correlation.\\\\   \\item We observe that the blue fraction exhibits a strong dependence with local density in the general field (Fig.~5) and that the cluster value represents a continuation of this trend to higher density regimes.  This suggests that the same processes are responsible for the Butcher-Oemler effect in all environments and that these mechanisms vary smoothly as a function of density. Again, this implies that cluster-specific processes are not likely causes of changes in the blue fraction, but local mechanisms whose efficiency varies smoothly with density, such as interactions \\citep{lh88,cou98}, may be viable explanations.\\\\  \\item The relative star formation rate for the blue galaxies does not appear to decrease as a function of density (Fig.~6) or as a function of luminosity or radius in clusters (Fig.~7).  This implies that changes in the blue fraction are solely due to changes in the relative fractions of red (quiescent) and blue (star-forming) galaxies. The trends we observe can be roughly reproduced by using the density-dependent luminosity functions of \\cite{cro04} for early- and late-type galaxies and the type-dependent cluster luminosity functions of \\cite{dep03} and simply assuming that early-type galaxies are red and late types are blue. This suggests that the simple model presented in \\cite{dep03}, where galaxies simply moved between spectral types without number or luminosity evolution, provides, heuristically, a good representation of galaxy evolution.\\\\  One possible caveat is that, by selecting blue galaxies, we have automatically selected for galaxies with H$\\alpha$ emission, while morphologically selected samples show declining star formation rates as a function of density \\citep{gom03,got03b}. However, the observation that blue cluster galaxies have the same H$\\alpha$ equivalent width as their counterparts in the field is non-trivial, because lower (but non-zero) star formation rates will lead to weaker H$\\alpha$ emission, while the colour would remain blue by our definition. This is in contrast, for instance, with models where galaxies are `choked' in clusters \\citep{bal00}.  \\end{itemize}   The above argues for a model where star formation declines over relatively short timescales, leading to the bimodal distribution in H$\\alpha$ observed by \\cite{bal03}, and galaxy colours evolve quickly to the red envelope, producing the colour bimodality observed here and in the SDSS \\citep{hogg03,bla03}. As in \\cite{dep03} this is possible if the optical colour is dominated by the young population but the majority of light is provided by the underlying old stars, so that once the star formation is extinguished galaxy colours quickly evolve on to the passive locus. \\cite{shi02} have shown that truncating the star formation of spiral galaxies in the field leads them on to the colour-magnitude relation defined by the Coma E/S0 galaxies, but that this process is inefficient for the more massive galaxies and is viable only for low mass spirals. When we consider the radial and luminosity trends, one possible interpretation is that a large fraction of the blue galaxies are actually dwarfs undergoing episodes of star formation, as also suggested by a number of other lines of evidence \\citep{ros97, cou98,dep04}. However, we caution the reader that since our sample of cluster and field galaxies is local, we cannot properly discuss the origin of the blue fraction at higher redshift.  A complication in interpreting these data is that the observed correlations of $f_b$ and H$\\alpha$ with density on large scales imply that the environment {\\it today} is not affecting the properties of the population and therefore we are unable to observe galaxy evolution in progress but just its end result. This implies that identifying the mechanisms responsible for the Butcher-Oemler effect in the local universe is problematic. Surveys at intermediate redshifts (e.g. DEEP2, VIMOS) may be able to witness the main phases of galaxy evolution."
    },
    "0402/astro-ph0402187_arXiv.txt": {
        "abstract": " ",
        "introduction": "With the release of the Wilkinson Microwave Anisotropy Probe (WMAP) results (Spergel et~al. 2003), the Cold Dark Matter (CDM) has essentially shifted from a ``favoured'' paradigm to what is now referred to as the ``concordance model''.  Hierarchical clustering is an important component of CDM models, one in which the first objects to collapse in the Universe were small, with subsequent merging of these objects coupled with collapse on increasingly larger scales as the Universe ages.  Such merger- and accretion-driven evolution appears to have peaked over the redshift range $\\sim 2$ -- 5 (e.g. Murali et~al. 2002), but equally important, continues to the present-day.  Indeed, our own Local Group provides several spectacular examples of hierarchical clustering ``in action'', including the disrupting Sagittarius dwarf (Ibata et~al. 1994), the putative Canis Major dwarf (Martin et~al. 2003), and perhaps the most visually stunning of all, the debris associated with the interacting Large and Small Magellanic Clouds (LMC and SMC, respectively, hereafter) -- the so-called Magellanic Stream (Mathewson et~al 1974).  Disrupting satellites such as these are the best local laboratory to understand the physical processes of ``galactic cannibalism'', as we have the luxury of obtaining detailed observations pertaining to the respective systems' star formation histories (e.g. Harris~\\& Zaritsky 2001; Smecker-Hane \\etal\\ 2002) and internal chemical evolution via stellar abundance patterns for individual stars within the satellites (Tolstoy et~al. 2003).  One of the most obvious of manifestations of cannibalism within the Local Group is that of the aforementioned Magellanic Stream. The Magellanic Stream (MS) is a remarkably colinear band of (primarily) neutral hydrogen (\\HI{}) stretching from horizon-to-horizon through the South Galactic Pole, emanating from the Magellanic System. van~Kuilenburg (1972)\\footnote{Anomalously high-velocity gas features near the South Galactic Pole had actually been known since the work of Dieter (1965), but the link to the Magellanic System was not fully appreciated until that of Mathewson et~al. (1974).} discovered a lengthy high-velocity gas stream near the South Galactic Pole, while Wannier \\& Wrixon (1972) noted that the feature had a large and smoothly varying velocity (from $\\vlsr \\sim 0$ to $-400 \\kms$, or $\\vgsr \\sim 0$ to $-200 \\kms$), and was over $60\\deg$ long (but only $\\sim 4\\deg$ wide).  Mathewson et~al.  (1974) finally confirmed the connection between this feature and the Magellanic Clouds (MCs), suggesting the stream was 180\\deg{} in length, lying on a great circle. They showed the stream was clumpy, and gave the designations MSI--VI to the six dominant clumps (Mathewson et~al.  1977).  Most recently, Putman \\etal\\ (1998) showed what is now considered to be the full extent of the stream, with the identification of a leading arm feature (LAF) definitively associated with the Magellanic System. Indirect supporting evidence for both the trailing and leading arm streams being associated directly with the disrupting Magellanic Clouds is also provided by the similarity in chemical ``fingerprints'' between the gas in the streams and the gas in the Clouds (Lu et~al. 1998; Gibson et~al. 2000).  Building upon the seminal work of Murai \\& Fujimoto (1980), recent observational and theoretical analyses are consistent with the suggestion that the Clouds are close to peri-Galacticon. For example, the Galactocentric radial velocities of the Clouds are small at $84 \\kms$ (van der Marel \\etal\\ 2002) and $7 \\kms$ (Hardy, Suntzeff \\& Azzopardi 1989; Gardiner, Sawa \\& Fujimoto 1994 -- hereafter GSF94) for the LMC and SMC, respectively, compared with their respective transverse velocities in the Galactocentric frame of $280 \\kms$ (van der Marel \\etal\\ 2002) and $\\sim 200 \\kms$ (Lin, Jones \\& Klemola 1995) consistent with this hypothesis.  The closest approach to date, both between the Clouds and between the MCs and the Milky Way occurred $\\sim 200 \\Myr$ ago, and the orbital period of the SMC about the LMC is $\\sim 900 \\Myr$, with the Clouds as a pair orbiting the Galaxy with a period of order $1.5 \\Gyr$ (\\eg{} GSF94).  Early models from Murai \\& Fujimoto (1980) and Lin \\& Lynden-Bell (1982), supplemented with the recent proper motion work from Jones, Klemola \\& Lin (1994) have provided us with an accurate representation of the present-day orbital characteristics of the Magellanic System.  The determination of the orbital sense of the system demonstrates clearly that the MS is an extension stretching beyond the present galactocentric distance of the MCs, rather than a bridge joining the MCs with the Galaxy, and leading them (Lin \\& Lynden-Bell 1982; Lin \\etal\\ 1995), and also that the MCs are close to peri-Galacticon.  Considerable debate exists within the literature as to whether the Magellanic Stream is the result of ram pressure stripping (Moore \\& Davis 1994; Mastropietro \\etal\\ 2004) or gravitational tidal effects in which the Stream material is either stripped off the LMC (Weinberg 2000), the SMC (\\eg{} Gardiner \\& Noguchi 1996 -- hereafter GN96), or a common envelope (the inter-Cloud region; \\eg{} Heller \\& Rohlfs 1994). The observations of the LAF (Putman \\etal\\ 1998) show that tidal forces account for at least some fraction of the ``force'' shaping the existence of the Stream, even as the observed H$\\alpha$ emission measured along the Stream suggest that some additional ram pressure heating effects may be present (Weiner \\& Williams 1996; cf. Putman et~al. 2003b).  Yoshizawa \\& Noguchi (2003; hereafter YN03) have provided recently a significant improvement to the now canonical ``tidal'' model of GN96, via the inclusion of gas dynamics and star formation. In a prescient forebearer to YN03, Gardiner (1999) also provided important extensions to his earlier GN96 work using new constraints introduced by the recent discovery of the LAF, the addition of a drag term into the particle force equations, and an improved modelling of the LMC's disk potential. These latter modifications have the beneficial effect of mildly deflecting the orientation of the LAF with respect to the Magellanic System in a manner more consistent with the Putman et~al. (1998) dataset.  Encouraged by the success of these earlier studies, we are undertaking a comprehensive computational program aimed at providing the definitive deconstruction of this Rosetta Stone of hierarchical clustering -- the disrupting Magellanic System.  We now have access to the full HIPASS South {\\it and } North dataset, data which was not available to Putman et~al. (1998), allowing us to improve upon the observational constraints on both the trailing Stream and leading arm. Our {\\it ultimate} product will be the construction of a model which includes all relevant physical processes, including gas dynamics, ram pressure, radiative cooling, star formation, and chemical enrichment, all treated {\\it self-consistently} for the first time, in a hope to understand the physical processes of galactic cannibalism.  Our cosmological chemodynamical code {\\tt GCD+} (Kawata \\& Gibson 2003a,b) affords the power and flexibility to attack this problem in a manner previously inaccessible.  What follows represents the first of a series of papers devoted to this system; this Paper~I shows preliminary results based solely upon very high-resolution N-body simulations undertaken without the gas component of {\\tt GCD+} implemented.  This first step was required in order to allow a full exploration of orbital parameter space prior to the introduction of gas into the modelling.  The reason for doing so is that current observational constraints on the system still allow one some flexibility in choosing a unique orbital configuration for the system, partly due to our less-than-optimal understanding of the LMC and SMC masses.  The spatial orientation and nature of any SMC disk is also poorly constrained. Since N-body simulations are less computationally ``expensive'', we can survey different orbits for the Clouds, and determine our best orbital configuration(s).  In what follows we present our current best N-body model for the Magellanic Stream, compare this model with the extant observational data, and provide a roadmap for our future work, highlighting the successes and failures of the currently accepted canonical tidal model for the formation of the Magellanic Stream. ",
        "conclusions": "The combined neutral hydrogen gas mass in the Magellanic Clouds, Stream, Leading Arm, and inter-Cloud region, is in excess of $10^9 \\msun$, within a factor of three or so of the \\HI{} mass of the Milky Way itself.  Within the framework of hierarchical clustering, this represents a significant reservoir of potential fuel for future generations of star formation.  Our aim is to properly model the formation, evolution, and ultimate fate of the gas (and stars) associated with the Magellanic System.  The past decade has seen a wealth of new observational data for the System appear in the literature, in addition to the benefits of Moore's Law currently governing the increase in computational power. In combination, the two have allowed us to formulate improved models for the Magellanic Stream, a ``Rosetta Stone'' for the hierarchical clustering scenario of galaxy formation.  We have presented here a simulation with $\\sim 40$ times the resolution of previous simulations of the Stream, enabling us to examine more subtle features (kinematically and spatially) not previously considered in the models. Our gross results parallel those of Gardiner \\& Noguchi (1996) and Yoshizawa \\& Noguchi (2003), partly by construction, but the improved resolution and parameter space coverage here are unique.  The gross features of both the trailing Stream and Leading Arm Feature are successfully recovered.  The bifurcation of the Stream observed both spatially (Morras 1983) and kinematically (Putman \\etal\\ 2003) had been previously suggested to be due to the ``twisting'' motion of the SMC's orbit about the LMC. Our simulations are consistent with this picture, with the thinnest part of the Stream corresponding to the location where the orbits cross at \\mbox{$(l,b) \\sim (45\\deg,80\\deg)$}. The presence of this helix-like structure is an important test for any simulation that wishes to model the Magellanic System.  Despite the successes of the model, the comparisons we have made to date have uncovered several unresolved problems. First, there is still too little flux in the modelled MS, with the observed Stream (within the MSI--VI clumps) having 3--4 times (Putman \\etal\\ 2003) more mass than the modelled stream. Second, both the MS and LAF are too ``long'' in the models, with the LAF also being somewhat displaced from that observed.  Third, the ratio of LAF-to-MS gas mass appears to be too high in the simulations, by a factor of at least 3--4.  Fourth, while the LAF seems to have the correct deviation angle (cf. Gardiner 1999) relative to the great circle traced by the MS, its origin is somewhat offset from that inferred by the Putman et~al. (1998) dataset. Finally, the velocity dispersion of the currently modelled inter-Cloud region remains too high, but we speculate that the inclusion of gas dissipation in our future studies may alleviate this discrepancy.  Paper~II of this series will contain the full details of the parameter space coverage undertaken in our work, as well as a thorough examination of the spatial and kinematical bifurcations alluded to in Section~3.  Paper~III incorporates gas dynamics, star formation, radiative cooling, feedback, and chemical enrichment throughout the Magellanic System.  We will re-examine the orbits of the Clouds coupled with improved potentials for the LMC and Galaxy, as additional data becomes available (particularly on the shape of the Galactic potential -- e.g.  Martinez-Delgado \\etal\\ 2003; Bellazzini 2003; Helmi 2003).  A drag force term will also be introduced into the model, akin to that adopted by Gardiner (1999), in an attempt to better reproduce the geometry of the LAF.  The tools employed in analysing the simulations described here will shortly be released to the public -- this software package affords any user the ability to project virtually any N-body simulation into various projections representative of the observer's ``plane'', including the production of FITS files suitable for further analysis by any other astronomical software package."
    },
    "0402/astro-ph0402464_arXiv.txt": {
        "abstract": "{We report on the detection of the lowest \\COJ{1}{0} transition in the powerful high-redshift radio galaxy 4C\\,60.07 at $z$ = 3.79. The CO emission is distributed in two spatially and kinematically distinct components as was previously known from the observations of the higher excitation \\COJ{4}{3} line. The total molecular gas mass in 4C\\,60.07 inferred from the \\COJ{1}{0} emission is $M(\\mbox{H}_2)\\simeq 1.3\\times 10^{11}\\,\\Msolar$, sufficient to fuel the inferred star-formation rate of $\\sim 1600\\,\\Msolar\\,$yr$^{-1}$ for $10^8$\\,yrs. From our high-resolution \\COJ{1}{0} VLA maps we find the dynamical mass of 4C\\,60.07 to be comparable to that of a giant elliptical at the present time. A significant fraction of the mass is in the form of molecular gas suggesting that 4C\\,60.07 is in an early state of its evolution. The merging nature of 4C\\,60.07 along with its large dynamical mass imply that this system is a giant elliptical caught in its formative stages. ",
        "introduction": "High-redshift radio galaxies (HzRGs) are amongst the most luminous objects known, and are believed to serve as tracers of the peaks of the primordial density field around which giant elliptical galaxies and clusters of galaxies form (Kauffmann 1996; West et al.~1994). In the radio, HzRGs typically display a double-lobe morphology and large radio luminosities ($P_{178\\,MHz} \\sim 5 \\times 10^{35}$\\,erg\\,s$^{-1}$\\,Hz$^{-1}$), indicating a highly active black hole.  Recently, evidence has been mounting that HzRGs are massive starburst galaxies. This has come about from sub-millimetre detections of a number of HzRGs, implying large rest-frame far-IR luminosities ($L_{FIR} \\simeq 10^{13}\\,\\Lsolar$) powered by intense star formation ($SFR \\simeq 1000-2000\\,\\Msolar$\\,yr$^{-1}$ - Dunlop et al.~1994; Hughes et al.~1997; Ivison et al.~1998; Archibald et al.~2001). In a recent SCUBA survey of seven HzRGs and their surroundings, Stevens et al.~(2003) not only found the star formation in the radio galaxies themselves to be extended on several tens of kilo-parsec scales but also found one or more previously undetected submm sources in the vicinity ($50-250\\,$kpc) of more than half of the targeted objects. It is difficult to see how the Active Galactic Nucleus (AGN) could power the far-IR luminosity on $\\ga 10$\\,kpc scales, and a massive starburst seems to be the natural explanation. Indeed, adequate \"fuel\" for such large star formation rates has been found in the four HzRGs which have been detected in CO to date (Papadopoulos et al.~2000; De Breuck et al.~2003a; De Breuck et al.~2003b). These observations revealed the presence of massive ($\\sim 10^{11}\\,\\Msolar$) reservoirs of molecular gas, enough to fuel a $\\sim 1000\\,\\Msolar$\\,yr$^{-1}$ starburst for $\\sim 10^8$\\,yr, and in in half of the cases the CO emission was found to be extended on tens of kpc scales (Papadopoulos et al.~2000; De Breuck et al.~2003a). Similar large molecular gas masses distributed in clumps on tens of kilo-parsec scales has been found in a number of QSOs at high redshifts (Carilli et al.~2002a; Carilli et al.~2002b). In general, HzRGs have the advantage over quasars that they are not gravitationally lensed since they are usually selected on the basis of extended lobe-emission whereas quasars are often found to be lensed. Furthermore, HzRGs are known to be associated with giant ellipticals in the local Universe (McLure \\& Dunlop 2000).  In this paper we present high-resolution observations of the \\COJ{1}{0} emission from 4C\\,60.07 at $z=3.788$ using the Very Large Array\\footnote{The Very Large Array (VLA) is operated by the National Radio Observatory, which is a facility of the National Science Foundation, operated under cooperative agreement by Associated Universities, Inc.}. Throughout we have assumed $H_0=65$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_M=0.3$ and $\\Omega_\\Lambda=0.7$. In this cosmology the luminosity distance of 4C\\,60.07 is 36.2\\,Gpc and $1\\arcsecs$ corresponds to 7.7\\,kpc. ",
        "conclusions": "\\subsection{CO Luminosity and Molecular Gas Mass} The observed \\COJ{1}{0} line fluxes for the broad and narrow components in 4C\\,60.07 imply intrinsic CO luminosities of $\\LOJ{1}{0} = (1.0 \\pm 0.2)\\times 10^{11}$\\,K\\,km\\,s$^{-1}$\\,pc$^2$ and $(6.0 \\pm 0.7)\\times 10^{10}$\\,K\\,km\\,s$^{-1}$\\,pc$^2$, respectively.  For the $4-3$ line Papadopoulos et al.~(2000) found $\\LOJ{4}{3}=(7\\pm 1)\\times 10^{10}$\\,K\\,km\\,s$^{-1}$\\,pc$^2$ and $\\LOJ{4}{3}=(3.5\\pm 0.8) \\times 10^{10}$\\,K\\,km\\,s$^{-1}$\\,pc$^2$ for the broad and narrow component, respectively, where we have computed the luminosities in the cosmology adopted here. From the \\COJ{1}{0} line the molecular gas mass can be found using the well-known relation $M(\\rm{H}_2) = X_{CO} \\LOJ{1}{0}$ which relates the \\COJ{1}{0} luminosity with the molecular gas mass (e.g.~Strong et al.~1988). \\XCO~is the \\COJ{1}{0} line luminosity to H$_2$-mass conversion factor which in the extreme UV-intense environments found in local Ultra Luminous Infra-Red Galaxies, and presumably also in high redshift galaxies such as 4C\\,60.07, has a value of about 0.8\\,(K\\,km\\,s$^{-1}$\\,pc$^2$)$^{-1} \\Msolar$ (Downes \\& Solomon~1998). In doing so we find molecular gas masses of $M(\\rm{H}_2) \\sim 8 \\times 10^{10}\\,\\Msolar$ and $M(H_2) \\sim 5 \\times 10^{10}\\,\\Msolar$ for the broad and narrow CO emitting components, respectively. Hence, even for a conservative, non-Galactic value of $X_{\\rm{CO}}$ we find that about $\\sim 10^{11}\\,\\Msolar$ of molecular gas is associated with 4C\\,60.07. The estimated gas masses are in very good agreement with those of Papadopoulos et al.~(2000).  The total gas mass in 4C\\,60.07 will of course be larger once the neutral hydrogen has been accounted for. Assuming a value of $M(\\mbox{HI})/M(\\mbox{H}_2) = 2$ which is typically found in IRAS galaxies (Andreani, Casoli, \\& Gerin 1995) we find a total gas mass of $M_{gas} = M(\\mbox{H}_2) + 2M(\\mbox{H}_2) = 2.4\\times 10^{11}\\,\\Msolar$ for the broad component and $1.5\\times 10^{11}\\,\\Msolar$ for the narrow component. In case that metal-poor gas is also present the gas mass can be even higher.   \\begin{table} \\caption[]{Observed and Derived Properties for 4C\\,60.07.} $$ \\begin{array}{p{0.4\\linewidth}p{0.25\\linewidth}p{0.25\\linewidth}} \\hline \\noalign{\\smallskip} Parameter &  4C\\,60.07 & 4C\\,60.07\\\\ &  (broad)   & (narrow) \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} R.A.(J2000)   & 05\\hours12\\mins54\\psec75         &  05\\hours12\\mins55\\psec30\\\\ Decl.(J2000)  & +60\\degs30\\arcmins50\\parcsec92   &  +60\\degs30\\arcmins52\\parcsec29\\\\ $S_{\\rm{CO(1-0)}}$\\,[mJy]  &  $(0.27 \\pm 0.05)$   & $(0.30 \\pm 0.10)$ \\\\ $S_{\\rm{CO(1-0)}}\\Delta V$\\,[Jy\\,km\\,s$^{-1}$]  &  $(0.15 \\pm 0.03)$   & $ (0.09 \\pm 0.01)$ \\\\ \\mbox{L$'_{\\rm{CO}(1-0)}$}\\,[K\\,km\\,s$^{-1}$\\,pc$^2$]  & $(1.0 \\pm 0.2) \\times 10^{11}$   &  $(6.0 \\pm 0.7)\\times 10^{10}$\\\\ $M(\\mbox{H}_2)$ [$\\Msolar$]   & $8\\times 10^{10}$        &  $5\\times 10^{10}$\\\\ \\noalign{\\smallskip} \\hline \\end{array} $$ \\end{table} \\begin{figure*} \\centering \\includegraphics[width=0.4\\hsize,angle=0]{0788fig7a.eps}\\includegraphics[width=0.4\\hsize,angle=0]{0788fig7b.eps} \\caption{{\\bf a)} The narrow component tapered at 200k$\\lambda$. The resolution is $1\\parcsec72 \\times 1\\parcsec42$ at P.A.~= $285\\degs$. Contours are at -2, 2, 3, and $4 \\times \\sigma$, where $\\sigma = 55\\,\\mu$Jy\\,beam$^{-1}$. {\\bf b)} The \\COJ{1}{0} emission tapered down to 60k$\\lambda$ ($FWHM = 3\\parcsec22 \\times 3\\parcsec07$) overlaid as contours on a gray-scale representation of the \\J{4}{3} detection by Papadopoulos et al.~(2000). The contours are -2, 2, 3, and $4\\sigma$ with $\\sigma = 0.6$\\,mJy\\,beam$^{-1}$, and the gray-scale range is $0.6 - 2.0$\\,\\,mJy\\,beam$^{-1}$.} \\label{Fig2} \\end{figure*}  \\subsection{FIR Luminosity and star-formation efficiency} In the case where the far-IR luminosity is powered by a starburst and not an AGN, and all the stellar radiation is absorbed by dust, the far-IR to CO luminosity ratio, $L_{FIR}/L_{CO}$, provides a rough measure of the integrated luminosity of massive stars responsible for heating the dust ($L_{FIR}$) relative to the amount of fuel available for star formation ($L'_{CO}$). We estimate the far-IR luminosity of 4C\\,60.07 to be $L_{FIR} \\simeq 4\\times 10^{13}\\,\\Lsolar$, where we have used the 850-$\\mu$m flux measurement of Archibald et al.~(2001) and adopted a dust temperature and spectral index of $T_{d}=50\\,$K and $\\beta = 1.5$, respectively. This yields a $L_{FIR}/L_{CO}$ ratio of $\\simeq (4\\times 10^{13}\\Lsolar)/(1.6\\times 10^{11}$\\,K\\,km\\,s$^{-1}$\\,pc$^2) = 250\\,\\Lsolar ($\\,K\\,km\\,s$^{-1}$\\,pc$^2)^{-1}$, which is similar to values found in ULIRGs (e.g.~Solomon et al.~1997). Carilli et al.~(2002a) found continuum-to-line ratios of 350 and 323 in the QSOs BRI\\,1202-0725 ($z=4.70$) and BRI\\,1335-0417 ($z=4.41$), respectively.  A somewhat lower value of 200 was found in the $z=4.12$ QSO PSS\\,J2322+1944 (Carilli et al.~2002b).  The current star-formation rate measured in $\\Msolar$\\,yr$^{-1}$ is given by \\begin{equation} SFR \\simeq (L_{FIR}10^{-10}/\\Lsolar) \\delta_{MF}\\delta_{SB}, \\end{equation} where $\\delta_{MF} \\sim 0.8-2$ depends on the mass function of the stellar population, and $\\delta_{SB}$ is the fraction of the far-IR luminosity which is powered by the starburst and not the AGN (Omont et al.~2001). If we adopt $\\delta_{MF}=0.8$ and conservatively assume that only 50\\% of $L_{FIR}$ is powered by the starburst ($\\delta_{SB}=0.5$) we find $SFR \\sim 1600\\,\\Msolar\\,\\mbox{yr}^{-1}$. Assuming that the total amount of molecular gas observed towards 4C\\,60.07 ends up as fuel for the starburst, we find that there is enough molecular gas to sustain the inferred star-formation rate for $\\sim 8 \\times 10^7\\,\\mbox{yr}$. While this is short compared to the Hubble time it is still sufficient to produce a giant elliptical with a stellar mass of $M_* \\simeq 10^{11}\\,\\Msolar$.  The efficiency with which stars are being formed, i.e.~the rate of star-formation per solar mass of molecular hydrogen, is given by $SFE = SFR/M(\\mbox{H}_2)$ or equivalently $L_{FIR}/M(\\mbox{H}_2)$. For 4C\\,60.07 we find $SFE = L_{FIR}/M(\\mbox{H}_2) \\simeq 300\\,\\Lsolar\\,\\Msolar^{-1}$, which is somewhat higher than the $\\sim 190\\,\\Lsolar\\,\\Msolar^{-1}$ reported by Papadopoulos et al.~(2000) who derived their value based on their flux density measurement of $S_{850\\mu m}=11\\,$mJy. Here we have used the measurement by Archibald et al.~(2001) which yields a somewhat higher 850\\,$\\mu$m flux density of 17\\,mJy, although still lower than the 22\\,mJy reported by Stevens et al.~(2003).  Local ULIRGs exhibit star-formation efficiencies comparable to that of 4C\\,60.07, once the same \\XCO-factor has been used (Solomon et al.~1997). It is worth stressing that our detection of \\COJ{1}{0} enables us to make a direct comparison of SFEs with that of local ULIRGs since the same gas mass measure (the \\J{1}{0} line) is used for ULIRGs.  The apparently high star-formation efficiencies found for the above systems, could be severely overestimated if the AGN contributes significantly to the far-IR luminosity. However, the detection of CO together with the fact that 4C\\,60.07 and other HzRGs appear extended on several tens of kilo-pc scales at submm-wavelengths (Ivison et al.~2000; Stevens et al.~2003) strongly suggests that the far-IR emission is powered by large-scale starburst and not the AGN. Here we must mention that while \\COJ{1}{0} may be a good indicator of the total metal-rich H$_2$ gas resevoir, it may be a poor one regarding the dense gas phase ($n \\ge 10^5\\,\\mbox{cm}^{-3}$) that \"fuels\" star-formation. The latter could be particular true in the tidally disrupted giant molecular clouds (GMCs) expected in ULIRGs where a diffuse phase may dominate the \\COJ{1}{0} emission but has little to do with star formation. This could be the reason why the $L_{FIR}/L'_{CO}$ ratio is found to be such a strong function of $L_{FIR}$, increasing for the merger systems associated usually with large far-IR luminosities. Interestingly, recent work shows that the SFE of dense gas, parametrised by the $L_{FIR}/L_{HCN(1-0)}$ ratio (the HCN \\J{1}{0} critical density is $2 \\times 10^5\\,\\mbox{cm}^{-3}$), remains constant from GMCs all the way to ULIRG system (Gao \\& Solomon 2003).  \\subsection{Excitation Conditions} From the detections of CO in 4C\\,60.07 we can infer the CO $(4-3)/(1-0)$ velocity/area-averaged brightness temperature ratio using $r_{43} = \\rm{L}'_{\\rm{CO}(4-3)}/\\rm{L}'_{\\rm{CO}(1-0)}$. The global line ratio, i.e.~the line ratio obtained by combining the flux from the two components, is $r_{43}=0.7^{+0.3}_{-0.2}$. The line ratios for the broad and narrow components are $0.7^{+0.3}_{-0.2}$ and $0.6^{+0.2}_{-0.2}$, respectively. Hence, given the large uncertainties we find no significant difference in the excitation conditions between the two components. Thus in this case H$_2$ mass estimates solely from \\COJ{4}{3} assuming full thermalisation and optical thickness of the latter (i.e.~$r_{43} \\simeq 1$) would not result in too large errors. However, the sub-thermal excitation of such high-J CO lines remains a possibility even in starburst environments (e.g.~Papadopoulos \\& Ivison 2002; Greve et al.~2003).   In Figure \\ref{fig:line_ratio} we plotted the velocity-integrated CO line flux densities (normalised to \\J{4}{3}) for the broad and narrow components in 4C\\,60.07 as well as the normalised line fluxes for PSS\\,J2322+1944. We have used a standard single-component large velocity gradient (LVG) code to interpret the observed line ratio. In fact, given the large range compatible with our measurements we can only put a lower limit on the density. Indeed, for both components the upper values for $r_{43}$ are compatible with LTE and optically thick emission where the ratio is no longer sensitive to the density.   A lower limit on the average gas density can be set by the lowest possible value of $r_{43}=0.4$ allowed by the observations. Adopting a typical CO abundance $\\Lambda = X_{\\rm{CO}}/(dV/dr) = 10^{-5}\\,$(km\\,s$^{-1}$\\,pc$^{-1})^{-1}$, a $T_{CMB} = (1 + z)\\times 2.73\\,\\mbox{K} = 13.07\\,$K and a $T_{kin} = 50\\,\\mbox{K}$, which is a typical dust temperature in starburst environments (Colbert et al.~1999; Hughes et al.~1997), we find $n(\\mbox{H}_2) \\geq 2 \\times 10^{3}\\,\\mbox{cm}^{-3}$. Adopting a higher $\\Lambda = 10^{-4}$ (unlikely in such kinematically violent, UV-aggressive environments) lowers the aforementioned limit to $\\sim 600\\,\\mbox{cm}^{-3}$. A lower assumed temperature of $T_{kin}=30$\\,K does not change these lower limits by much. For the most likely value of $r_{43}=0.7$ we find $n(\\rm{H}_2) = 3 \\times 10^{3}$\\,cm$^{-3}$ ($T_k = 50$\\,K, $\\Lambda  = 10^{-5}\\,$(km\\,s$^{-1}$\\,pc$^{-1})^{-1}$).   Comparing with the $z=4.12$ QSO PSS\\,J2322+1944 which has a (4-3)/(1-0) line ratio of $1.4\\pm0.6$ (Carilli et al.~2002b; Cox et al.~2002) it appears that the excitation conditions in 4C\\,60.07 are less extreme. Hence, these results may indicate that the molecular gas in 4C\\,60.07 is not as dense as that seen in PSS\\,J2322+1944. \\begin{figure} \\centering \\includegraphics[width=1.0\\hsize,angle=0]{0788fig8.ps} \\caption{Velocity-integrated CO line flux densities from 4C\\,60.07 are shown as filled squares. The integrated line fluxes have been normalised to the \\J{4}{3} line. Also shown are the line fluxes from the $z=4.12$ QSO PSS\\,J2322+1944 (filled triangles - Carilli et al.~2002b; Cox et al.~2002), the $z=4.69$ QSO BRI\\,1202-0725 (open circles - Omont et al.~1996; Carilli et al.~2002a), the local starburst galaxy M\\,82 (open triangles - Mao et al.~2000), and the integrated emission from within the solar radius of the the Milky Way (filled circles - Fixsen et al.~1999). The red dotted line shows the line flux increasing as frequency squared which is expected for optically thick conditions. The green lines show results from a LVG-model with $T_{kin} = 50\\,$K, $X_{\\rm{CO}}/(dV/dr) = 10^{-5}\\,$(km\\,s$^{-1}$\\,pc$^{-1})^{-1}$, and $n(\\rm{H}_2)$ equal to $3.0\\times 10^3$\\,cm$^{-3}$ (solid line) and $2.2\\times 10^3$\\,cm$^{-3}$ (dashed line).} \\label{fig:line_ratio} \\end{figure} Such a comparison, however, is prone to the effects of strong lensing. While 4C\\,60.07 is unlikely to be a strongly lensed system because of its CO components with widely different line widths, it is not the case for PSS\\,2322+1944 which is known to be lensed (Carilli et al.~2003). If the distribution of the \\JJ{1}{0} and \\J{4}{3} emitting gas in the source plane differs, differential magnification of the two lines can result in erroneous intrinsic line ratios.   \\subsection{What is the Evolutionary Status of 4C\\,60.07?} A comparison between the molecular gas mass and the dynamical mass allows for the determination of the evolutionary status of a galaxy, while a comparison of its dynamical mass with that of present-day spiral or elliptical can point toward its possible descendant. Typically, dynamical masses are calculated assuming the gas is distributed in a disk in Keplerian rotation (e.g.~Carilli et al.~2002a). In the case of 4C\\,60.07, however, we have sufficient spatial and kinematical information to conclude that the two gas components do not belong to such a structure. The detection of two kinematically distinct gas resevoirs in 4C\\,60.07, each with a large gas mass, suggest a major merger event. Hence, a more plausible scenario might be that the two clouds are part of spherical system in the process of collapsing. Assuming the system is virialised, one can apply the virial theorem to derive the following expression for the dynamical mass \\begin{equation} M_{dyn} = 2.33\\times 10^9 \\left ( \\frac{R}{\\mbox{kpc}} \\right ) \\left ( \\frac{\\Delta V_{FWHM}}{100\\, \\mbox{km\\,s}^{-1}} \\right ) ^2 \\left ( \\frac{1}{\\alpha (1+q)} \\right ), \\label{eq:Mdyn} \\end{equation} where $q$ is a factor which described the influence of non-gravitational forces on the virial equilibrium, and $\\alpha$ ranges from 0.55 to 2.4 with a typical value of 1.5 which we shall adopt here (see Bryant \\& Scoville (1996) for details). A possible caveat to this argument comes from the fact that if the assumption of a virialised system is not true, eq.~\\ref{eq:Mdyn} will tend to overestimate the true dynamical mass. Adopting a radius of the sphere corresponding to half the separation between the two components, i.e. $R\\sim 15$\\,kpc (corresponding to $2\\arcsecs$), and a FWHM width equal to the velocity offset between the line centers of the two gas resevoirs, i.e.~$\\Delta V_{FWHM}\\sim 700$\\,km\\,s$^{-1}$, leads to an enclosed dynamical mass of $M_{dyn}\\la 1.1\\times 10^{12}\\,\\Msolar$, where we have assumed that non-gravitational forces are negligible, i.e.~$q << 1$. Note that since Bryant \\& Scoville (1996) derived eq.~\\ref{eq:Mdyn} by examining a large number of clouds filling up a spherical region, we have in the above assumed that the velocity and spatial separation between the two CO components is a fair approximation for $\\Delta V_{FWHM}$ and $R$ in eq.~\\ref{eq:Mdyn}. Alternatively, we can do the more simple calculation of finding the virial mass required to keep the two clouds bound by a spherical potential. In this case the velocity dispersion becomes $\\Delta V = (1/2) 700\\,\\mbox{km\\,s}^{-1} = 350\\,\\mbox{km\\,s}^{-1}$, and we find $M_{virial} = G^{-1} R \\Delta V^2 \\simeq 4.3 \\times 10^{11}\\,\\Msolar$, where we have assumed $R=15$\\,kpc. Projection effects will only tend to make $R$ and $\\Delta V_{FWHM}$, and therefore also the above mass estimates, larger. Thus, 4C\\,60.07 is a massive system with a dynamical mass comparable to that of giant elliptical galaxies seen in the present day Universe. The inferred ratio of the molecular-to-dynamic mass for the system as a whole is $M(\\mbox{H}_2)/M_{dyn} \\sim 0.12$, and the total gas fraction is $M_{gas}/M_{dyn} \\sim 0.35$. The high gas fraction, along with the large velocity dispersion of the system suggest that 4C\\,60.07 is a very young system in its early stages of formation.   The fact that such a massive ($M_{dyn} \\sim 10^{12}\\,\\Msolar$) system has assembled at $z=3.8$, which corresponds to a time when the Universe was only $\\sim 10$\\% of its present age, seems to favour the monolithic collapse model (e.g.~Eggen, Lynden-Bell \\& Sandage 1962; Tinsley \\& Gunn 1976) of massive ellipticals over the hierarchical formation scenario (e.g.~Baron \\& White 1987; Baugh et al.~1996; Kauffmann \\& Charlot 1998). Further in support of the monolithic collapse scenario is the large star-formation rate ($SFR\\sim 1600\\,\\Msolar$yr$^{-1}$) we infer from the far-IR luminosity. If such a large star-formation rate can be sustained it is capable of producing a giant elliptical in a time scale comparable to the dynamical time. The dynamical time-scale for a system with a mass $M$ within a radius $R$  is \\begin{equation} t_{dyn} = 7.42 \\times 10^{11} \\left ( \\frac{R}{\\mbox{kpc}} \\right ) ^{3/2} \\left ( \\frac{M}{\\Msolar} \\right )^{-1/2}, \\end{equation} where $t_{dyn}$ is measured in years (see p.~37 of Binney \\& Tremaine 1987). For a giant elliptical with a mass $M_{ell} = 10^{12}\\,\\Msolar$ and a radius $R=30$\\,kpc, we find $t_{dyn} \\simeq 1.2 \\times 10^8$\\,yr. Hence, in order to form a giant elliptical within a dynamical time scale, star-formation rates of the order $M_{ell}/t_{dyn} \\simeq 8000\\,\\Msolar\\,$yr$^{-1}$ are required, which is comparable to what we find in 4C\\,60.07. It is important to note that even if the true dynamical mass or the star-formation rate of this system are less than the above estimates, the fact that at $z=3.8$ about $1.3 \\times 10^{11}\\,\\Msolar$ in H$_2$ gas mass alone has assembled within 15\\,kpc points towards a major galaxy forming merger rather than a gradual accumulation of mass over time and redshift. This mass can only be revised upwards and typically reach $\\sim 4 \\times 10^{11}\\,\\Msolar$ if HI is also accounted for, thereby making the case for a monolithic collapse scenario even stronger.  4C\\,60.07 is not the only system with extended CO emission. Molecular gas distributed on tens of kpc scales, sometimes in separate gas components, have been observed in high redshifts quasars (e.g.~Carilli et al.~2002a; Papadopoulos et al.~2001) and radio galaxies (De Breuck et al.~2003; De Breuck et al.~in prep.). Thus, there is evidence to show that at least some luminous active galaxies in the early Universe, in addition to harbouring supermassive black holes in their centres, are associated with massive reservoirs of molecular gas distributed on tens of kpc scales which are in the process of merging. This observed coevality between large scale mergers of gaseous subsystems and the epoch of AGN-activity might provide clues to the origin of the tight relationship observed locally between the velocity dispersion of spheroids and the mass of their central black holes (Ferrarese \\& Merritt 2000)."
    },
    "0402/astro-ph0402522_arXiv.txt": {
        "abstract": "Within the standard big bang nucleosynthesis (BBN) and cosmic microwave background (CMB) framework, the baryon density measured by the Wilkinson Microwave Anisotropy Probe (WMAP) or the primordial D abundance is much higher than the one measured by the $^4$He or $^7$Li abundances. To solve the discrepancy, we propose a scenario in which additional baryons appear after BBN. We show that simply adding the baryons can not be a solution but the existence of a large lepton asymmetry before BBN makes the scenario successful. These extra baryons and leptons, in addition to the initial baryons which exist before the BBN, can be all produced from $Q$-balls. ",
        "introduction": "\\label{sec:intro}  The baryon density is one of the most important cosmological parameters. Especially, it is the only one input parameter for the standard big bang nucleosynthesis (BBN) theory, which predicts the abundances of light elements, D, $^4$He and $^7$Li. Meanwhile, for the cosmic microwave background (CMB), it also plays an important role in determining the shape of the acoustic peaks. Observations of the three light elements and the CMB on the whole indicate equal amount of baryons and make us confirm the validity of the standard cosmology.  Recently, following the precise measurement of the CMB by the Wilkinson Microwave Anisotropy Probe (WMAP), its concordance with the BBN and the light elements observations has been investigated in Refs~\\cite{Cyburt:2003fe,Coc:2003ce,Coc:2004ij,Cyburt:2004cq}. The baryon density measured from the WMAP data is $\\omega_b \\equiv \\Omega_b h^2 = 0.024 \\pm 0.001$ (with the power-law $\\Lambda$CDM model)~\\cite{Spergel:2003cb}, where $\\Omega_b$ is the baryon energy density divided by the critical energy density today and $h$ is the Hubble constant in units of 100 km s$^{-1}$ Mpc$^{-1}$. The uncertainty is very small because the WMAP has detected the first and second peaks accurately in the temperature angular spectrum \\cite{Page:2003eu}. This corresponds to $\\eta \\equiv n_b/n_{\\gamma} = (6.6 \\pm 0.3) \\times 10^{-10}$, where $n_b$ and $n_{\\gamma}$ are baryon and photon number densities, via the relation $\\eta = \\omega_b/(3.65 \\times 10^7)$. Refs~\\cite{Cyburt:2003fe,Coc:2003ce,Coc:2004ij,Cyburt:2004cq} take this well-determined WMAP $\\omega_b$ as the BBN input\\footnote{To be more precise, they adopt $\\eta = (6.14 \\pm 0.25) \\times 10^{-10}$ which is obtained from the running spectral index model value $\\omega_b = 0.0224 \\pm 0.0009$ via the $\\omega_b$-$\\eta$ relation. Since this value is inferred from the combination of CMB, galaxy and Lyman $\\alpha$ forest data, we adopt the value quoted in the text which is inferred using only the WMAP data.} and calculate the light elements abundances and their theoretical errors using improved evaluations of nuclear reaction rates and uncertainties. The results are compared with the received measurements of the primordial abundances of three light elements, D (Ref.~\\cite{Kirkman:2003uv}), $^4$He (Refs.~\\cite{Fields1998} or \\cite{Izotov:2003xn}) and $^7$Li (Refs.~\\cite{Ryan2000} or \\cite{Bonifacio:2002yx}). Although there are small differences concerning their adopted reaction rates or observation data, their conclusions agree: from the WMAP baryon density, the predicted abundances are highly consistent with the observed D but not with $^4$He or $^7$Li. They are produced more than observed. Especially, the $^7$Li-WMAP discrepancy is severer and it may require an explanation.  The most conservative and likely interpretation of such discrepancy is that systematic errors in the primordial $^4$He and $^7$Li measurements are underestimated in spite of the thorough analysis hitherto. Or, as Ref.~\\cite{Coc:2003ce} has pointed out, it is possible that since the cross section of the reaction $^7$Be(d,p)2$^4$He has not been measured for the BBN energy range, it could reduce the $^7$Li yield to match with the observation if the rate turns out to be hundreds times more than the often neglected value obtained by the extrapolation.  If the discrepancy can not be attributed to systematic errors, we would have to invoke new physics to reduce their abundances. Actually, some approaches to reduce $^7$Li by astrophysical means, such as non-standard depletion mechanism inside stars, are discussed in Refs.~\\cite{Cyburt:2003fe, Coc:2004ij,Romano:2003gv} and references therein.  What we seek in this paper is a cosmological solution. For a long time, not a few non-standard models are known to affect $^4$He abundance but none of them seem to be able to solve the discrepancy. For example, non-standard expansion rate (extra relativistic degrees of freedom) and/or large lepton asymmetry change $^4$He too much while adjusting $^7$Li to its observed value \\cite{Barger:2003zg,Barger:2003rt,Steigman:2003gy}. A varying fine structure constant (electromagnetic coupling constant) can relieve the tension between either D and $^4$He or D and $^7$Li but not the three elements together~\\cite{Nollett:2002da}. At this time, only bolder attempt such as to combine non-standard expansion rate and a varying fine structure constant investigated by two of the authors in Ref.~\\cite{Ichikawa:2004ju} seems to be able to accommodate D, $^4$He and $^7$Li simultaneously, but whether the introduction of these non-standard ingredients can fit the WMAP data is not fully checked yet \\footnote{A varying deutron binding energy may have the capacity to render internal agreement between the light element abundances and with WMAP \\cite{Dmitriev:2003qq}.}.  In this paper, we investigate a solution which allows different amount of baryons for the BBN and the CMB. Unfortunately this only reconciles the WMAP, D and either $^4$He or $^7$Li. We choose to make $^7$Li consistent with the observation in this way and consider a large lepton asymmetry in addition in order to accommodate $^4$He too. One of the advantages of this scenario is that both the additional baryons and lepton asymmetry can be produced from $Q$-balls, which also generate the original baryons before the BBN.  In next section, we explain the discrepancy briefly and investigate how the prediction of the primordial light element abundances is affected by the additional baryons after the BBN. After we show that is not enough to solve the discrepancy, we introduce the lepton asymmetry and calculate how much extra baryons and lepton asymmetry are needed to be a solution. In section~\\ref{sec:qball}, we describe how to generate required baryons and leptons. We conclude in section~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion} Concerning recent cosmological observations, there seems to be conflicting measurements of the baryons in the universe, namely, the baryon density deduced from the CMB observation by the WMAP or the primordial D abundance is much higher than the one derived from $^4$He or $^7$Li abundances. In this paper, we have proposed a scenario to reconcile such inconsistency by adding baryons after the BBN and assuming a lepton asymmetry before the BBN. The introduction of the additional baryons leads to success in explaining the observed low abundance of $^7$Li without recourse to special models of stellar depletion. Also it was shown that the scenario can be naturally implemented in the AD mechanism, where $Q$-balls play an essential role in preserving extra lepton and baryon asymmetries. We should note that  this is the first cosmological scenario that can completely remove the long-standing tension among the three light elements and CMB. According to our result, amazingly, about one third of the baryon we observe in the present universe might have been ``sterile\" during the BBN epoch. Such illuminating history of the universe might be confirmed through further observational results with better accuracy in the future."
    },
    "0402/astro-ph0402086_arXiv.txt": {
        "abstract": "We report calculations of the stellar and gaseous  response to a Milky Way mass distribution model including a spiral pattern with a locus as traced by K-band observations, over imposed on the axisymmetric components in the plane of the disk. The stellar study extends calculations from previous work concerning the self-consistency of the pattern. The stellar response to the imposed spiral mass is studied via computations of the central family of periodic and nearby orbits as a function of the pattern rotation speed, $\\Omega_p$, among other parameters. A fine grid of values of $\\Omega_p$ was explored ranging from 12 to 25 $km~s^{-1}~kpc^{-1}$. Dynamical self-consistency is highly sensitive to $\\Omega_p$, with the best fit appearing at 20 $km~s^{-1}~kpc^{-1}$. We give an account of recent independent pieces of theoretical and observational work that are  dependent on the value of $\\Omega_p$, all of which are consistent with the value found here; the recent star formation history of the Milky Way, local inferences of cosmic ray flux variations and Galactic abundance patterns. The gaseous response, which is also a function of $\\Omega_p$, was calculated via 2D hydrodynamic simulations with the ZEUS code. For $\\Omega_p = 20~km~s^{-1}~kpc^{-1}$, the response to a two-armed pattern is a structured pattern of 4 arms, with bifurcations along the arms and interarm features. The pattern resembles qualitatively the optical arms observed in our Galaxy and other galaxies. The complex gaseous pattern appears to be linked to resonances in stellar orbits. Among these, the 4:1 resonance plays an important role, as it determines the extent of the stellar spiral pattern in the self-consistency study here presented. Our findings seemingly confirm predictions by Drimmel and Spergel (2001), based on K band data. ",
        "introduction": "The comparison of near-infrared and optical images of external galaxies reveals interesting differences. Striking examples are M81 and NGC 2997 (see pictures in Elmegreen, D., 1981; and Block et al 1994, respectively). It is common to observe a smooth, simple 2-armed K band pattern but a more complex pattern in the optical blue band, often suggesting more arms and bifurcations (segments of arms that appear to be connected to a K band arm but are not detectable in the infrared). A two-armed smooth structure underlying a more complex morphology also appeared in the work of Grosb{\\o}l, Pompei and Patsis (2002): in a K band study of 53 nearby spirals, most galaxies displayed a grand-design, two-armed, symmetric pattern in their inner regions which often breakups into tighter winded, multiple arms further out. Nonlinear effects were invoked to explain such morphology.  In recent work, data from COBE-DIRBE have shed light into the Milky Way spiral pattern. Drimmel (2000), and Drimmel \\& Spergel (2001) have presented a comprehensive picture of how this pattern is like, presenting emission profiles of the Galactic plane in the K band and at 240 $\\mu m$. The former data set, which suffers little absorption and traces density variation in the old stellar population, is dominated by a two-armed structure with a minimum pitch angle of 15.5$^\\circ$. At 240 $\\mu m$, the pattern is consistent with the standard four-armed model, that corresponding to the distribution of the youngest stellar populations delineated by HII regions.  The conventional picture of the spiral pattern of our Galaxy maps at least 4 arms, named Norma, Crux-Scutum, Carina-Sagittarius and Perseus (for a recent review see Vall\\'ee 2002, who also reports a likely pitch angle of 12$^\\circ$ for this pattern). Additionally, features such as the Orion spur at the Solar neighborhood have been revealed (Georgelin \\& Georgelin 1976). Drimmel (2000) laid down, from the comparison, the hypothesis that the 4-armed structure is the gas response to the 2-armed ``stellar\" pattern.  Assuming that indeed the K band data is by far a better tracer of mass than the optical data of spiral structure, in this work we model the spiral pattern from the locus and pitch angle of Drimmel (2000) and study its self-consistency, a requirement we consider it must satisfy. As the answer strongly depends on the pattern speed, this study should yield a value for this fundamental Galactic parameter.  The value of the pattern speed of the Galaxy has been a matter of controversy for a long time. From the values proposed by Lin, Yuan \\& Shu (1969) of $\\Omega_p = 11-13~km~s^{-1}~kpc^{-1}$, numbers in the range of 10-60 $km~s^{-1}~kpc^{-1}$ have been used in the literature. In a previous paper (Pichardo et al 2003, hereafter P1), we explored the stellar dynamics in a full axisymmetric model for our Galaxy superimposing our modeling of mass distribution for the spiral pattern: the locus of Drimmel (2000), the optical locus of Vall\\'ee (2002), and a superposition of both. P1 did not assume the usual simple perturbing term (a cosine term for the potential) that had been used in the literature in the modeling of spiral arms: it is precisely the very prominent spiral structure in red light what suggests to us that such structure should be considered an important galactic component worthy of a modeling effort beyond the simple perturbing term. P1 modeling consists of a superposition of oblate spheroids for the spiral. Two different values of $\\Omega_p = 15, 20~km~s^{-1}~kpc^{-1}$, were tried. The best self-consistency was achieved with $\\Omega_p = 20~km~s^{-1}~kpc^{-1}$ for the locus and pitch angle of Drimmel (2000). Figure 7 of P1, a mosaic of our self-consistency test -- explained below -- for parameters including the global spiral mass and loci,  exhibits a remarkably good response for this case which rules out  the other 3 cases in the panel. In fact, the behavior at $\\Omega_p = 20~km~s^{-1}~kpc^{-1}$ is so good that one can hardly envision an improvement on general theoretical grounds. The question that comes to mind is, given the spiral locus and pitch angle of Drimmel (2000), our adopted mass distribution modeling, and the spiral mass implied by observations of external galaxies applied to those parameters, now fixed, are there other values of $\\Omega_p$ which also satisfy our self-consistency criterion to that accuracy? In this paper, we extend the self-consistency calculations to a finer range of values of $\\Omega_p$ in order to answer that question. Once the best value is found, we put it to the test in two ways: calculating the gaseous response to find out whether it is consistent with the observed optical Galactic spiral pattern, which amounts to a test of Drimmel's hypothesis. Finally, we review recent work using independent data sets which are sensitive to the value of $\\Omega_p$. This is another way of testing our prediction against \"nature\", and not only versus different modeling approaches, as the subjects of those independent studies are quite different from our Galactic modeling effort.  A continued line of work by Contopolous and collaborators (see, v.g. Patsis, Grosb{\\o}l and Hiotelis 1997 and references therein) has provided the framework to study the response of gaseous disks to spiral perturbations. In that paper, a comparison between SPH models with Population I features observed on B images of normal, grand design galaxies, showed that the 4:1 resonance generates a bifurcation of the arms and interarm features. Furthermore, Contopoulos \\& Grosb{\\o}l (1986,1988) had shown that the central family of periodic orbits do not support a spiral pattern beyond the position of the 4:1 resonance, which thus determines the extent of the pattern. Weak spirals can extend their pattern up to corotation, from linear theory. A phenomenological link between resonances, the angular speed, and the stellar and gas patterns in spirals is complemented  by the study of Grosb{\\o}l and Patsis (2001) using deep K band surface photometry to analyze spiral structure in 12 galaxies. They find that the radial extent of the two-armed pattern is consistent with the location of the major resonances: the inner Lindblad resonance (ILR), the 4:1 resonance, corotation and the outer Lindblad resonance (OLR). For galaxies with a bar perturbation, the extent of the main spiral was better fitted assuming it is limited by corotation and the OLR. Using $H_0 = 75~km~s^{-1}~Mpc^{-1}$, the pattern speed was found to be for the entire sample of the order of 20 $km~s^{-1}~kpc^{-1}$, and remarkably, not a sensitive function of morphological type or total mass.  In the following section  we describe our results for the stellar orbital response to the imposed spiral pattern, through which $\\Omega_{p}$ is determined. ",
        "conclusions": "As found by Contopoulos \\& Grosb{\\o}l (1988), self-consistency is improved by introducing velocity dispersion; this is a realistic effect that can only be neglected by arguing that in strong spirals nonlinearity dominates. For our Galaxy, the observations of Drimmel \\& Spergel (2001) suggesting the termination of the spiral at the position of the 4:1 resonance indicate a marginally strong spiral in Contopoulos \\& Grosb{\\o}l's (1988) framework. On the opposite side, our best self-consistent model is found at the lowest spiral mass considered in P1, 0.0175 $M_D$, for which no stochastic motion was found. From this fact, a weak spiral and a linear regime come to mind. A possible solution to this issue could be that, while the strength of the spiral begins to fall at the 4:1 resonance, termination occurs at corotation, as predicted for the weak case in that framework. However, the quite different modeling of the galactic -- particularly, spiral -- mass distribution used in their studies and ours could make an interpretation of our results in their framework an unfair one. At any rate, the response at our best $\\Omega_p$ is so flat that there appears no need to invoke velocity dispersions. In an analysis based on periodic orbits, such as ours, such dispersion will be small. Other close values of $\\Omega_p$ could be improved by this effect in a study out of the scope of this work, involving many orbits departing from the periodic ones and hence subjected to larger velocity dispersions.  We conclude that a two-armed density spiral imposed taking into account observational restrictions from K-dwarf distributions yields an optimal dynamically self-consistent model, for values close to $\\Omega_p$ = 20 $~km~s^{-1}~kpc^{-1}$. The fact that various independent estimates of this quantity, sensitive to highly distinct physics, yield values for $\\Omega_p$ in agreement with our estimation, gives us confidence in the result. In regard to the gaseous response, we notice that the independent studies we quote for comparison are not only consistent with the value found here for $\\Omega_p$, but also with a density distribution corresponding to a 4-armed gaseous pattern with structural features reminiscent of optical observations of our Galaxy and other spirals.     \\begin{figure} \\psfig{file=fig1.eps,angle=0,width=9.0cm} \\caption{Self-consistency analysis for $\\Omega_p = 20 ~km~s^{-1}~kpc^{-1}$. The proposed spiral locus is shown with open squares. A set of periodic orbits are traced with continuous lines, and the maxima in the response density are the filled (black) squares. The frame of reference is the rotating one where the spiral pattern is at rest} \\label{fig1} \\end{figure}  \\begin{figure} \\psfig{file=fig2.eps,angle=0,width=9.0cm} \\caption{Simulation with the ZEUS code of the gas response to a spiral pattern with $\\Omega_p = 20 ~km~s^{-1}~kpc^{-1}$, open squares, shown in the rotating frame of the spiral pattern after 2.55 Gyr of evolution. The arrows give the velocity field, their size being proportional to the speed, with the maximum speed shown being 212 $km~s^{-1}$. Dense zones correspond to dark regions.} \\label{fig2} \\end{figure}"
    },
    "0402/astro-ph0402565_arXiv.txt": {
        "abstract": "{We present an {\\it XMM-Newton} analysis of the cataclysmic variable T~Leo. The X-ray light curve shows sinusoidal variation on a period $P_x$ equal to $0.89^{+0.14}_{-0.10}$ times the previously spectroscopically determined orbital period. Furthermore, we find a signal in the power spectrum at 414~sec that could be attributed to the spin period of the white dwarf. If true, T~Leo would be the first confirmed superoutbursting intermediate polar (IP).  The {\\it spin profile} is double-peaked with a peak separation of about 1/3 spin phases. This appears to be a typical feature for IPs with a small magnetic field and fast white dwarf rotation.  An alternative explanation is that the 414~sec signal is a Quasi-periodic Oscillation (QPO) that is caused by mass transfer variation from the secondary, a bright region (``blob'') rotating in the disc at a radius of approximately~9$\\Rwd$ or -- more likely -- a travelling wave close to the inner disc edge of a dwarf nova with a low field white dwarf.  The {\\it XMM-Newton} RGS spectra reveal {\\it double peaked emission} for the O~VIII Ly $\\alpha$ line. Scenarios in the IP and dwarf nova model are discussed (an emitting ring in the disc, bright X-ray spot on disc edge, or emitting accretion funnels), but the intermediate polar model is favoured. Supported is this idea by the finding that only the red peak appears to be shifted and the `blue' peak is compatible with the rest wavelength.  The red peak thus is caused by emission from the northern accretion spot when it faces the observer. Instead, the peak at the rest wavelength is caused when the southern accretion funnel is visible just on the lower edge of the white dwarf -- with the velocity of the accreting material being perpendicular to the line of sight. ",
        "introduction": "Cataclysmic variables (CVs) are close, interacting binaries, consisting of a white dwarf receiving matter from its companion, usually -- and so in our case -- a main sequence star. In non-magnetic systems, the matter is transfered via an accretion disc, while, if the white dwarf has a significant magnetic field, the inner part of the disc is disrupted (intermediate polar) or the formation of a disc is prevented altogether (polar). In both cases the transfered material couples to the magnetic field lines before falling onto the white dwarf. CVs have been found to emit radiation across the whole electromagnetic spectrum, as various energy sources are present within the system.  This paper focuses on {\\it XMM-Newton} X-ray observations of the dwarf nova T~Leo, i.e., we are looking at high energy sources within the system. These are usually the boundary layer (e.g., Wood et al.\\ 1995), a hot corona of the white dwarf (Mahasena \\& Osaki 1999), or accretion funnels and accretion spots on the surface of the white dwarf in magnetic systems (e.g., Schwarz et al.\\ 2002).  \\begin{figure*} \\centering \\includegraphics[width=8cm]{0846f01.ps} \\includegraphics[width=8cm]{0846f02.ps}  \\caption{OM (B and UVW1, left) and PN light curves at original resolution phased on the orbital period of 84~min. (Note, the data set is nearly exactly 2 orbital periods long, no repetition of data is plotted). } \\label{Figpnbom} \\end{figure*}  T~Leo is a very short period dwarf nova of SU~UMa type showing superoutbursts (Slovak, Nelson \\& Shafter 1987, Kato \\& Fujino 1987). With $\\Porb = 84$ min (Shafter \\& Szkody 1984) the orbital period is in fact close to the period minimum of about 76 min (e.g., Thorstensen et al.~2002) and makes T~Leo therefore particularly interesting. So far, T~Leo was only studied twice in detail in X-rays, once in superoutburst (with {\\it RXTE} by Howell et al.\\ 1999) and once in quiescence (with {\\it ASCA} by Szkody et al.\\ 2001).  Howell et al.\\ detected T~Leo only on one of five days during the outburst and on that day the X-ray emission was five times lower than during quiescence (compared with previous pointed {\\it ROSAT} PSPC observations by van Teeseling et al.\\ 1996 and Richman 1996). The timing of the detection coincides with a re-brightening of the system.  It seems to be a typical behaviour of dwarf novae that hard (3-12.2 keV) X-rays are suppressed during outburst as McGowan, Priedhorsky \\& Trudolyubov (2003) find also for SS~Cyg. They suggest that this lack of hard X-rays, originating in an optically thin disc corona, is caused by radiation driven disc winds stripping the coronal layers from the disc during an outburst (Meyer \\& Meyer-Hofmeister 1994).  Szkody et al.\\ (2001) find a variation on the orbital period for their quiescent {\\it ASCA} X-ray data: just before upper conjunction of the white dwarf the system brightens slightly. Otherwise the X-ray emission appears constant.  In this study we investigate spectral properties and the variability in the {\\it XMM-Newton} data of the quiescent T~Leo in order to learn more about the source of the X-ray emission. Hereby, two basic models are discussed to explain the behaviour of T~Leo: the intermediate polar model, as previously suggested by Shafter \\& Szkody (1984) or the dwarf nova model as otherwise assumed (e.g., Kato 1997). ",
        "conclusions": "\\subsection{Origin of the photometric variation} \\label{origin_phot}  \\begin{figure} \\centering \\includegraphics[width=8cm]{0846f10.ps} \\caption{{\\it ASCA} SIS0 light curve with overplotted sinusoidal fit of period 10.5 h. } \\label{Figasca} \\end{figure}   We discuss the observed frequencies in the light curve in order to investigate the nature of the system T~Leo.  In order to investigate if $P_x$ is caused by a beat period with $\\Porb$, we search for a variation on a timescale of about half a day. For this purpose we examined {\\it ASCA} data from 1998 (Szkody et al.\\ 2001) that were taken over a period of about 20 hours. Fig.~\\ref{Figasca} shows that there might indeed be a sinusoidal variation on a period of $P_y = 10.5$ h. However, the sinusoid with an amplitude of 0.042 counts s$^{-1}$ has only a marginally better $\\chi^2$ of 2.24 compared to a constant flux ($\\chi^2 = 2.62$) and the peak at $\\sim 13$~h might also be a flare.  However, apart from the fact that the period $P_x$ has an error range that is large enough to include the orbital period, it would be difficult to explain it. Neither a warped disc model (e.g., Patterson 1999) nor a jet model (Livio 1998) can explain the periods $P_x$ or $P_y$ reasonably well or can be seriously considered. A sudden and extreme decrease of the orbital period in the last couple of decades can also be excluded (Kolb 2003, private communication). Therefore, we consider $P_x$ to be identical to $\\Porb$.  For clarity we collected the definitions of all discussed periods in Table~\\ref{Tabperiods}.  In the next two Sections we discuss two models that could explain the observed period at 414~s: T~Leo is either an intermediate polar or a dwarf nova showing Quasi-Periodic Oscillations.  \\begin{table*} \\caption[]{Discussed periods. All but the first period refer to T~Leo.} \\label{Tabperiods} \\begin{tabular}{lrl} \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} Symbol{\\hspace{1ex}}  & Value & Explanation \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} $P_{\\SSS min}$ & $\\sim$ 76 min & observed period minimum for Cataclysmic variables (e.g.\\ Thorstensen et al. 2002)\\\\ $\\Porb$ & 84.69936$\\pm$0.00068 min & orbital period (Shafter \\& Szkody 1984) determined from radial velocity measurements of H$\\alpha$\\\\ $P_{\\SSS sh}$ & 86.7 $\\pm$ 0.1 min & super hump period (Lemm et al.~1993)\\\\ $P_x$ & $75^{\\SSS+12}_{\\SSS-8}$ min & $0.89^{\\SSS+0.14}_{\\SSS-0.10} \\Porb$ as determined from our EPIC PN data, see Fig.~\\ref{Figchi2}\\\\ $P_y$ & 10.5 h & determined from {\\it ASCA} data, see Fig.~\\ref{Figasca}\\\\ $\\Pspin$ & 414 s & spin period of white dwarf(?), see Fig.~\\ref{Figpower}\\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{table*}     \\begin{figure} \\centering \\includegraphics[width=7cm]{0846f11.ps} \\caption{The EPIC PN light curve folded on the period 0.1163~h = 414.42~s. The error bars are calculated from the scatter in the light curve. The light curve is plotted twice for clarity of the plot. } \\label{Figspin} \\end{figure}  \\subsubsection{Intermediate Polar model} \\label{spin} This scenario presumes that the white dwarf is slightly magnetic, just enough to disrupt the inner part of the disc up to no more than a few white dwarf radii (however, not to a complete absence of the disc, as (super) outbursts prove that a disc must be present). In such an intermediate polar (IP) the accreting material from the disc is forced to follow the magnetic field lines onto the two poles on the surface of the white dwarf.  The X-ray emission then originates from the accretion column close to the white dwarf (e.g., King 1995).  Such a scenario was already proposed by Shafter \\& Skzody~(1984). While the lack of optical polarisation can be easily explained with a weak magnetic field, it is more difficult to explain the lack of high excitation emission lines in the optical range. The lack of obvious orbital modulation in the optical and UV must then be related to a relatively low inclination angle of the system whose accretion disc mainly radiates at these long wavelengths. In contrast, the X-rays then originate to a large part from the accretion funnels and the accretion spots on the white dwarf which are much more prone to viewing angle changes than the disc, even in a low inclination system. It then seems likely that there is actually a small amount of variation visible in the {\\it UV} as suggested in Fig.~\\ref{Figpnaom}, because the accretion funnels will be hot and emit a bluer spectrum, thus be more pronounced in the UV than in the optical. However, in both the optical and UV the disc might be considerably brighter than the accretion funnel.  In IPs the magnetized white dwarf is usually desynchronized with a spin period of several hundred to thousand seconds. The peak in the power spectrum of the EPIC PN light curve at 414~s (Fig.~\\ref{Figpower}) could thus indeed be interpreted as the spin period of the white dwarf. In Fig.~\\ref{Figspin} we plot the EPIC PN light curve folded onto the 414~s period. The simultaneous, independent EPIC MOS data are consistent with the spin profile, however, the shape is less pronounced due to much higher noise level. In particular, the MOS2 data show a better agreement with the PN data than the MOS1 data.  The spin profile is clearly double peaked with a peak-to-peak separation of about 0.3 to 0.4 (definitely $< 0.5$) spin phases or roughly 1/3 of the orbit. Norton et al.\\ (1999) observed a similar situation in the intermediate polar V709 Cas. However, while they see a peak in the power spectrum at a frequency 3/$\\Pspin$, our power spectrum does not show any obvious peak at this frequency (3/$\\Pspin = 7.24\\times10^{-3}$~Hz). On the other hand, considering the relative strength of the 1/$\\Pspin$ and 3/$\\Pspin$ peaks in Norton et al.'s power spectrum of about 5:1, a 3/$\\Pspin$ signal of an amplitude of the order 0.025 units might be hidden in the profile of another peak at $7.17\\times10^{-3}$~Hz (see Fig.~\\ref{Figpower}).  A more detailed look reveals that in our case the full amplitude of the variation is only 10\\% of the signal (compared to 50\\% in V709 Cas). The system is brighter for about 2/3 of the orbit and shows a sharp drop after the second maximum. This {\\em could} indicate a partial occultation of the X-ray source.  We do not find any significant signal at the beat frequency ($1/\\Pspin - 1/\\Porb$), but possibly at ($1/\\Pspin - n/\\Porb$) with $n = 5,6$, and at ($1/\\Pspin - 1/P_x$) as well as ($1/\\Pspin - n/P_x$) with $n = 3,4,6,-1,-2,-3$. As Norton et al.\\ (1999) point out, the beat frequency is usually a signal of stream-fed or disc overflow accretion and therefore more likely be observed in IPs with stronger magnetic fields. For T~Leo this means it is one of the few systems that is disc-fed and seen at a relatively low angle. So far it is unclear what the observed frequencies represent, if they are significant.  In comparison to other IPs Norton et al.\\ (1999) note that such double peaked profiles appear to arise particularly in white dwarfs with short spin periods (of less than about 700 s). This appears to be very well fulfilled. They argue further that the double peaked spin profile is caused by the small magnetic field (as we propose for T~Leo's white dwarf). This either leads to large accretion areas and thus broad accretion curtains with optical depth reversal or (less likely because of the small inclination angle) accretion columns with conventional optical depth behaviour but that are tall enough to be partially visible from behind the white dwarf. In both scenarios the X-ray pulse maximum is seen whenever one of the poles is closest to the observer (see also Norton et al.`s Fig.~8).  Since the maxima in the light curve folded on the spin period are less than 0.5 spin phases apart, the geometry of the magnetic field must be somewhat asymmetric.  All of this is most striking as the observed variations are only visible in X-rays, while at optical wavelengths no periodic signal has been observed (Shafter \\& Szkody 1984) and our UV data show only a very marginal periodicity if at all.   \\begin{figure} \\centering \\includegraphics[width=7cm]{0846f12.ps} \\caption{The EPIC PN light curve folded on the period 20.567~s (bottom) and 22.112~s (top, shifted upwards by 1 counts s$^{-1}$ for clarity of the plot). } \\label{Figdno} \\end{figure}  \\subsubsection{Dwarf nova model} \\label{QPO}  An alternative model for T~Leo is that it is a dwarf nova showing Quasi-periodic Oscillations (QPOs) and Dwarf Nova Oscillations (DNOs, for a review on QPOs and DNOs see Warner 2004). Dwarf novae show often short-lived QPOs with a period of a few hundred seconds. In DNO-related QPOs one can typically see DNOs with a period ratio $P_{\\SSS QPO}/P_{\\SSS DNO} \\sim 15$. The frequencies around 21~s -- if significant-- could thus actually be the DNOs related to a 414~s QPO. As the nominal period deduced from the Keplerian velocity at the radius of the white dwarf is 16~s, it would mean that either the 20.567-s-signal arises from a disc annulus very close to the white dwarf or that the white dwarf has a high rotational velocity of 800 km s$^{-1}$.  A test for phase shifts between the first and second half of the data set does not help in deciding in favour of the QPO model as the measured phase shift of 0.1 is compatible with zero within the error range.  Problematic for this model is, however, that the lightcurve folded on any of the two ``DNO'' periods reveals a double peaked curve instead of the expected highly sinusoidal signal (Fig.~\\ref{Figdno}), although there might be exceptions (Warner 2004).  However, a number of dwarf novae have been found to show rapid oscillations in X-rays in quiescence usually of the order a few tens to a hundred seconds (Table 2 in Warner 2004). In particular, WZ~Sge is the only one showing persistent DNOs and QPOs in quiescence (Warner \\& Woudt 2002).  The question arises, whether T~Leo is similar to WZ~Sge. T~Leo certainly shows large amplitude outbursts, but is not generally considered a candidate for a WZ~Sge type dwarf nova as it also shows normal outbursts (Kato, Sekine \\& Hirata 2001).  WZ~Sge shows DNOs at 27.868~s and 28.952~s (mostly one of the two, but occasionally both simultaneously) and a period close to the beat period of the two DNOs at 742~s. These periods are explained as radiation from the central source with a rotation period of 27.868~s and a thickened region of the disc with a prograde rotation of period 744~s, the beat period of the DNOs, that reprocesses and obscures the radiation from the white dwarf (Warner \\& Woudt 2002).  In trying to apply this model to T~Leo, we first notice that close to the expected beat period of the 20.567~s and 22.112~s periods (294~s) is only a minor peak with a maximum at 297~s in the power spectrum. We therefore dismiss the idea that T~Leo displays DNO-related QPO and consider it likely that the frequencies around 21~s are not significant.  A more likely model, however, is that the QPOs in unrelated to any DNOs.  One option is that the QPO is either caused by a modulated mass transfer rate through the inner Lagrangian point due to oscillations of the secondary similar to the 5-min oscillations on the sun (Warner 2004).  Another possibility is that the source of the variation is in the outer disc (Warner 2004). The period of 414~s corresponds to the Keplerian velocity of the disc material at a distance of $9\\Rwd$. However, as T~Leo is a SU~UMa type dwarf nova and has therefore a rather large disc, this distance is likely at an intermediate radius, rather than at the disc edge (the 3:1 resonance radius $R_{3:1}$ is at $22\\Rwd$, see Section~\\ref{ring} for system parameters).  We can better imagine a model in which a travelling front originates close to the inner disc edge of a slightly disrupted disc. In this scenario, proposed by Warner \\& Woudt (2002), the white dwarf has a low magnetic field, just large enough to disrupt the very inner part of the accretion disc. The stream from the secondary overflows the disc and impacts onto the inner disc edge. This can excite a buldge travelling retrogradely in the rotating disc (in fact a low frequency prograde $m=1$ g-mode). Using Warner \\& Woudt's eq.~19 and 21 and the system parameters as described in Section~\\ref{ring}, the impact radius is between 3.1 and 3.8 $\\Rwd$ and the Keplerian period at this radius is between 160~s to 205~s (for mass ratios $q=1/4$ to $1/6$). Thus, if the buldge is close to the disc edge it has a speed between a third to half the velocity of the disc material.   \\subsection{Origin of the emission line profile}  \\label{explspec} The broadened, double peaked emission line we observe for O~VIII could be caused by a) an accretion disc or ring, similarly to double peaked optical line emission in high inclination systems; b) a spot on the disc edge of an otherwise X-ray dark disc; in the dwarf nova model or c) accretion curtains in which the matter from the inner disc radius falls down onto a magnetic white dwarf, in the intermediate polar model.  In the following section we investigate which of these scenarios apply to T~Leo.   \\subsubsection{Accretion Disc or Ring} \\label{ring}  The explanation that the O~VIII emission arises in an accretion disc seems unlikely since the optical emission lines show only a very marginal double peaked structure (Shafter \\& Szkody 1984), particularly not as clear as the separation in the O~VIII line. The inclination angle is with $i<65^\\circ$ (Shafter \\& Szkody 1984, Szkody et al.\\ 2001) too small to produce double peaks.  It is, however, noticeable that the separation of the peaks in the O~VIII line of $\\sim$950 km s$^{-1}$, i.e., $\\pm475$ km s$^{-1}$, fits to the velocities in the accretion disc in the Doppler Images of Shafter \\& Szkody.  It is, therefore, worth investigating if the emission in O~VIII might be originating in a ring of narrow radial width in the disc. Simple model calculations can produce a double peaked U-shaped profile. While the model profiles have significant flux at the system velocity of the line, the observed profiles appear more clearly separated. However, due to the large noise in the data, we cannot exclude that the emission line profiles are produced in this way.  For a Keplerian ring in the accretion disc we can estimate the radius from the velocity $v \\sin i$ and the white dwarf mass $\\Mwd$. The system parameters of T~Leo are quite uncertain, various values can be found in the literature. According to Shafter \\& Szkody (1984) the white dwarf has a mass $\\Mwd < 0.4 \\Msol$ and an inclination angle $i < 65^\\circ$. This leads to an upper value for the ring radius of $0.31 \\Rsol \\sim 31 \\Rwd$ (assuming $\\Rwd = 0.01 \\Rsol)$. The values derived by Belle et al. (1998) from model fitting of an IUE spectrum during an outburst of $\\Mwd = 0.6 \\Msol$ and $i = 40^\\circ$ give a ring radius of $0.23 \\Rsol \\sim 20 \\Rwd$ using their $\\Rwd$ of $8\\times10^8$ cm.  Using a secondary mass of $M_2 = 0.1 \\Msol$ (period-mass relation, Warner 1995) we obtain a mass ratio of $q = \\frac{1}{4}$ or $\\frac{1}{6}$, respectively. This leads to a distance from the white dwarf to the inner Lagrangian point $L_1$ in relation to the binary separation $a$ of $\\Rl/a = 0.65 \\pm 0.02$ (determined as the average of the $\\Rl/a$'s derived for the two values of $q$; Silber 1992, after Warner 1995). Using Warner's (1995) mass-radius relationship for the secondary gives a secondary radius of $R_2 \\sim 0.136 \\Rsol$. With this we can derive the absolute distance from the white dwarf to the inner Lagrangian point $\\Rl$ to $0.25 \\Rsol$ or 22$\\Rwd$, i.e., for the following discussion we use Belle et al.'s (1998) system parameters as their parameters allows the proposed ring to be inside the Roche-lobe of the white dwarf. However, the true value must be somewhat larger, as the 3:1 resonance radius $R_{3:1}$ is at also $22\\Rwd$, and during super outbursts the disc becomes larger than this. The true system parameters must therefore be slightly different from the quoted values.  \\begin{figure} \\centering \\includegraphics[width=8cm]{0846f13.ps} \\caption{O~VIII emission line with overplotted models for an emitting ring in the accretion disc. The model profiles were convolved with a Lorentzian of $0.04\\AA$ width. The ring is either narrow at a radius of $8\\Rwd$ ($0.37\\Rl$, dashed line) or has a radius range between maximally $4\\Rwd$ to $14\\Rwd$ ($0.18$ to $0.65\\Rl$, dotted line). The normalization of the model profiles are chosen simply for illustrative purposes.} \\label{Figprof} \\end{figure}   As an example we constructed a simple model of an emitting ring and calculated the emission line profile folded with a Lorentzian with the observationally given resolution of $0.04\\AA = 630$ km s$^{-1}$. We could only achieve a double peaked profile with the peaks at the observed velocities ($\\pm 475$ km s$^{-1}$) and the observed line width for either a very narrow ring ($< 1$\\%) at about $8\\Rwd$ ($0.37\\Rl$) or a broader ring between maximal $4\\Rwd$ and $14\\Rwd$ ($0.18$ to $0.65\\Rl$) (Fig.~\\ref{Figprof}). However, it is impossible to separate the two peaks clearly by assuming a ring structure.   This calculation shows that the ring model does not convincingly reproduce the observed results. Furthermore, the ring would have to be quite far from the hot inner disc ($>4\\Rwd$) and it is unclear why only there the O~VIII emission would be produced, as O~VIII has a maximum at a temperature of $3\\times10^6$~K. Such temperatures are only expected for the inner part of the disc (i.e., very close to the white dwarf).   In principle we can save this idea, if we assume that the disc is {\\em warped} in the inner part of the disc. This would mean that the hot, inner disc has a different and variable inclination angle than the rest of the disc. This could explain the lack of correlation between the gradients in the light curve and the presence of double peaks as well as the single peaked optical emission lines. When we see it edge on, it becomes double peaked, however, at times when we face the inner part of the disc, the emission line becomes {\\it narrow} not broad. However, this scenario is also problematic.   \\begin{figure*} \\centering \\includegraphics[width=8cm]{0846f14.ps} \\includegraphics[width=8cm]{0846f15.ps}  \\caption{Spectra during rising and falling (decreasing) phases ({\\em left}) and spectra during high and low phases ({\\em right}). The spectra for the decreasing branch phases and the high phases are shifted upwards by 2 counts for clarity of the plots. See Fig.\\ref{Figpnaom} for identification of phase ranges.}  \\label{Figreblhilo} \\end{figure*}   \\subsubsection{X-ray bright hot spot}  In order to investigate option b), we compared the timings of the spectra with those of the X-ray light curve. Hereby, we assume that the sinusoidal variation in the PN light curve is caused by the bright spot moving in and out of view. Then we can separate the light curve in a {\\em falling} portion, when the X-rays fade, and a {\\em rising} portion, when the X-ray brighten according to Fig.~\\ref{Figpnaom}. If the beaming hot spot theory is correct, we would expect that the spectra produced from timings of the {\\it falling} portion of the light curve only show a red peak in the O VIII emission line, while the other spectra created from timings of the {\\it rising} portions of the light curve should show only a blue peak. However, this is not the case (Fig.~\\ref{Figreblhilo}, left panel). While the ``{\\it blue}'' spectrum clearly shows a double peaked structure for the O VIII line, the ``{\\it red}'' spectrum simply shows a broadened emission line. Therefore, we have to dismiss this scenario.     \\subsubsection{Accretion curtains in an Intermediate Polar}  Eventually, we investigate, if (and how) the double peaked emission line can be produced within an IP model. This may include a scenario in which the white dwarf has a very low field as proposed for the QPO model (end of Section~\\ref{QPO}). This goes hand in hand with the question, if thus we should expect such {\\em emission line profiles} for all IPs with small magnetic fields (and double peaked {\\em spin profiles}).  In an IP the matter from the secondary is first fed into an accretion disc. This disc is disrupted at a certain inner radius from where the matter couples to the magnetic field lines of the white dwarf. This leads to accretion curtains through which the matter from the secondary is dumped onto the surface of the white dwarf.  Whenever the magnetic north pole (by definition the pole that is on the hemisphere closer to the observer) is facing the observer, the material in the accretion curtain is falling down onto the white dwarf, causing red shifted emission. However, at the timing of the second maximum the white dwarf has turned such that the material in the curtains moves perpendicular to the line of sight.  Even if the accretion columns are tall enough so that the southern one is visible at the time of the first maximum from behind the white dwarf we would not expect blue-shifted emission, as the X-rays are produced too close to the white dwarf to be visible and the curtain bends too strongly to see any emitting material falling in the direction towards the observer.  We can expect both blue and redshifted emission, if the poles are moving in and out of view due to the rotation of the white dwarf. However, the rotational velocity of the white dwarf (using Belle et al.'s (1998) radius of $8\\times10^8$~cm) is $2\\pi \\Rwd / \\Pspin = 121$~km s$^{-1}$, much less than the observed $\\pm 475$ ~kms$^{-1}$ in the emission line profile.  As mentioned in Section~\\ref{spectro} emission line peaks are not shifted symmetrically to the rest wavelength, where the blue peak is compatible with the rest wavelength and the red one shifted to a positive velocity of about $v = +630\\pm340$ km s$^{-1}$. This means, we see the red peak originating in the accretion funnel of the northern hemisphere in which the material is moving away from the observer onto the white dwarf (involving an optical depth reversal as proposed by Norton et al.~1999). The blue peak is not seen, as the infalling material with velocity towards the observer is hidden from the view behind the white dwarf. The peak at the rest wavelength comes from material moving perpendicular to the line of sight, e.g., at the phase when the southern pole is closest to the observer.  An analysis of spectra extracted during the phases of the two spin profile maxima and the single minimum ($0.5-0.925, 0.925-1.15, 0.15-0.5$, respectively, see Fig.~\\ref{Figspin}) seems to support this idea (e.g.\\ the emission during the broad maximum between spin phases 0.5 to 0.925 appears to show only the red peak), however, the count rates are too low to make any firm statements. Further observations with a more secure wavelength calibration and better S/N ratio are necessary to confirm this model.  The IP model for T~Leo was first suggested by Shafter \\& Szkody (1984). In order to explain their measured radial velocity curve shift Warner (1995) proposes that the accretion stream from the secondary overflows the accretion disc and hits the inner disc edge similarly as in EX~Hya. Shafter \\& Szkody's observed line width of H$\\alpha$ in fact is compatible with an inner disc radius of 6-7 $\\Rwd$. However, as the disc is much cooler than the acretion funnels we do not expect any implication of the stream overflow model on our X-ray data.   T~Leo is not the only suspect for an SU~UMa type dwarf nova showing characteristics of an IP, however, it is the strongest candidate. SW UMa also shows superoutbursts with the typical features like superhumps (Robinson et al. 1987) and Shafter, Szkody \\& Thorstensen (1986) discovered a period of 15.9 min that could be the rotation period of the magnetic white dwarf. However, Rosen et al.\\ (1994) could not confirm the 15.9~min period with their data.  While the superoutbursting (Kato \\& Nogami 1997) VZ Pyx has been suspected to be an IP (de Martino et al.\\ 1992, Remillard et al. 1994), Szkody \\& Silber (1996) point out, this system might not be an IP, as it shows coherent pulses only during decline from outburst. In contrast, our spin detection was made outside of any outburst. Furthermore, Warner, Woudt \\& Pretorius (2003) find QPO and DNO signals in their photometry of VZ~Pyx, excluding the possibility that it is an IP. Another candidate is HT Cam (RX J0757.0+6306) which shows oscillations (Kemp et al.\\ 2002), but whose SU~UMa nature is not confirmed, yet (Tovmassian et al.~1998).  Further objects listed in Warner's (1995) Table 7.2 being SU~UMa type dwarf nova and simultaneously IP candidates are AL~Com and HT~Cas. However, while both are clearly SU~UMa stars showing super outbursts, there is no clear evidence that HT~Cas is an IP (but possibly a VY Scl type object, Robertson \\& Honeycutt 1996), and AL~Com shows only some properties similar to EX~Hya (Abbott et al. 1992). All in all both are rather unusual objects, possibly linked to the WZ~Sge-type dwarf novae."
    },
    "0402/astro-ph0402109_arXiv.txt": {
        "abstract": "\\hspace{5.3cm}{\\bf ABSTRACT}  \\vspace{0.1cm} A large number of recent observational data strongly suggest that we live in a flat, accelerating Universe composed of $\\sim$ 1/3 of matter (baryonic + dark) and $\\sim$ 2/3 of an exotic component with large negative pressure, usually named {\\bf Dark Energy} or {\\bf Quintessence}. The basic set of experiments includes: observations from SNe Ia, CMB anisotropies, large scale structure, X-ray data from galaxy clusters, age estimates of globular clusters and old high redshift galaxies (OHRG's). Such results seem to provide the remaining piece of information connecting the inflationary flatness prediction ($\\Omega_{\\rm{T}} = 1$) with astronomical observations. Theoretically, they have also stimulated the current interest for more general models containing an extra component describing this unknown dark energy, and simultaneously accounting for the present accelerating stage of the Universe. An overlook in the literature shows that at least five dark energy candidates have been proposed in the context of general relativistic models. Since the cosmological constant and rolling scalar field models have already been extensively discussed, in this short review we focus our attention to the three remaining candidates, namely: a decaying vacuum energy density (or  ${\\bf \\Lambda(t)}$ {\\bf models}), the {\\bf X-matter}, and the so-called {\\bf Chaplygin-type gas}. A summary of their main results is given and some difficulties underlying the emerging dark energy paradigm are also briefly examined. ",
        "introduction": "\\vspace{0.1cm} In 1998, some results based on Supernovae (SNe) type Ia observations published independently by two different groups, drastically changed our view about the present state of the universe \\cite{perlmutter,riess}. In brief, the Hubble-Sandage diagram describing the observed brightness of these objects as a function of the redshift lead to unexpected and landmark",
        "conclusions": "\\vspace{0.1cm} The search for cosmologies driven by dark energy is presently in vogue. The leitmotiv is the observational support for an accelerated Universe provided by the type Ia supernovae (SNe) experiments at intermediate and high redshifts.  This short review focused on some alternative candidates to dark energy. This ubiquitous component plus the dark matter are responsible for nearly 95\\% of the matter-energy content filling the Universe. However, different from dark matter, the extra dark (energy) component is intrinsically relativistic and its negative pressure is required by the present accelerating stage of the Universe. Its tiny density and weak interaction presumably preclude the possibility of identification in the terrestrial laboratory. Unfortunately, even considering that we are in the golden age of empirical cosmology, the existing data are still unable to discriminate among the different dark energy candidates, thereby signaling that we need better observations in order to test the basic predictions. This means that the determination of cosmological parameters will continue to be a central goal in the near future. The fundamental aim is to shed some light on the nature of the dark energy, but it is not clear if it can be revealed using background tests with basis only in a different equation of state. Another possibility is to add some hypothesis concerning the nature itself (is it formed by massive or massless particles?), and to follow examining its consequences. It is also worth notice that the energy of this relativistic dark component grows in the course of an adiabatic expansion. Macroscopically, the energy increases on account of the thermodynamic work done on the system (negative pressure). This intriguing behavior is in marked contrast to what happens to any component with positive or null pressure. Naturally, a possible microscopic explanation for such a fact is of great interest because it depends on the intrinsic nature of the dark energy, and may also have important consequences to the ultimate fate of the Universe.  On the other hand, since the current models are more complicated than the Einstein-de Sitter Universe, such a situation is somewhat uncomfortable either from theoretical or observational viewpoints. It has also to be admitted that none dark energy model has been successful enough to deserve the status of a ``standard model\". However, the present time for many cosmologists is very exciting because although preserving some aspects of the basic physical picture, the new invisible actor (dark energy) which has not been predicted by the standard model of Particle Physics, and is responsible for repulsive gravity, may alter profoundly the traditional view of space-time and matter.  \\vspace{0.05cm} \\vspace{0.1cm}{\\bf Acknowledgments:} This work was partially supported by the CNPq (62.0053/01-1-PADCT III/Milenio) and  FAPESP (00/06695-0). I am also indebted to G. Steigman, J. S. Alcaniz, N. Pires, J. Santos, R. Silva, J. V. Cunha and R. C. Santos for many helpful discussions. \\vspace{0.01cm}"
    },
    "0402/astro-ph0402423_arXiv.txt": {
        "abstract": "We present Keck LWS images of the Orion BN/KL star forming region obtained in the first multi-wavelength study to have $0\\rlap{.}''3$-$0\\rlap{.}''5$ resolution from $4.7\\mu$m to $22\\mu$m. The young stellar objects designated infrared source\\,{\\it n} and radio source\\,{\\it I} are believed to dominate the  BN/KL region.  We have detected extended emission from a probable accretion disk around source\\,{\\it n} but infer a stellar luminosity on the order of only 2000\\,L$_\\odot$.  Although source\\,{\\it I} is believed to be more luminous, we do not detect an infrared counterpart even at the longest wavelengths.   However, we resolve the closeby infrared source, IRc2, into an arc of knots $\\sim 10^3$\\,AU long at all wavelengths.  Although the physical relation of source\\,{\\it I} to IRc2 remains ambiguous, we suggest these sources mark a high density core ($10^7$-$10^8$\\,pc$^{-3}$ over $\\sim10^3$\\,AU) within the larger BN/KL star forming cluster. The high density  may be a consequence of the core being young and heavily embedded.  We suggest the energetics of the BN/KL region may be dominated by this cluster core rather than one or two individual sources. ",
        "introduction": "Orion BN/KL is the closest region of high-mass star formation ($\\sim 450$\\,pc) and though also the archetype, it is not necessarily well understood \\citep[e.g.,][]{genzel89}.  For instance,  the sources that drive extensive outflow activity in BN/KL \\citep[e.g.,][]{wright96}  and contribute most to its luminosity have not been identified with certainty \\citep[e.g.,][]{mr95}.  This is true because BN/KL is crowded, exhibiting $\\sim 20$ compact infrared (IR) peaks distributed over $\\sim 20''$ and because extinction is high and variable across the region \\citep{gezari98}.  As high-mass stars often form in clusters \\citep[][and references therein]{gl99},  there are presumably numerous embedded protostars and young stellar objects (YSOs) in the region, many of which have probably not been recognized yet.  Two notable sources in BN/KL exhibit centimeter-wave continuum and as a result are believed to be luminous YSOs, IR source\\,{\\it n}  \\citep{lonsdale82} and radio source\\,{\\it I} \\citep{churchwell87}. Source\\,{\\it n}, also known as  radio source\\,{\\it L} \\citep{churchwell87}, exhibits a bipolar radio jet \\citep{mr95}, hard X-ray emission \\citep{garmire00,feigelson02}, and $\\lambda 2.3\\mu$m CO overtone emission probably from the hot surface of an accretion disk \\citep{luhman00}.  Source\\,{\\it I} has been difficult to study because it has not been detected at IR wavelengths, probably due to high extinction. It lies close to the mid-IR source IRc2 \\citep{dyg92,mr95} for which $A_V\\sim 60$ \\citep{gezari98} and the peak of the Orion Hot Core, for which probably $A_V\\gg10^2$ \\citep{genzel89,migenes89,plambeck95}. However, source\\,{\\it I} excites H$_2$O and SiO maser emission in dense, hot material at radii $<10^3$\\,AU \\citep{genzel89,dyg92}, which may be used to study gas dynamics in detail \\citep[e.g.,][]{greenhill98,greenhill03}.  We have imaged BN/KL in the mid-IR at resolutions of $0\\rlap{.}''3$-$0\\rlap{.}''5$  with the intention to characterize sources of excitation.  This is the first multiwavelength study with sub-arcsecond resolution up to $22\\mu$m. \\citep[cf.,][]{dougados93,rouan04,shuping04}.  We concentrate on a largely qualitative interpretation of the data. Quantitative modeling and analysis of the whole region will be discussed in a separate paper. ",
        "conclusions": "We have detected the probable signature of an accretion disk around source\\,{\\it n} at mid-IR wavelengths, and we estimate the stellar luminosity is $\\sim 2000$\\,L$_\\odot$.  We have not identified any counterpart to radio source\\,{\\it I}, despite higher angular resolution and broader wavelength coverage than previous studies.  We resolve IRc2 into knots at up to $\\lambda22.0\\mu$m and propose that they are sites of star formation.  Together, source\\,{\\it I} and IRc2 may mark the dense core ($10^7$-$10^8$\\,M$_\\odot$\\,pc$^{-3}$) of the larger BN/KL infrared cluster. Many past studies have sought to demonstrate that luminosity and outflow associated with BN/KL are driven by one or two sources.  This seems increasingly unlikely. We note that source\\,{\\it n} has only a modest luminosity and its jet axis is poorly aligned with outflow axes observed on scales of $10^4$\\,AU \\citep[e.g.,][]{schultz99}.  Moreover, members of the putative IRc2-source\\,{\\it I} subcluster are ready candidates to drive outflow collectively.  In a later paper, we will present a radiative model of the LWS data, with estimates of dust temperatures, line-of-sight optical depths, and intrinsic luminosities for the dominant mid-IR sources in BN/KL."
    },
    "0402/astro-ph0402490_arXiv.txt": {
        "abstract": "We study the big-bang nucleosynthesis (BBN) scenario with late-decaying exotic particles with lifetime longer than $\\sim 1$ sec.  With a late-decaying particle in the early universe, predictions of the standard BBN scenario can be significantly altered. Therefore, we derive constraints on its primordial abundance.   We pay particular attention to hadronic decay modes of such particles.   We see that the non-thermal production process of D, ${\\rm ^{3}He}$ and ${\\rm ^6Li}$ provides a stringent upper bound on the primordial abundance of late-decaying particles with hadronic branching ratio. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0402/astro-ph0402173_arXiv.txt": {
        "abstract": "Part of the CMB polarization signal in the direction of galaxy clusters is produced by Thomson scattering of the CMB temperature quadrupole. In principle this allows measurement of the CMB power spectrum harmonic $C_2(z)$ with higher accuracy (at $z>0$) than the cosmic variance limit imposed by sample variance on one CMB sky. However the observed signals are statistically correlated if the comoving separation between the clusters is small enough. Thus one cannot reduce the sample variance by more than roughly the number of separate regions available which produce uncorrelated signals, as first pointed out by Kamionkowski and Loeb. In this paper we analyze in detail the procedure outlined by Kamionkowski and Loeb,  computing the correlation of the polarization signals by considering the variation of the spherical harmonic expansion coefficients of the temperature anisotropy on our past light cone.  Given a hypothetical set of Stokes parameter measurements of the CMB polarization in the directions of galaxy clusters, distributed at random on a given redshift shell, we show how to construct an estimator of the angular power spectrum harmonic $C_2$ at that redshift. We then compare the variance of this estimator with the cosmic variance of the CMB multipole on our sky which probes the same scale. We find that in fact the cosmic variance is not reducible below the single sky CMB value using the cluster method. Thus this method is not likely to be of use for reconstruction of the primordial power spectrum. However the method does yield a measurement of $C_2$ as a function of redshift with increasing accuracy at higher redshift, and thus potentially a probe of the mechanism which may have suppressed the quadrupole.  We also examine to what extent the redshift dependence of $C_2$ can be used to probe the time changing potential anisotropy as the universe evolves into the vacuum dominated phase (the late-time integrated Sachs-Wolfe  effect). We find that this effect is not observable in the time dependence of $C_2$ since it is swamped by cosmic variance, but there is an observable signature in the correlation functions of the Stokes parameters. ",
        "introduction": "The CMB radiation incident on galaxy clusters has an intrinsic intensity quadrupole $Q_2$ created by inhomogeneity at the surface of last scattering. Thomson scattering of the CMB in a galaxy cluster with typical line of sight optical depth $\\tau_{\\rm C}$ generates polarization of order $Q_2 \\tau_{\\rm C}$. Thus a measurement of this polarization signal would allow an estimate of the CMB quadrupole at non-zero redshift. This is of interest because it would potentially allow us to get around the restriction of cosmic variance. To elaborate, at $z=0$ we only have one CMB sky to observe, with $(2l+1)$ independent real data points for each spherical harmonic mode of the CMB on our sky, to compare with the ensemble average prediction of the variance. There is thus an intrinsic fractional sample variance of the harmonic $C_l$ of $2/(2l+1)$ (see section \\S\\ref{ch5:gencorr}), which severely limits comparison with the ensemble averaged theory at low $l$. This restriction limits the accuracy of measurements of the primordial power spectrum on the largest scales. The theoretical predictions thus obtained for the CMB power spectra are fundamentally limited by this sample variance, commonly termed the \\emph{cosmic variance}.  Thomson scattering of the $l=2$ part of the CMB anisotropy in a cluster generates a secondary polarization anisotropy which depends on the spherical harmonic components $a_{2m}$ as seen by a (hypothetical) observer at the cluster. Since this polarization signal produced by a cluster is sensitive to the density perturbations on a last scattering surface different to our own, this in principle allows one to make more accurate comparison to the theoretical predictions for CMB angular power spectra at low $l$ than allowed by the cosmic variance limit. However the observed signals are correlated if the comoving separation between the clusters is small, and many strongly correlated signals are no more useful for reducing the sample variance than one signal. The variance in the estimated quadrupole can be reduced by roughly the number of regions available which produce uncorrelated signals.  This method of using the CMB polarization signal produced by galaxy clusters to get around cosmic variance limits was first pointed out by Kamionkowski and Loeb \\cite{Kamionkowski:1997na} (we usually refer to it as the ``cluster method'' in what follows). However Kamionkowski and Loeb did not actually compute the correlation of the cluster signals in a particular cosmological model in their paper, and did not therefore demonstrate explicitly that the cosmic variance is reduced with a given set of clusters, nor did they develop any formalism for converting measurements of the Stokes parameters of the CMB to statistical estimators which get around cosmic variance. In \\cite{2003PhRvD..67f3505C}, estimators of $C_2(\\tau)$ were constructed (taking into account the kinematic SZ contamination of the polarization signal also), but the effect of statistical variation in the polarization signal on the estimator variance was not included. This variation was considered by \\cite{2000PhRvD..62l3004S}, but they computed the variation of the quadrupole as an expansion in small cluster separations, and their analysis is not applicable to a general set of clusters in arbitrary locations. In this paper we compute the correlation of the cluster signals in the case of an idealized set of measurements from clusters distributed in random directions on a given redshift shell. We describe an explicit procedure for carrying out the program outlined in  \\cite{Kamionkowski:1997na}, and study how the correlations die off  as clusters of increasingly high redshift are used. Information about the correlation of the polarization signals is contained in the generalized correlation functions of the CMB temperature anisotropy coefficients, $\\langle a_{lm}(\\mbox{\\boldmath $x$},\\tau) a^*_{l^{\\prime}m^{\\prime}}(\\mbox{\\boldmath $x$}^{\\prime},\\tau^{\\prime}) \\rangle$, which contain all of the statistical information (assuming Gaussianity) about the variation of the $a_{lm}$ coefficients as the observation point and associated last scattering surface change. With these functions, we can derive an estimator for $C_2(z)$ in terms of Stokes parameters, and find its variance.  We should clarify exactly what we mean by reducing cosmic variance. It is true that the polarization signals provide an estimator \\( \\hat C_2(\\tau) \\) of the remote quadrupole at a given redshift which has a smaller fractional cosmic variance than the local quadrupole. However, this is not the most useful comparison, since this estimator probes smaller physical scales than the local quadrupole. Cosmologists already have estimators of the power on these scales, namely the WMAP angular power spectrum harmonics $C_l$ with $l>2$. So the interesting question to ask is whether \\( \\hat C_2(\\tau) \\) has smaller cosmic variance than \\emph{the CMB multipole on our sky that probes the same physical scale}. This determines whether or not the cluster technique is capable in principle of providing a better reconstruction of the primordial potential than the WMAP data. The results are presented in \\S\\ref{ch5:sec_QUstat}.  We now outline the organization of the paper. In \\S\\ref{ch5:gencorr} we derive the two-point generalized correlation functions of the spherical harmonic coefficients, assuming a Gaussian primordial perturbation spectrum. In \\S\\ref{ch5:sec_transfer} we discuss the CMB transfer functions used to compute the generalized correlation functions, and examine the time dependence of $C_2(\\tau)$. In \\S\\ref{ch5:quadscatt} we derive expressions for the Stokes parameters $Q, U$ (defined in an appropriate all-sky basis) of the CMB radiation scattered into the line-of-sight by the cluster gas, in terms of the local $a_{lm}$ at the cluster, for a general line-of-sight. Note that in this section we found it convenient to use the ``density matrix'' formalism for polarization calculations \\citep{1994PhDT........24K,2000PhRvD..62d3004C}, which is outlined in Appendix \\ref{appb}. Then in \\S\\ref{ch5:sec_QUstat} we consider the the statistical variation of these Stokes parameters with the comoving position of the cluster. We construct a simple estimator for $C_2(\\tau)$ and compute its variance for a number of simulated sets of clusters. In \\S\\ref{ch5:sec_discuss} there is a discussion and summary. Note that we restrict the discussion to the case of a flat FRW universe throughout, for simplicity. ",
        "conclusions": "}  We have developed a statistical theory of the part of the polarization signal in the CMB in the direction of galaxy clusters produced by scattering of the CMB temperature quadrupole. We have shown explicitly that it is possible to use the indirect information about the last scattering surfaces of distant observers contained in these polarization measurements to constrain the $l=2$ angular power spectrum harmonic of the CMB, $C_2$,  as a function of redshift with greater statistical accuracy. We also showed however that it not possible to use the cluster polarization measurements to probe the power on a given scale with higher accuracy than the limits imposed by cosmic variance on the single sky CMB data. Thus power spectrum estimation cannot be improved using this method.  But we believe the cluster method is still of considerable value though, since it serves as a probe of the physical mechanism which might have suppressed the quadrupole. It has been noted that the quadrupole seems to be lower than might be expected due to a statistical fluctuation alone \\citep{deOliveira-Costa:2003pu}. The quadrupole may be anomalously low, that is lower than the standard models predict, for various reasons: a cutoff at large scales in the fluctuation power spectrum, different-than-expected behavior of the transfer function at large scales, or effects associated with the large-scale topology or geometry of the universe. There could conceivably be other explanations, but these are the simplest. Signatures of these effects could be present in the time evolution of $C_2$. For example if the quadrupole is low because of some topological suppression, we should see a rise in $C_2$ with redshift as the scale probed falls below the local quadrupole scale (but if the standard models are correct, there would be no such rise, at least in the $n=1$ case).  One might be worried that since the accuracy of the cluster measurement of $C_2$ is not higher than the accuracy of the measurement of the corresponding WMAP harmonic, as shown is \\S\\ref{ch5:sec_QUstat}, the cluster method may not provide any more information that that already contained in the WMAP data. However the cluster technique yields information which is not obtainable with WMAP (about perturbations on last scattering surfaces different to our own) and is thus complementary. In future work it would be desirable to perform a full analysis of the feasibility of using the cluster method to detect the effects of topology (or other quadrupole suppression mechanisms) in the time evolution of $C_2$, and a comparison with what we can already learn from the WMAP data.  We also showed that in the standard $\\Lambda$CDM model the ISW effect produces a small ($\\approx 2$\\%) bump in CMB harmonic $C_2(\\tau)$, which is swamped by the high cosmic variance at low redshift even using large numbers of clusters. However the ISW effect produces a significant feature in the two-point correlation function of the Stokes parameters which might be detectable. Detection of the ISW effect would provide additional information about the acceleration of the universe and the dark energy. We also note that this method is a rather sensitive probe of deviations from the scale-invariant power spectrum, because if $n\\ne 1$ then the Sachs-Wolfe contribution to $C_2(\\tau)$ either grows or decays rapidly with conformal time.  The procedure we have outlined is something of an idealization. We assumed that the polarization signal induced by the quadrupole is obtainable from many clusters at the same redshift, and we ignored noise and contamination of the signal. In practice there are contaminating polarization signals from the kinetic and thermal Sunyaev-Zeldovich effects \\citep{Diego:2003dp}, and distortions in the polarization field due to lensing \\citep{2000ApJ...538...57S}. Clearly separating the quadrupole signal from the other contaminants would be a major experimental challenge. However the signal to noise may be increased to some extent by combining signals from clusters located at similar directions and redshifts, since the signal from sufficiently nearby clusters is strongly correlated. It remains to be seen if this technique will be a practically useful cosmological probe."
    },
    "0402/astro-ph0402459_arXiv.txt": {
        "abstract": "We have measured the present accretion rate of roughly 800 low-mass ($\\sim1-1.4~M_\\sun$) pre-Main Sequence stars in the field of Supernova~1987A in the Large Magellanic Cloud (LMC, $Z\\simeq 0.3\\ Z_\\sun$). It is the first time that this fundamental parameter for star formation is determined for low-mass stars outside our Galaxy. The Balmer continuum emission used to derive the accretion rate positively correlates with the \\ha\\ excess. Both these phenomena are believed to originate from accretion from a circumstellar disk so that their simultaneous detection provides an important confirmation of the pre-Main Sequence nature of the \\ha\\ and UV excess objects, which are likely to be the LMC equivalent of Galactic Classical T~Tauri stars. The stars with statistically significant excesses are measured to have accretion rates larger than $\\sim 1.5\\times10^{-8}M_\\sun\\ yr^{-1}$ at an age of 12-16~Myrs. For comparison, the time scale for disk dissipation observed in the Galaxy is of the order of 6~Myrs. Moreover, the oldest Classical T~Tauri star known in the Milky Way (TW~Hydr\\ae, with 10~Myrs of age) has a measured accretion rate of only $5\\times 10^{-10}~M_\\sun/yr$, \\ie 30 times less than what we measure for stars at a comparable age in the LMC. Our findings indicate that metallicity plays a major role in regulating the formation of low-mass stars. ",
        "introduction": "} The processes at play during star formation determine much of the appearance of the visible Universe. The shape of the stellar Initial Mass Function (IMF) and its normalization (the star-formation rate) are, together with stellar evolution theory, key ingredients in determining the chemical evolution of a galaxy and its stellar content. Yet, our theoretical understanding of the processes that lead from diffuse molecular clouds to stars is still very tentative, as many complex physical phenomena concur in producing the final results. While clear variations in the star-formation rate are observed in different regions of the Milky Way and in external galaxies, with their histories showing bursts and lulls \\citep[e.g.,][]{tol00}, the observational evidence for (or against) variations in the IMF is often contradictory \\citep[see the review of] [or the Gilmore vs Eisenhauer debate in ``Starbursts: Near and Far', 2001]{scalo98}. Yet, variations in the IMF can dramatically alter the chemical evolution of a galaxy \\citep[e.g.,][]{wyse98}.  From an observational standpoint, most of the effort has traditionally been devoted to nearby Galactic star-forming regions, such as the \\objectname{Taurus-Auriga} complex, \\objectname{Orion}, etc. If this, on the one hand, permits one to observe very faint stars at the best possible angular resolution, on the other it is achieved at the expense of probing only a very limited set of initial conditions for star formation \\citep[all these clouds have essentially solar metallicity, e.g.,][]{padg96}.  Studying the effects of a lower metallicity on star formation is also essential to understand the evolution of both our own Galaxy, in which a large fraction of stars were formed at metallicities below solar, and what is observed at high redshifts. As a matter of fact, the global star formation rate appears to have been much more vigorous (a factor of 10 or so) at $z\\simeq 1.5$ than it is today \\citep[][and subsequent incarnations of the so-called ``Madau plot'']{mad96}.  At that epoch the mean metallicity of the interstellar gas was similar to that of the \\objectname{Large Magellanic Cloud} at present \\citep[LMC, e.g.][]{pei99}. This fact makes the study of star forming regions in the LMC especially important for the understanding of galaxy evolution.  With a distance modulus of $18.57\\pm0.05$ \\citep[see the discussion in][]{rom00}, the LMC is our closest galactic companion after the \\objectname{Sagittarius} dwarf galaxy. At this distance one arcminute corresponds to about 15~pc and, thus, one pointing with a typical imaging instrument comfortably covers almost any star forming region in the LMC \\citep[10~pc see, e.g.,][]{hod88}. In particular, the field of view of about $2.7\\arcmin\\times2.7\\arcmin$ of the WFPC2 on board the HST corresponds to $37~\\mathrm{pc}\\times37~\\mathrm{pc}$, and leads to the detection of several thousands of stars per pointing \\citep[e.g., ][]{rom02}. The LMC is especially suited for stellar populations studies for two additional reasons. First, the depth of the LMC along the line of sight is negligible, at least in the central parts we consider \\citep{mar01}. All of the stars can, then, effectively be considered at the same distance, thus eliminating a possible spurious scatter in the Color-Magnitude Diagrams. Second, the extinction in its direction due to dust in our Galaxy is low, about $\\ebv\\simeq 0.05$ \\citep{bes91,sch91} and, hence, our view is not severely obstructed.  There is currently a widespread agreement that low mass stars form by accretion of material until their final masses are reached \\citep[e.g.][and references therein]{bon01}. As a consequence, the accretion rate is arguably \\emph{the} single most important parameter governing the process of low-mass star formation and its final results, including the stellar Initial Mass Function.  Ground and HST-based studies show that there may be significant differences between star formation processes in the LMC and in the Galaxy. For example, \\citet{lam99} and \\citet{dewit02} have identified by means of ground-based observations high-mass pre-Main Sequence stars (Herbig AeBe stars) with luminosities systematically higher than observed in our Galaxy, and located well above the ``birthline'' of \\citet{ps91}. They attribute this finding either to a shorter accretion timescale in the LMC or to its smaller dust-to-gas ratio. Whether such differences in the physical conditions under which stars form will generally lead to differences at the low mass end is an open question, but \\citet{pan00} offer tantalizing evidence of a higher accretion also for LMC low mass stars.  In this paper we present the first measurement of the accretion rate onto low-mass pre-Main Sequence stars outside of our Galaxy. The data and the reduction process are presented in the next section, while the detection of Balmer continuum excess and its conversion into an accretion rate are described in section~\\ref{sec:measuring_acc}. Section~\\ref{sec:pms_age} is devoted to deriving the stellar parameters of our sample, \\ie the stars' masses and ages. Finally, the conclusions are drawn in section~\\ref{sec:sum}. ",
        "conclusions": "} We have identified about 800 stars with statistically significant Balmer continuum excess in the field of SN~1987A in the Large Magellanic Cloud. This excess positively correlates with \\ha\\ emission, as derived from the comparison of broad and narrow-band photometry. Whereas the quality of the data does not allow to derive the detailed \\emph{shape} of the correlation, the existence of the correlation itself is proven with a very high statistical significance (Spearman's test returns a probability of less than $10^{-4}$ that the two variables are uncorrelated). Both the Balmer continuum and the \\ha\\ excesses are well above the levels expected from chromospheric activity. These facts lead us to interpret these objects as pre-Main Sequence stars.  In this framework, given their location in the HR diagram and the emission in the Balmer continuum and \\ha\\ line, the objects in our sample are the equivalent of Galactic Classical T~Tauri stars. This is the first time that such a fundamental parameter as the accretion rate is measured for this class of objects in a galaxy other than our own Milky Way. Doing so in the Large Magellanic Cloud provides a unique opportunity to sample astrophysical conditions not represented in local star-forming regions, but that were common in the early Universe.  When interpreted as pre-Main Sequence stars, the comparison of the objects' location in the HR diagram with theoretical evolutionary tracks allows one to derive their masses ($\\sim1-1.4~M_\\sun$) and ages ($\\sim12-16$~Myrs). At such an age and with an accretion rate in excess of $\\sim 1.5\\times10^{-8}M_\\sun\\ yr^{-1}$, these candidate pre-Main Sequence stars in the field of SN~1987A are both older and more active than their Galactic counterparts known to date. In fact, the overwhelming majority of T~Tauri stars in Galactic associations seem to dissipate their accretion disks before reaching an age of about 6~Myrs \\citep{hai01,arm03}. Moreover, the oldest Classical T~Tauri star know in the Galaxy, \\objectname{TW~Hydr\\ae} at an age of 10~Myrs, \\ie comparable to that of our sample stars, has a measured accretion rate some 30 times lower than the stars in the neighborhood of SN~1987A \\citep{muz00}.  The situation is summarized in Figure~\\ref{fig:muz}, adapted from \\citet{muz00}, where we compare the position in the age-$\\dot{M}$ plane of the stars described in this paper with that of members of Galactic star-forming regions. An obvious selection bias that affects our census is that we only detect those stars with the largest Balmer continuum excesses, \\ie highest accretion rates. There might be stars in the field with smaller accretion rates, either intrinsically or because they were observed when the accretion activity was at a minimum, which fall below our detection threshold. This selection effect is rather hard to quantify, but it is clear that the locus of the accreting stars that we do detect in the neighborhood of SN~1987A is significantly displaced from the one defined by local pre-Main Sequence stars.  There are essentially two ways to reconcile the position of the LMC point in Figure~\\ref{fig:muz} with the mean locus observed in the Galaxy: significantly reduce the age of the stars in the SN~1987A region and/or significantly decrease their inferred accretion rate. Neither option, however, seems viable. In fact, while the uncertainties on either quantity can be quite large for any given star, the LMC point reflects the mean values of several hundred stars, thus making the random errors, for all practical purpose, negligible. Let us now consider the systematic errors.  If we move the LMC point \\emph{horizontally} in Figure~\\ref{fig:muz}, in order to occupy the same area as Galactic stars, the LMC stars would have to be younger than about 4~Myrs or less, the age of the oldest Galactic star with a measured accretion rate of $1.5\\times10^{-8}M_\\sun\\ yr^{-1}$. This would require our estimate of the ages of the LMC stars to be systematically wrong by a factor of three or more. Such a large shift is rather implausible on accounts of two main considerations. First, \\citet{rom02} have shown that the dereddening technique we have adopted here does not introduce any significant systematic errors on the stellar parameter, hence, given a set of pre-Main Sequence evolutionary tracks, on the derived age. Second, the age of 14~Myrs inferred from the peak of the histogram in Figure~\\ref{fig:age-mass} agrees very well both with that of the most massive stars in the field \\citep[$12\\pm2$~Myrs,][]{pan00} and that of the progenitor of SN~1987A \\citep[10-12~Myrs see, for example,][and references therein]{scud96}.  These age estimates are, of course, completely independent from that of the low-mass stars we are discussing here, lending support to it.  The other possibility is that, given the age of about 12-16~Myrs, we have overestimated the accretion rate onto the candidate pre-Main Sequence stars in the surroundings of SN~1987A. A factor of about 30 reduction is required in order to make our value for the LMC to agree with TW~Hydr\\ae, the oldest accreting T~Tauri known in our Galaxy (see Figure~\\ref{fig:age-mass}). According to equation~(\\ref{eq:mdot_ef336w}) this corresponds to reducing the Balmer continuum excess $L_{\\mathrm{F336W},exc}$ by a factor of almost 20 (for comparison, this would be so small as to be even smaller than the bin size of the histogram in Figure~\\ref{fig:huflex}). Such a reduction, however, is clearly incompatible with the observed distribution of the excesses, in particular with its asymmetry, which marks the stars with a Balmer continuum excess beyond statistical errors (see also Figure~\\ref{fig:hubex}).  A potential source of concern is the effect on the derived accretion rate of the presence of circumstellar dust in the immediate surroundings of the candidate pre-Main Sequence stars (flared disks, remnants of the cocoon they were formed out of, etc.). This extra source of reddening is not accounted for by our method and the candidate pre-Main Sequence stars could be affected by a higher value of $E(B-V)$ than mean one from the 4 closest neighbors with a direct determination, \\ie have bluer intrinsic $(m_\\mathrm{F336W,0}-m_\\mathrm{F439W,0})$ colors. According to the data in Figure~\\ref{fig:ubuvex}, however, this would correspond to increasing the inferred Balmer continuum excess and, consequently, the accretion rate, further confirming the difference with what is observed in the Galaxy.  Finally, let us notice that the accretion rate for several of the Galactic stars in Figure~\\ref{fig:age-mass} was derived with the same technique we have used in this paper, \\ie the relation between Balmer continuum excess and accretion rate by \\citet{gull98}. As such, they are affected by the same issues on the calibration of this relation as the LMC stars presented here (if any). The discrepancy between Galactic stars and those in the neighborhood of SN~1987A in the LMC, then, appears to be real.  Of course, spectroscopic data on these stars would be highly desirable to better characterize their properties.  A higher accretion rate in the LMC than in the Galaxy was also suggested for A and B spectral type pre-Main Sequence stars by \\citet{lam99} and \\citet{dewit02}. The combination of this result with ours points towards a significantly higher accretion activity over a factor of about 6 in mass for pre-Main Sequence stars at lower metallicity."
    },
    "0402/astro-ph0402345_arXiv.txt": {
        "abstract": "We present seven epochs of spectroscopy on the quadruply imaged quasar SDSS~J1004+4112, spanning observed-frame time delays from 1 to 322 days.  The spectra reveal differences in the emission lines between the lensed images.  Specifically, component A showed a strong enhancement in the blue wings of several high-ionization lines relative to component B, which lasted at least 28 days (observed frame) then faded.  Since the predicted time delay between A and B is $\\lesssim$30 days, our time coverage suggests that the event was not intrinsic to the quasar.  We attribute these variations to microlensing of part of the broad emission line region of the quasar, apparently resolving structure in the source plane on a scale of $\\sim\\!10^{16}\\,{\\rm cm}$ at $z=1.734$.  In addition, we observed smaller differences in the emission line profiles between components A and B that persisted throughout the time span, which may also be due to microlensing or millilensing.  Further spectroscopic monitoring of this system holds considerable promise for resolving the structure of the broad emission line region in quasars. ",
        "introduction": "Microlensing in the images of a multiply-imaged quasar was first reported by \\markcite{iwh+89}{Irwin} {et~al.} (1989) for the quadruple lens Q2237+0305 \\markcite{huchra85}({Huchra} {et~al.} 1985).  Most quasar microlensing studies have been based on broad-band photometric monitoring \\markcite{wus+00,sus+03,cs03}(e.g., {Wo{\\' z}niak} {et~al.} 2000; {Schechter} {et~al.} 2003; {Colley} \\& {Schild} 2003), which is sensitive primarily to variations in the continuum. Microlensing of the continuum is expected since the optical/UV continuum emission is thought to originate in a region that is comparable in size to the Einstein radius of a typical star in a typical lens galaxy.  Microlensing of the broad emission line region (BELR) is also possible, if the BELR has structure on scales comparable to the Einstein radius of a star \\markcite{nem88,sw90}({Nemiroff} 1988; {Schneider} \\& {Wambsganss} 1990).  The possibility of BELR microlensing seemed rather remote until recent reverberation mapping work revised the estimate of the BELR size downward from $\\sim\\!10^{18}$ cm to $\\sim\\!10^{16}$ cm \\markcite{wpm99,ksn+00}({Wandel}, {Peterson}, \\& {Malkan} 1999; {Kaspi} {et~al.} 2000). Inspired by these numbers, \\markcite{amm+02}{Abajas} {et~al.} (2002) and \\markcite{la03}{Lewis} \\& {Ibata} (2003) revived the idea of looking for microlensing of the BELR and computed possible line profile variations for various BELR models.  Possible examples of microlensing of a quasar emission line have been presented by \\markcite{fil89}{Filippenko} (1989) for Q2237+0305, and by \\markcite{cea+04}{Chartas et al.} (2003) for H1413+117.  In particular, \\markcite{cea+04}{Chartas et al.} (2003) detected a strong, redshifted Fe K$\\alpha$ emission line in the X-ray spectrum of only one of the components in the quadruple lens H1413+117.  Although they did not have multiple epochs to look for variability, \\markcite{cea+04}{Chartas et al.} (2003) invoked a short predicted time delay between the components to argue that microlensing is the preferred explanation for seeing the Fe K$\\alpha$ line in only one component.  In this paper we present results from spectroscopic monitoring of the recently discovered quadruple lens SDSS~J1004+4112 \\markcite{inada03,oguri03}({Inada et al.} 2003; {Oguri et al.} 2003).  We observe variability in the broad emission line profiles of one of the lensed images that provides strong evidence for microlensing of the BELR, suggesting that the theoretical predictions for microlensing were correct and confirming that the BELR has structure on the scale of the Einstein radius of a star. ",
        "conclusions": "We have presented seven epochs of spectroscopic data on the two brightest components of the wide-separation, quadruply imaged quasar SDSS J1004+4112.  Although the simplest lensing scenarios predict that the four components should have identical spectra, the data reveal significant differences in the emission line profiles of the components.  In particular, the \\ion{C}{4} emission line profile in components A and B show both variable differences and differences that are constant over 322 observed-frame days.  The \\ion{He}{2} and \\ion{Si}{4}/\\ion{O}{4}] lines in component A also show variability similar to that seen in the \\ion{C}{4} line.  Because the predicted time delay between A and B is $\\lesssim$30 days, we argue that the differences are not due to intrinsic variability in the quasar coupled with a lensing time delay.  Instead, we suggest that the variability in the blue wing of component A is best explained by microlensing of part of the broad emission line region, resolving BELR structure on the order of a few light days.  This represents the first robust detection of BELR microlensing, with evidence based on multiple emission lines and involving observed variability.  The nature of the time-independent differences is less clear, but they may also be the result of a lensing event.  In any case, it is clear that continued spectroscopic monitoring of SDSS~J1004+4112 should be carried out in an attempt to map the structure of its broad emission line region through additional microlensing events."
    },
    "0402/astro-ph0402035_arXiv.txt": {
        "abstract": "{I review  X-ray observations of classical and recurrent novae in outburst, some of them recently done with {\\sl Chandra} and {\\sl XMM-Newton} for 12 objects. Significant X-ray flux is emitted by  the nova shell, with a peak luminosity up to L$_{\\rm x}$=10$^{35}$ erg s$^{-1}$ in the 0.2-10 keV range. In recurrent nova systems, or in novae hosting a red giant, the source of X-rays may be previous circumstellar matter shocked by the nova wind. However, for most classical novae, X-rays originate inside the nebula ejected in the outburst.  The data indicate a very high fraction of shocked material, and a non-smooth, varying wind outflow. A nebular emission line spectrum is also observed at late phases. In about half of the observed novae, the central white dwarf appears as a very luminous supersoft X-ray source for 1 to 9 years after the outburst. It is the best type of object to study the characteristics of shell hydrogen burning on white dwarfs in single degenerate systems. Still incomplete statistics indicate that  the duration of the supersoft X-ray phase is peaked around $\\simeq$2 years. The correlation of the X-ray light curve with the nova properties is not quite clear. Recently,  ``template grating spectra'' with high S/N have been obtained for  V4743 Sgr.  The X-ray light curve of this nova reveals a rich and complex power spectrum, with signatures of non-radial g-mode oscillations of the white dwarf.  The oscillations  and the spectra allow to determine the properties of the shell hydrogen burning white dwarf. }   \\addkeyword{Stars: binaries, white dwarfs} \\addkeyword{X-rays: stars}   \\begin{document} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0402/astro-ph0402203_arXiv.txt": {
        "abstract": "The VIMOS VLT Deep Survey (VVDS) is an on-going program to map the evolution of galaxies, large scale structures and AGNs from the redshift measurement of more than 100000 objects down to a magnitude $I_{AB}=24$, in combination with a multi-wavelength dataset from radio to X-rays. We present here the first results obtained from more than 20000 spectra. Dedicated effort has been invested to successfuly enter the ``redshift desert'' $1.5<z<3$, by producing high S/N galaxy templates in the rest wavelength domain probed by VIMOS. We have produced the luminosity function (LF) of galaxies for the first time out to $I_{AB}=24$, and the evolution of the LF for each morphological type is clearly established, indicating that most of the LF evolution is concentrated in the later types. We present for the first time the 3D spatial distribution of a well controlled sample of faint galaxies, and show that significant clustering is detected out to $z\\sim1.5$. ",
        "introduction": "The VIMOS VLT Deep Survey is dedicated to study galaxies, large scale structures and AGN evolution up to $z\\sim5$. To determine the evolutionary status of the galaxy and AGN population as a function of lookback time, it has been recognized early in the survey design phase that the large range in redshifts and the complex physical phenomena probed require to expand the core observations of the VVDS from the visible domain to a broad multi-wavelength spectrum. Observing the rest-frame SED of galaxies (e.g. in the B-band rest frame) at the same wavelength from $z\\sim0$ to $z\\sim5$, and having a sufficiently accurate estimate of the k-correction necessary to compute absolute magnitudes, requires to obtain deep photometric data from the U band to the K band. Furthermore, to get a complete picture of e.g. the star formation process, for galaxies with variable amount of dust extinction, or to study the evolution of large scale structures or AGNs, it is necessary to expand the wavelength baseline of observations from the far UV to the far IR, from the radio-waves to the X-rays.   We describe here the multi-wavelength approach followed by the VVDS, and present early results from the first epoch VVDS observations. We describe the effort to assemble for the first time an unbiased magnitude limited sample of galaxy redshifts and spectra $I_{AB}\\leq24$ up to $z\\sim5$, going through the ``redshift desert'' $1.5<z<2.5$,  and present preliminary results on the evolution of galaxies and large scale structures up to $z\\sim1.5$. ",
        "conclusions": "The VVDS is providing for the first time a magnitude limited sample of $\\sim15000$ galaxy redshifts, with $\\sim11000$ redshift in the $I_{AB}\\leq24$ sample, and $\\sim4000$ in the $I_{AB}\\leq22.5$ sample, going successfuly through the ``redshift desert'' $1.5<z<2.5$. The evolution of the luminosity function up to $z\\sim1.5$ is traced for several spectral type, with most of the evolution concentrated in the later types. Strong overdensities in  the distribution of galaxies are visible over scales in excess of 30Mpc up to $z\\sim1.5$. This provides evidence that the galaxy distribution was already strongly structured when the universe was half of its present age.\\\\  Acknowledgements: We thank C. Cesarsky and ESO for the early allocation of VLT Garanteed Observing Time, and A. Shapley who has provided fits files for her nice UV templates for $\\lambda\\leq1700\\AA$."
    },
    "0402/astro-ph0402217_arXiv.txt": {
        "abstract": "{ We perform three-dimensional numerical simulations of surface flows on the companion star in a semi-detached binary system and construct the corresponding Doppler maps. The mass ratio of the binary system, $q=M_2/M_1$, considered here is $q=0.05, 0.33, 0.5, 1, 2$, and $3$. For all cases, we obtain the H-, L1-, and L2-eddies, as found in our previous work, and confirm that the flow pattern does not heavily depend on the mass ratio.  We propose that this kind of problem be dubbed ``{\\it stellar meteorology}.'' The Doppler maps at the position of the companion show a structure tilted towards clockwise direction and presenting deviations from the critical Roche surface due to the L1-eddy and the L2-eddy on the companion star.  We apply our results to the Galactic supersoft X-ray source RX J0019.8+2156 and try to attribute the low radial velocity component of the emission lines of He II ${\\lambda}4686$ observed recently to the irradiated spot on the surface of the companion rather than that of the white dwarf or the accretion disc. Based on the comparison between the observations and our constructed Doppler map, we estimate the mass of the companion star in RX J0019.8+2156 to be $\\sim 2 M_{\\odot}$ assuming the mass of the white dwarf star to be around $0.6 M_{\\odot}$. ",
        "introduction": "A semi-detached binary system consists of a Roche-lobe-filling companion star and a (compact) accreting object. Part of the gas of the companion star is gravitationally attracted by the primary, flows through the L1 point towards the primary star and, if the latter is smaller than the circularisation radius, forms an accretion disc around it.  Many studies on the accretion disc itself exist, but until recently, only a very few studies considered the surface flow of the companion, except for the pioneering work by Lubow \\& Shu (\\cite{lub75}). A possible reason for the lack of studies on the companion is that there had been no effective way to observe the surface of the companion star as well as the flow on it.  However, Davey \\& Smith (\\cite{dav92}) demonstrated the surface mapping method of the companion surface in cataclysmic variables (CVs) using Na I line absorption. Dhillon \\& Watson (\\cite{dhi00}) demonstrated a technique of imaging companion stars of CVs using Roche tomography. Besides, Doppler map technique also can potentially give some trace of the surface flow (as considered in this paper).  Oka et al. (\\cite{oka02}) performed three-dimensional simulations of surface flows on the companion star and predicted the existence of three kinds of eddy associated with a high/low pressure on the companion star: the H-, L1-, and L2-eddies as schematically shown in Fig.1. The notations H, L1 and L2 denote the high pressure around a (north) pole, the low pressure around the L1 point and the low pressure at the opposite side to the L1 point, respectively.  In a rotating fluid, the pressure gradient force balances the Coriolis force, and the wind blows thus along isobaric lines. On the Earth, such a flow is called the geostrophic wind, while on a star it was called the astrostrophic wind by Lubow \\& Shu (\\cite{lub75}). Therefore,  a clockwise-rotating eddy forms around a high pressure region, while a counter-clockwise rotating eddy appears around a low pressure region on the northern hemisphere in our counterclockwise rotating system (viewed from the north). We dub this kind of problem ``\\textit{stellar meteorology}.''  With the observational progress mentioned above as well as theoretical ones in hand, we are now in the state of investigating the flow on the surface of the companion star in some more detail.  Oka et al. (\\cite{oka02}) considered only the mass ratio of unity. In reality, however, the range of mass ratio is wide, and it is thus interesting to investigate the dependence of the flow pattern on the mass ratio, $q=M_2/M_1$. In the present study, we consider six cases for the mass ratio: $q=0.05, 0.33, 0.5, 1, 2$, and $3$. As described above, a Doppler map can potentially give some traces of the surface flow, and so we construct Doppler maps in the present study as well.  \\begin{figure} \\centering \\resizebox{\\hsize}{!}{\\includegraphics{explain3_2.eps}} \\caption{A schematic diagram of the surface flow of a companion star in a semi-detached binary system: H, L1 and L2 denote eddies associated with a high/low pressure, respectively. Typical velocity vectors are shown with numbers.} \\label{fig1} \\end{figure}  \\subsection{Observational implication of the surface flow on the companion star}  Davey \\& Smith (\\cite{dav92}) found an asymmetric irradiation pattern in the secondary star of dwarf novae (e.g. IP Peg); the trailing hemisphere has stronger Na I line absorption than the leading hemisphere (see also Dhillon \\& Watson \\cite{dhi00}). They considered that this asymmetry could be due to the circulation current on the companion star (but see Watson et al. \\cite{wat03}).  As was stated above, Oka et al. (\\cite{oka02}) found three types of eddies on the companion star. They considered that the H-eddy as well as the L1-eddy could transport the heat due to the irradiation to the leading hemisphere, and argued that these currents naturally explained the asymmetry of the irradiation pattern. Note, however, that magnetic CVs, on the other hand, show the irradiation to be on the trailing hemisphere (Davey \\& Smith \\cite{dav92},\\cite{dav96}; Watson et al. \\cite{wat03}). To understand the irradiation pattern of magnetic CVs, we need another mechanism, which could lead to the effective heating of the trailing hemisphere.  \\subsection{Supersoft X-ray sources}  Supersoft X-ray sources (SSS) consist, like CVs, of a Roche lobe filling companion star and a white dwarf. X-rays are believed to be produced by a steady nuclear burning occurring on the surface of the white dwarf. This requires a high mass accretion rate (${\\ga} 10^{-7} M_{\\odot}/$yr) that is only possible with a high mass companion star (van den Heuvel et al. \\cite{van92}).  In our galaxy four SSSs are known (see e.g. G\\\"ansicke et al. \\cite{gan00}); one of which is RX J0019.8+2156. This object and its 12.2 mag optical counterpart were discovered by Beuermann et al. (\\cite{beu95}). Recent observations revealed a low radial velocity component to the He II ${\\lambda}$4686 emission line (e.g. Becker et al. \\cite{bec98}; Cowley et al. \\cite{cow98}; Deufel et al. \\cite{deu99}; McGrath et al. \\cite{mcg01}). These authors considered that the He II ${\\lambda4686}$ emission line was emitted from the white dwarf or from the inner part of the accretion disc. With this assumption they estimated the mass of the companion to be smaller than the white dwarf. However, this estimate certainly contradicts the massive companion theory developed by van den Heuvel et al. (\\cite{van92}).  Hachisu \\& Kato (2003), on the other hand, considered that the low radial velocity of the He II ${\\lambda }$4686 emission line was emitted from the inner hemisphere of the companion star. The inner hemisphere is irradiated by the companion and becomes hot enough to emit the He II ${\\lambda}$4686 line. The Doppler map of He II ${\\lambda 4686}$ observed by Deufel et al. (\\cite{deu99}) showed that the emission line peak is located around $V_{\\rm x}{\\sim }~100$ km/s and $V_{\\rm y}{\\sim}~100$ km/s (Note that the plus sign of the above velocity is due to our choice of the coordinate frame). This emission line peak does not correspond to the region of the white dwarf or the inner part of the accretion disc. With the irradiated companion model Hachisu \\& Kato (2003) argued that such low radial velocity emission lines were naturally explained by the massive companion star in a system with a mass ratio of $q{\\sim }3-4$ and a white dwarf of $0.6M_{\\odot}$.  Matsumoto \\& Mennickent (\\cite{matm00}) constructed the Doppler map of He II ${\\lambda4686}$ of RX J0925.5-4758 which is another known Galactic SSS and also found a peak located around $V_{\\rm x}\\sim 100$ km/s and $V_{\\rm y}\\sim 100$ km/s (in our velocity sign): i.e., the low radial velocity region. The low velocity of the He II ${\\lambda4686}$ emission line seems therefore to be a common feature of supersoft X-ray sources.  Hachisu \\& Kato (2003) tried to interpret the low radial velocity component of the He II ${\\lambda4686}$ emission line in the Doppler map of RX J0019.8+2156. However, their numerical model was two-dimensional, and could therefore not handle flows on the surface of the companion. A three-dimensional calculation is necessary in this respect, and this is the aim of the present research.  This paper is organized as follows. In Sect. 2 we describe the assumptions of the model and the numerical method. In Sect. 3 we show the flow patterns and construct Doppler maps. In Sect. 4 we discuss the problem of the supersoft X-ray sources. A discussion and a summary are given in Sect. 5.   \\begin{table} \\caption{Mass ratio, computational region, and number of grid points of models} \\label{Tab:Simu} \\begin{tabular}{cccc} \\hline Mass ratio &  & Computed region & Number of grid points \\\\ \\hline\\hline 0.05 &  & 1.0 $\\times$ 1.0 $\\times$ 0.5 & 101 $\\times$ 101 $\\times$ 51\\\\ 0.33 &  & 1.0 $\\times$ 1.0 $\\times$ 0.5 & 101 $\\times$ 101 $\\times$ 51\\\\ 0.5 &  & 1.1 $\\times$ 1.0 $\\times$ 0.5 & 111 $\\times$ 101 $\\times$ 51\\\\ 1 &  & 1.2 $\\times$ 1.0 $\\times$ 0.5 & 121 $\\times$ 101 $\\times$ 51\\\\ 2 &  & 1.4 $\\times$ 1.2 $\\times$ 0.6 & 141 $\\times$ 121 $\\times$ 61\\\\ 3 &  & 1.4 $\\times$ 1.2 $\\times$ 0.6 & 141 $\\times$ 121 $\\times$ 61\\\\ \\hline \\end{tabular} \\end{table} ",
        "conclusions": "We investigate the surface flow of the companion star in a semi-detached binary system and construct model Doppler maps. Six cases with mass ratios of $q=0.05, 0.33$, $0.5$, $1$, $2$, and $3$ are considered. We apply our results to the Galactic supersoft X-ray source RX J0019.8+2156. The results are summarized as follows:  \\begin{enumerate} \\item We obtain the H-, L1-, and L2-eddies, which were found by Oka et al. (\\cite{oka02}), for all mass ratios. We confirm that the flow pattern does not heavily depend on the mass ratio.  \\item We construct model Doppler maps, and we find an asymmetric velocity distribution, due to the surface flow.  \\item The companion mass of RX J0019.8+2156 is estimated to be around $2 M_{\\odot}$. This massive companion is consistent with the theoretical model proposed by van den Heuvel et al. (\\cite{van92}).  \\item We argue that the Doppler map of supersoft X-ray sources shows indications of the surface flow on the companion star. \\end{enumerate}"
    },
    "0402/astro-ph0402021_arXiv.txt": {
        "abstract": " ",
        "introduction": " ",
        "conclusions": "In our yet preliminary study we have looked at clouds with a particular ratio between the non-thermal and thermal pressure {\\it at the cloud center:} we investigated clouds with $(P_{\\rm w}/ P_{\\rm th})_{r=0}=0$ or 1, and with $(P_{\\rm m}/ P_{\\rm th})_{r=0}=0$ or 1, which is meant to give a first impression on the effect of non-thermal pressure support on the cloud stability. We adopt an extinction due to material surrounding the clouds of 5 magnitudes in the visual.  As a function of cloud mass we computed the maximum central column density for a stable hydrostatic equilibrium. These are plotted in Figure \\ref{fig:1} for clouds with and without non-thermal pressure. A comparison with the column densities measured for all starless low-mass cores we found in the literature shows that some of the more massive cores have column densities which are larger than what could be accounted for by a stable cloud with pure thermal pressure support.  With a magnetic or turbulent pressure at least comparable to the thermal pressure in the cloud center, higher central condensations comparable to those observed would be possible for stable clouds.  If this non-thermal pressure were due to turbulent motions, then the observed cores' non-thermal velocity dispersions should be comparable to the sound speed, but observations typically show smaller non-thermal linewidths \\cite{caselli2002}. Thus turbulence is unlikely to be the dominant support in such cores, and magnetic fields remain as the stabilizing force.  A second argument for the need of non-thermal support in the more massive low-mass cores derives from their maximum possible surface pressure. A core of given mass in a stable hydrostatic equilibrium can only exist if the kinetic surface pressure due to thermal and turbulent particle motions is below some critical value. Significant non-thermal pressure is necessary to stabilize the more massive cores ($M > {2 ~ \\rm to ~ 3} \\, M_\\odot$) in such environments.  To summarize, our preliminary analysis of the stability of dense, self-gravitating cores and our comparison with observed cores and their environments suggests the need for non-thermal pressure support, most likely provided by magnetic fields of less than $100 ~ \\rm \\umu G$."
    },
    "0402/astro-ph0402351_arXiv.txt": {
        "abstract": "We present a covariant framework of kinematics in the dS spacetime, which is a natural postulate of recent astronomical observations ($\\Lambda>0)$. One-particle states are presented explicitly. It is noticed that the dispersion relation of free particles is  dependent on the degrees of freedom of angular momentum and spin. This fact can be referred to as the effects of the cosmological constant on kinematics of particles. The kinematics in dS spacetime is used to investigate the phenomenon of ultra high energy cosmic rays.  We emphasize the possibility of solving the threshold anomalies of the interactions between ultra high energy cosmic rays and soft photons in the covariant framework of kinematics in the dS spacetime. \\newline \\newline PACS numbers: 98.70.Sa, ~95.30.Cq, ~95.85.-e. \\vspace{1.0cm} ",
        "introduction": "The origin of the ultra-high energy cosmic rays (UHECR) is one of the outstanding puzzles of modern astrophysics. Today's understanding of the phenomena responsible for the production of UHECR is still limited. Currently, there are generally two categories of production mechanisms of the UHECR. One is the ``bottom-up\" acceleration scenario with some astrophysical objects as sources\\cite{acc1,acc2}. The other is called ``top-down\" scenario in which UHECR particles are from the decay of certain sufficiently massive particles originating in the early Universe\\cite{top1}.  Decades ago, Greisen, Zatsepin and Kuzmin\\cite{GZK} discussed the propagation of the UHECR particles through the cosmic microwave background radiation (CMBR). Due to the photopion production process by the CMBR, the UHECR particles will lose their energies drastically down to a theoretical threshold, which is about $5\\times 10^{19}$eV. That is to say, the mean free path for this process is only a few Mpc\\cite{GZK1}. This is the so-called GZK cutoff. However, we have observed indeed hundreds of events with energies above $10^{19}$eV and about 20 events above $10^{20}$eV\\cite{data1}-\\cite{data5}. At the same time, there is another paradox\\cite{data6} in the terrain of cosmic ray, which comes from the detected $20$TeV photons from the MrK $501$ (a BL Lac object at a distance of $150$Mpc). Similar to the case of UHECR, due to pair production process by the IR background photons, the $20$TeV photons should have disappeared before arrival at the ground-based detectors. Both of the puzzles can be considered to be some cosmic ray threshold anomalies: energy of an expected threshold is reached but the threshold has not been observed yet.  Recently, there is growing interest on the theory of doubly special relativity (DSR)\\cite{DSR1, DSR2}. In the DSR, there exist two observer-independent scales. One of them is a scale of velocity, which is identified with the velocity of light. The other is a scale of mass/length, which is expected to be the order of the Planck mass/length. In fact, the violation of the Lorentz invariance and the Planck scale physics have long been studied as possible solutions of the cosmic ray threshold anomalies\\cite{loren2}-\\cite{loren8}. However, all of these scenarios are far beyond the standard cosmological theory and the standard model of particle physics. Furthermore, it is well-known that we are still very far from a theory of quantum gravity, in spite of extensive investigations on candidates such as the supergravity, Kluza-Klein, noncommutative geometry and superstring theory.  The recent astronomical observations on supernovae \\cite{constant-super-1, constant-super-2} and CMBR \\cite{costant-cmbr} show that about two thirds of the whole energy in the Universe is contributed by a small positive cosmological constant $(\\Lambda)$. An asymptotic de Sitter (dS) spacetime is premised naturally. The physics in an asymptotic dS spacetime has been discussed extensively\\cite{ds-paper-1}-\\cite{ds-paper-3}.  In this paper, we discuss kinematics in an asymptotic dS spacetime. The framework of classical as well as quantum kinematics in the dS spacetime is set up carefully. We get a general form of dispersion relation for free particles in the dS spacetime. This formalism is used to describe the UHECR propagating in the cosmic microwave background as well as the TeV-$\\gamma$ propagating in the infrared background. We obtain explicitly the corrections of the GZK threshold for the UHECR particles interacting with soft photons, which  are dependent on the cosmological constant as supposed in the beginning of the paper. We show how the threshold varies with a positive cosmological constant and additional degrees of freedom of the angular momentums of interacting particles. It should be noticed that, for a positive cosmological constant, the theoretic threshold tends to be above the energies of all the observed events. Thus, we may conclude that the tiny but nonzero cosmological constant is a possible origin of the threshold anomalies of the UHECR and the TeV-$\\gamma$.  The paper is organized as follows. In Section 2, we discuss the classical kinematics in the dS spacetime. Conservation laws of momentum and angular momentum are obtained along the geodesics. Section 3 is devoted to the investigation of the quantum kinematics in the dS spacetime. By solving equations of motion of a free particle, we present a remarkable dispersion relation for free particles in the dS spacetime, which includes degrees of freedom of angular momentum and spin. In Section 4, by taking effects of a tiny but nonzero positive cosmological constant into account, we show that the theoretic threshold is above the energies of all the observed UHECR events. Similar discussion is made for the TeV-$\\gamma$ in Section 5. In the end, we present conclusions and remarks. ",
        "conclusions": "In this paper, we discussed kinematics in the dS spacetime. The kinematic invariance group in the dS spacetime is $SO(1,4)$ instead of the Poincar\\'{e} one in the Minkowski sapcetime. Kinematics, which is based on the de Sitter group $SO(1,4)$, was set up formally.  It should be noticed that the dispersion relation for a free particle in the dS spacetime is related with degrees of freedom of angular momentum and spin. With the help of this deformed dispersion relation, and thanks to the positive cosmological constant, a possible origin of the threshold anomalies was proposed naturally.  We would like to point out that, if the cosmological constant can vary notably in different period of the Universe (at least in some scenarios it is), we may get a higher threshold for the UHECR propagating in the early Universe. And this could be a good news for the ``top-down\" scenario of the UHECR. \\newline \\newline {\\large\\bf Acknowledgement:}\\\\ We would like to thank Prof. H. Y. Guo and C. J. Zhu for useful discussion. The work was supported partly by the Natural Science Foundation of China. One of us (C.B.G.) is supported by grants through the ICTS (USTC) from the Chinese Academy of Sciences."
    },
    "0402/astro-ph0402398_arXiv.txt": {
        "abstract": "{We present high precision CCD photometry of 1791 objects in 20 open clusters with an age from 10\\,Myr to 1\\,Gyr. These observations were performed within the $\\Delta a$ photometric system which is primarily used to detect chemically peculiar stars of the upper main sequence. Time bases range between 30 minutes and up to 60 days with data from several nights. We describe the way of time series analysis reaching a detection limit of down to 0.006\\,mag. In total, we have detected 35 variable objects from which four are not members of their corresponding clusters. The variables cover the entire Hertzsprung-Russell-diagram, hence they are interesting targets for follow-up observations. ",
        "introduction": "The detection of variable members in open clusters is very important since these objects have fairly well known astrophysical parameters, such as luminosity and effective temperature. Several theories (e.g. pulsational and evolutionary models) can be tested with these variable stars.  In the literature, a huge amount of papers dedicated to the search for new variable stars in open clusters can be found. In general, two different kinds of surveys are conducted: 1) the search for special types of variables (Viskum et al. 1997, Jerzykiewicz et al. 2003) or 2) selected open clusters are searched for all kinds of variable objects (Kafka \\& Honeycutt 2003, Mochejska et al. 2003).  Our search for new variable stars in open clusters is a serendipity result from already published CCD $\\Delta a$ photometry (Bayer et al. 2000, Paunzen \\& Maitzen 2001, 2002 and Paunzen et al. 2002, 2003). The intermediate band, three filter $\\Delta a$ system investigates the flux depression at 5200\\,\\AA\\, found for magnetic chemically peculiar objects (Maitzen 1976). Our observations span widely different time intervals (0.02 to 60 days) yielding different possibilities for detecting the whole set of variations. We want to emphasize that these observations are not optimized for the detection of variable stars but are capable to find even very low amplitude variables (the typical detection limit reached is between 0.006 and 0.022 mag).  We describe the way to define the variability limit and present all bona-fide variable stars within the Hertzsprung-Russell-diagram. Four objects are probably not members of the corresponding open clusters. We give a discussion about the possible nature of the detected variability.  \\begin{table*} \\caption[]{Open clusters observed at ESO and UTSO in 1995 (upper panel) as well as CASLEO in 1998 and 2001 (lower panel). The ages (log\\,$t$) and distance moduli ($V_0-M_V$) were taken from the literature. The limit of apparent variability (Limit) is according to Sect. \\ref{time}. The errors in the final digits of the corresponding quantity are given in parentheses.} \\label{clusters} \\begin{center} \\begin{tabular}{lcrccccrccr} \\hline \\hline \\multicolumn{2}{c}{Designation} & N$_{S}$ & N$_{V}$ & N$_{F}$ & N$_{N}$ & JD (start) & \\multicolumn{1}{c}{$\\Delta t$} & Limit & log\\,$t$ & $V_0-M_V$ \\\\ & & & & & & & \\multicolumn{1}{c}{[d]} & [mag] \\\\ \\hline NGC 2439 & C0738$-$315 & 115 & 3 & 18 & 4 & 2449816.51736 & 7.995 & 0.022 & 7.30 & 13.00(10) \\\\ NGC 2489 & C0754$-$299 & 53 & 1 & 13 & 4 & 2449818.57083 & 34.937 & 0.012 & 8.45 & 10.80(10)  \\\\ NGC 2567 & C0816$-$304 & 34 & $-$ & 17 & 4 & 2449818.61111 & 44.868 & 0.012 & 8.43 & 11.10(10) \\\\ NGC 2658 & C0841$-$324 & 84 & 1 & 12 & 3 & 2449817.56597 & 18.026 & 0.018 & 8.50 & 12.90(15) \\\\ Melotte 105 & C1117$-$632 & 122 & $-$ & 15 & 2 & 2449816.61111 & 3.072 & 0.010 & 7.77 & 11.30(10) % \\\\ NGC 3960 & C1148$-$554 & 32 & $-$ & 17 & 3 & 2449828.60625 & 59.865 & 0.014 & 8.88 & 11.10(20) \\\\ NGC 5281 & C1343$-$626 & 16 & 1 & 55 & 6 & 2449816.78125 & 48.010 & 0.008 & 7.04 & 10.60(15) \\\\ NGC 6134 & C1624$-$490 & 82 & 2 & 32 & 6 & 2449821.75069 & 18.003 & 0.010 & 8.84 & 9.80(15) \\\\ NGC 6192 & C1636$-$432 & 64 & $-$ & 38 & 6 & 2449822.79583 & 44.949 & 0.006 & 7.95 & 11.15(20)  \\\\ NGC 6208 & C1645$-$537 & 15 & $-$ & 12 & 2 & 2449880.69236 & 2.010 & 0.020 & 9.00 & 10.00(15) \\\\ NGC 6396 & C1734$-$349 & 48 & 2 & 18 & 5 & 2449821.88958 & 31.939 & 0.014 & 7.40 & 10.60(15)  \\\\ NGC 6451 & C1747$-$302 & 41 & 2 & 12 & 2 & 2449883.78125 & 0.988 & 0.010 & 8.30 & 11.65(20) \\\\ NGC 6611 & C1816$-$138 & 45 & 2 & 42 & 5 & 2449849.79306 & 38.989 & 0.016 & 6.48 & 11.65(10) \\\\ NGC 6705 & C1848$-$063 & 275 & 1 & 43 & 5 & 2449822.86250 & 59.911 & 0.014 & 8.40 & 11.65(20) \\\\ NGC 6756 & C1906+046 & 33 & 3 & 38 & 5 & 2449823.90069 & 53.944 & 0.008 & 8.11 & 12.60(15) \\\\ \\hline NGC 3114 & C1001$-$598 & 181 & 7 & 50 & 4 & 2451138.82296 & 2.979 & 0.022 & 8.48 & 9.60(15)  \\\\ Collinder 272 & C1327$-$610 & 45 & 2 & 22 & 1 & 2452144.48472 & 0.020 & 0.008 & 7.11 & 11.85(15)  % \\\\ Pismis 20 & C1511$-$588 & 178 & 2 & 80 & 2 & 2452143.52872 & 1.045 & 0.022 & 6.70 & 12.55(20)  \\\\ NGC 6204 & C1642$-$469 & 268 & 3 & 55 & 1 & 2452143.63456 & 0.067 & 0.020 & 8.30 & 10.40(25)  \\\\ Lyng\\aa\\, 14 & C1651$-$452 & 60 & 3 & 70 & 1 & 2452144.58135 & 0.087 & 0.008 & 6.00 & 12.05(15)  % \\\\ \\hline \\multicolumn{11}{l}{N$_{S}$ $\\dotfill$ number of investigated stars; N$_{V}$ $\\dotfill$ number of variable objects; N$_{F}$ $\\dotfill$ number of frames;} \\\\ \\multicolumn{7}{l}{N$_{N}$ $\\dotfill$ number of nights; $\\Delta t$ $\\dotfill$ time base of the observations} \\\\ \\end{tabular} \\end{center} \\end{table*} ",
        "conclusions": ""
    },
    "0402/astro-ph0402167_arXiv.txt": {
        "abstract": "We present first results of our NAOS-CONICA search for close substellar companions around young nearby stars. This program was started only a few months ago. We have obtained 1$^{st}$ epoch images of several targets which are unpublished young stars ($<$100\\,Myrs), hence ideal targets to look for planetary companions. For one target star we could even take a 2$^{nd}$ epoch image. By comparing both images we could look for co-moving companions of the target star. Those data show clearly that the detection of planetary companions (m$<$13\\,$M_{Jup}$) inward a saturn-like orbit (r$<$10\\,AU) is feasible with NAOS-CONICA and in addition that the astrometric confirmation of those companions is doable with only a few weeks of epoch difference. ",
        "introduction": "We have taken deep, high dynamic range images of most stars in the TW Hydra, Horologium, Tucana and $\\beta$ Pic groups (age=10..40\\,Gyrs, d=10..60\\,pc). Such young nearby stars are well suited for direct imaging of substellar companions, both brown dwarfs and massive planets (Neuh\\\"auser et al., 2003[1]). The target stars are located in young star forming regions, hence companions are young and therefore still self-luminous. Due to the proximity of the target stars a high spatial resolution can be achieved. We started our search by using seeing limited imaging with SOFI and speckle technique with the MPE speckle camera SHARP, both at the ESO 3.5m NTT on La Silla. Companion candidates detected in 1$^{st}$ epoch images must be confirmed by 2$^{nd}$ epoch images (proper motion) and follow-up spectra, in the latter case done with ISAAC at the VLT. From non-detection and detection limits, we could conclude that massive planets (m$>$5\\,$M_{Jup}$) in wide orbits (a$>$50\\,AU) are rare and appear around less than 9\\,\\% of the stars. Finally our results have shown that the detection of extrasolar planets around young nearby stars was indeed already possible without AO, i.e. with SHARP at the 3.5m\\,NTT, but only at large separation outside 50\\,AU. ",
        "conclusions": ""
    },
    "0402/astro-ph0402484_arXiv.txt": {
        "abstract": "We use high resolution collisionless $N$-body simulations to study the secular evolution of disk galaxies and in particular the final properties of disks that suffer a bar and perhaps a bar-buckling instability.  Although we find that bars are not destroyed by the buckling instability, when we decompose the radial density profiles of the secularly-evolved disks into inner S\\'ersic and outer exponential components, for favorable viewing angles, the resulting structural parameters, scaling relations and global kinematics of the bar components are in good agreement with those obtained for bulges of late-type galaxies. Round bulges may require a different formation channel or dissipational processes. ",
        "introduction": "Evidence has accumulated in the past decade showing that many bulges, especially at low-masses, have a disk-like, almost-exponential radial fall-off of the stellar density \\citep{as94,cdjb96,dj96,csm98,csdzsd01,c99,mch03}, and in some cases disk-like, cold kinematics \\citep{k93,kbb02}.  Comparisons of bulge and disk parameters have furthermore shown a correlation between the scale-lengths of bulges and disks \\citep{dj96,mch03} and, on average, similar colors in bulges and inner disks \\citep{tdfdw94,pb96,cdjb96}.  The disk-like properties of bulges and the links between bulge and disk properties have been suggested to indicate that bulges may form through the evolution of disk dynamical instabilities such as bars, which are present in about $70\\%$ of nearby disk galaxies \\citep{k99,efp00}.  It has long been known that bars lead to angular momentum redistribution and to an associated increase in the central mass density.  Already \\citet{h71} found that an initially single component disk evolved a double exponential density profile under the influence of a bar; as a result, the scale-length of the outer disk increases. It has been suggested that bars may be efficient at building {\\it three-dimensional} stellar bulge-like structures via scattering of stars at vertical resonances \\citep{cdfp90}, or by the collisionless buckling instability (which weakens the bar; Raha et al. 1991), or via bar destruction due to the growth of a central mass concentration \\citep{pn90,nsh96}.  As buckling occurs relatively easily, it could be a particularly promising avenue for bulge formation.  This instability leads to the large-scale, coherent bending of the bar perpendicular to the plane of the disk \\citep{rsjk91}.  Buckling, which is a result of vertical anisotropy \\citep{a85,fp84,mh91,ms94}, thickens the stellar system and weakens the bar.  Evidence that buckling occurs in nature has relied on the fact that buckled bars are boxy/peanut-shaped when viewed edge-on and on the unmistakable gas-kinematic signature of a bar observed in such bulges \\citep{km95,mk99,bf99}.  The relevance of buckling to bulge formation remains however rather anectdotal, as little work has been done to check that the structural and kinematic properties of buckled bars are quantitatively consistent with those observed in bulges.  In this Letter, we report on high resolution $N$-body simulations of bar-unstable disks, some of which buckle, and compare their secularly-evolved structural and kinematic properties with those of bulges in local galaxies. We describe our simulations in \\S 2, present our results in \\S 3 and discuss their implications in \\S 4. ",
        "conclusions": "Although \\citet{rsjk91} only reported bar-weakening, not destruction, by the buckling instability, it has become common wisdom that buckling destroys bars.  Our simulations have failed to turn up a single instance in which the bar was destroyed by buckling, although we cannot exclude that it is in some extreme case.  The channel of bulge formation by the dissipationless destruction of bars during buckling, therefore, is not viable.  However, the minimal collisionless secular evolution present in our simulations must also occur in nature, resulting in systems that exhibit double component mass density profiles.  For certain viewing orientations, the spread in structural parameters and kinematic properties are indistinguishable from those observed in systems that are classified as bulges.  Nonetheless, the existence in nature of round bulges inside low inclination galaxies, which our simulations cannot reproduce, requires that other processes are also involved.  Possibly secular evolution including dissipative gas \\citep{mw04} will result in rounder bulges, but this requires extended central objects with masses of $\\sim 10-20\\%$ that of the disk \\citep{ss04}.  The higher interaction rate in the early universe may have triggered the large gas inflows needed to build such objects.  Finally, because the halos of our simulations are rigid, the baryonic component cannot but conserve its angular momentum.  In (semi-) analytic models of disk galaxy formation \\citep{fe80,mmw98,ls91,wf91,dss97,vdb98}, the distribution of disk scale-lengths, $R_d$, is set by that of their angular momenta. \\citet{djl96} found that the width of the observed $R_d$ distribution at fixed luminosity is smaller than that predicted by analytic theory. Our simulations show that, under the influence of a bar, $R_d$ may increase by a factor of 2 or more at constant global angular momentum. Since less extended disks are likely to be more bar-unstable, the secular evolution of these disks may be responsible for at least part of this discrepancy."
    },
    "0402/hep-ph0402102_arXiv.txt": {
        "abstract": "s{ After a short review of the ultrahigh energy cosmic ray puzzle  -- the apparent observation of cosmic rays originating from cosmological distances with energies above the expected Greisen-Zatsepin-Kuzmin cutoff $4\\times 10^{19}$~eV -- we consider strongly interacting neutrino scenarios as an especially interesting solution. We show that all features of the ultrahigh energy cosmic ray spectrum from $10^{17}$~eV to $10^{21}$~eV can be described to originate from a simple power-like injection spectrum of protons, under the assumption that the neutrino-nucleon cross-section is significantly enhanced at center of mass energies above $\\approx 100$~TeV. In such a scenario, the cosmogenic neutrinos produced during the propagation of protons through the cosmic microwave background initiate air showers in the atmosphere, just as the protons. The total air shower spectrum induced by protons and neutrinos shows excellent agreement with the observations. We shortly discuss TeV-scale extensions of the Standard Model which may lead to a realization of a strongly interacting neutrino scenario. We emphasize, however, that such a scenario may even be realized within the standard electroweak model: electroweak instanton/sphaleron induced processes may get strong at ultrahigh energies. Possible tests of strongly interacting neutrino scenarios range from observations at cosmic ray facilities and neutrino telescopes to searches at lepton nucleon scattering experiments. } ",
        "introduction": "The Earth's atmosphere is continuously bombarded by cosmic particles (``rays''). Their measured flux extends over many orders of magnitude in energy (cf. Fig.~\\ref{cr-spectrum}). At energies above $10^{15}$~eV, they are observed in the form of extensive air showers (EAS's), initiated by inelastic scattering processes of cosmic particles off atmospheric nucleons. Ground-based observatories have measured EAS's with energies up to $E\\,\\lwig\\, 3\\times 10^{20}$~eV, corresponding to center-of-mass (CM) energies $\\sqrt{s}=\\sqrt{2 m_p E}\\,\\lwig\\, 750$~TeV, where $m_p$ is the proton mass. Therefore, the highest energy cosmic rays probe physics beyond the reach of the (Very\\cite{vlhc}) Large Hadron Collider\\cite{lhc} ((V)LHC), with a projected CM energy of $14\\,(200)$~TeV. In this context, it is interesting that the measured cosmic ray flux at the highest energies, $E\\gwig 10^{20}$~eV, represents a puzzle. What is this puzzle about?  \\begin{figure}% \\centerline{\\epsfxsize=3.2in\\epsfbox{cr_rev_f1.eps}} \\caption[...]{Compilation of measurements of the flux of cosmic rays. The dotted line shows an $E^{-3}$ power-law for comparison. Approximate integral fluxes (per steradian) are also shown (adapted\\cite{Olinto:2000fz} from Ref.\\cite{Cronin:vy}). \\label{cr-spectrum}} \\end{figure}    It hinges on the circumstantial evidence that the cosmic rays above $10^{17.5\\div 18.5}$~eV originate from cosmological distances (for a recent review, see Ref.\\cite{Anchordoqui:2002hs}). This evidence is largely based on the apparent large-scale isotropy in the arrival directions of cosmic rays (cf. Fig.~\\ref{arr-dist}). Moreover, whereas there are only very few -- if any -- nearby source candidates, plausible astrophysical sources are most likely to be found only at cosmological distances.  \\begin{figure}% \\vspace{-0.2cm} \\centerline{\\epsfxsize=3.6in\\epsfbox{arr_dir_agasa_21_01_04.eps}} \\caption[...]{ Arrival directions of cosmic rays detected by the AGASA and Akeno (A20) experiments in equatorial coordinates. Open circles and open squares represent cosmic rays with energies $(4 \\div 10) \\times 10^{19}$~eV, and $\\geq 10^{20}$~eV, respectively. The galactic and super-galactic planes are shown by the red and blue curves, respectively. Large shaded circles indicate event clusters within $2.5^{\\circ}$. The shaded regions indicate the celestial regions excluded by a zenith angle cut of $\\leq 45^\\circ$. Update\\cite{AGASA} (June 24, 2003) of the published data from Ref.\\cite{Takeda:1999sg}. \\label{arr-dist}} \\end{figure}  If the highest energy cosmic rays are nucleons (or nuclei), if their sources are indeed uniformly distributed at cosmological distances, and if their injection spectra are power-laws in energy -- a reasonable assumption, in view of the measured spectrum in Fig.~\\ref{cr-spectrum} which appears to be approximately of (broken) power-law type over many order of magnitude in energy -- then their total flux arriving at Earth should show a pronounced drop above the Greisen-Zatsepin-Kuzmin\\cite{Greisen:1966jv} (GZK) ``cutoff'' $E_{\\rm GZK}=4\\times 10^{19}$~eV. This is due to the fact that, above this energy, the universe becomes opaque to high energy nucleons (and nuclei), due to inelastic hadronic scattering processes with the cosmic microwave background (CMB) photons. The GZK cutoff is, however, not seen in the data, at least not in a significant manner (cf. Fig.~\\ref{uhecr-data}). Correspondingly, the events above $10^{20}$~eV in Fig.~\\ref{uhecr-data} should originate from small distances below $50$~Mpc, the typical interaction length of nucleons above $E_{\\rm GZK}$. However, no source within a distance of $50$~Mpc is known in the arrival directions of the post-GZK events\\footnote{The dominant radio galaxy M87 in the Virgo cluster, at a distance of about $20$~Mpc, has been a source candidate for a long time\\cite{Ginzburg:63}. The major difficulty with this idea is the isotropy of the arrival distribution. It might be overcome by invoking a particular galactic magnetic field originating from a ``galactic wind''\\cite{Ahn:1999jd}. Criticisms of this model\\cite{Billoir:2000wi} have been addressed in Ref.\\cite{Biermann}.}. The basic puzzle is: if the sources of ultrahigh energy cosmic rays are indeed at cosmological distances, how could they reach us with energies above $10^{20}$~eV?   \\begin{figure}% \\vspace{-0.9cm} \\centerline{\\epsfxsize=3.8in\\epsfbox{e3flux_pow.eps}} \\caption[...]{ Ultrahigh energy cosmic ray data with their statistical errors (top: combination of Akeno\\cite{Nagano:1991jz} and AGASA\\cite{Takeda:1998ps} data; bottom: combination of Fly's Eye\\cite{Bird:yi} and HiRes\\cite{Abu-Zayyad:2002ta} data) and the predictions arising from a power-law emissivity distribution~(\\ref{source-emissivity}) corresponding to sources which are uniformly distributed at cosmological distances. The best fits between $E_-=10^{17.2}$~eV and $E_+=10^{20}$~eV are given by the solid lines and correspond to the indicated values of the parameters $\\alpha$ and $n$ in the source emissivity distribution. The 2-sigma variations corresponding to the minimal (dotted) and maximal (dashed) fluxes are also shown. Other parameters of the analysis were $E_{\\rm max} = 3 \\times 10^{21}$~eV, $z_{\\rm min} = 0.012$, and $z_{\\rm max}=2$. From Ref.\\cite{Fodor:2003ph}. \\label{uhecr-data}} \\end{figure}  At the relevant energies, among the known particles only neutrinos can propagate without significant energy loss from cosmological distances to us. It is this fact which led, on the one hand, to scenarios invoking hypothetical -- beyond the Standard Model -- strong interactions of ultrahigh energy cosmic neutrinos\\cite{Beresinsky:qj} and, on the other hand, to the Z-burst scenario\\cite{Weiler:1997sh,Fodor:2001qy}.  In the latter, ultrahigh energy cosmic neutrinos (UHEC$\\nu$'s) produce Z-bosons through annihilation with the relic neutrino background from the big bang. On Earth, we observe the air showers initiated by the protons and photons from the hadronic decays of these Z-bosons. Though the required ultrahigh energy cosmic neutrino flux\\cite{Fodor:2001qy} is smaller than present upper bounds\\cite{Kravchenko:2003tc}, it is not easy to conceive a production mechanism yielding a sufficiently large one. In the near future, UHEC$\\nu$ detectors, such as the Pierre Auger Observatory\\cite{Auger}, IceCube\\cite{IceCube}, ANITA\\cite{ANITA}, EUSO\\cite{EUSO}, OWL\\cite{OWL}, and SalSA\\cite{Gorham:2001wr} can directly confirm or exclude this scenario\\cite{Eberle:2004ua}.  Scenarios based on strongly interacting neutrinos, on the other hand, are based on the observation that the flux of neutrinos originating from the decay of the pions produced during the propagation of nucleons through the CMB\\cite{Beresinsky:qj,Stecker:1979ah,Yoshida:pt,Protheroe:1995ft} -- the cosmogenic neutrinos -- shows a nice agreement with the observed ultrahigh energy cosmic ray (UHECR) flux above $E_{{\\rm GZK}}$. Assuming a large enough neutrino-nucleon cross-section at these high energies, these neutrinos could initiate extensive air showers high up in the atmosphere, like hadrons, and explain the existence of the post-GZK events\\cite{Beresinsky:qj}. This large cross-section is usually ensured by new types of TeV-scale interactions beyond the Standard Model, such as arising through gluonic bound state leptons\\cite{Bordes:1997bt}, through TeV-scale grand unification with leptoquarks\\cite{Domokos:2000dp}, through Kaluza-Klein modes from compactified extra dimensions\\cite{Domokos:1998ry} (see, however, Ref.\\cite{Kachelriess:2000cb}), or through $p$-brane production in models with warped extra dimensions\\cite{Ahn:2002mj} (see, however, Ref.\\cite{Anchordoqui:2002it}); for earlier and further proposals, see Ref.\\cite{Domokos:1986qy} and Ref.\\cite{Barshay:2001eq}, respectively.  In this review, we discuss strongly interacting neutrino scenarios as a possible solution to the GZK puzzle. We present  a detailed statistical analysis of the agreement between observations and predictions from such scenarios\\cite{Fodor:2003bn}. Moreover, we emphasize an example which -- in contrast to previous proposals -- is based entirely on the Standard Model of particle physics. It exploits non-perturbative electroweak instanton-induced processes for the interaction of cosmogenic neutrinos with nucleons in the atmosphere, which may have a sizeable cross-section above a threshold energy $E_{\\rm th}={\\mathcal O}( (4\\pi m_W/\\alpha_W )^2)/(2 m_p) = {\\mathcal O}( 10^{18})$~eV, where $m_W$ denotes the W-boson mass and $\\alpha_W$ the electroweak fine structure constant\\cite{Aoyama:1986ej,Morris:1991bb,Ringwald:2002sw}.   Our scenario is based on a standard power-like primary spectrum of protons injected from sources at cosmological distances. After propagation through the CMB, these protons will have energies below $E_{{\\rm GZK}}$, so they can well describe the low energy part of the UHECR spectrum. The cosmogenic neutrinos interact with the atmosphere and thus give a second component to the UHECR flux, which describes the high energy part of the spectrum. The relative normalization of the proton and neutrino fluxes is fixed in this scenario, so the low and high energy parts of the spectrum are explained simultaneously without any extra normalization. Details of this analysis can be found in Ref.\\cite{Fodor:2003bn}.  The structure of this review is as follows. In the next section, we review our procedure to infer the fluxes of protons and cosmogenic neutrinos at Earth, from an assumed injection spectrum at the sources. In Sect.~\\ref{inst-spect}, various possibilities, including the one exploiting electroweak instantons, for a large neutrino-nucleon cross-section at high energies are discussed, and the induced air shower rate is calculated. In Sect.~\\ref{comparison}, we present a comparison of the predictions with the observations and a determination of the goodness of the fit. Possible further tests are mentioned in Sect.~\\ref{tests},  while conclusions are given in Sect.~\\ref{conclusions}. ",
        "conclusions": "Summary and conclusions}   We have shown that a simple scenario with a single power-law injection spectrum of protons can describe all the features of the UHECR spectrum in the energy range $10^{17\\div 21}$~eV, provided the neutrino-nucleon cross-sections becomes of hadronic size at energies above $\\approx 10^{19}$~eV. In such a strongly interacting neutrino scenario, the cosmogenic neutrinos, which have been produced during proton propagation through the CMB, initiate air showers high up in the atmosphere and give thus rise to a second, neutrino-induced EAS component, extending well beyond the GZK energy. As examples giving rise to the necessary enhancement in $\\sigma_{\\nu N}$, we discussed $p$-brane production in TeV-scale gravity scenarios and Standard Model electroweak instanton-induced processes. The model for the proton injection spectrum has few parameters from which only two -- the power index $\\alpha$ and the redshift evolution index $n$ -- has a strong effect on the final shape of the spectrum. We found that, for certain values of $\\alpha$ and $n$, strongly interacting neutrino scenarios are compatible with the available observational data from the AGASA and HiRes experiments (combined with their predecessor  experiments, Fly's Eye and Akeno, respectively) on the 2-sigma level (also 1-sigma for HiRes). There are a number of astrophysical source candidates, notably  neutron stars and GRB's, which may provide the necessary conditions to accelerate protons to the required energies, $E_{\\rm max}\\gwig\\, 3\\times 10^{21}$~eV, by conventional shock acceleration.  The predicted ultrahigh energy cosmic neutrino component can be experimentally tested by studying the zenith angle dependence of the events in the range $10^{18\\div 20}$~eV and possible correlations with distant astrophysical sources at cosmic ray facilities such as the Pierre Auger Observatory and EUSO, and by looking for bumps in neutrino-initiated shower spectra at neutrino telescopes such as ANTARES and IceCube. As laboratory tests, one may search for a enhancements in (quasi-)elastic lepton-nucleon scattering or for signatures of QCD instanton-induced processes in deep-inelastic scattering, e.g. at HERA.  In summary, strongly interacting neutrino scenarios provide a viable and attractive solution to the ultrahigh energy cosmic ray puzzle and may be subject to various crucial tests in the foreseeable future."
    },
    "0402/astro-ph0402601_arXiv.txt": {
        "abstract": "The close-in extrasolar giant planets (CEGPs) reside in irradiated environments much more intense than that of the giant planets in our solar system. The high UV irradiance strongly influences their photochemistry and the general current view believed that this high UV flux will greatly enhance photochemical production of hydrocarbon aerosols. In this letter, we investigate hydrocarbon aerosol formation in the atmospheres of CEGPs. We find that the abundances of hydrocarbons in the atmospheres of CEGPs are significantly less than that of Jupiter except for models in which the CH$_4$ abundance is unreasonably high (as high as CO) for the hot (effective temperatures $\\gtrsim 1000$~K) atmospheres. Moreover, the hydrocarbons will be condensed out to form aerosols only when the temperature-pressure profiles of the species intersect with the saturation profiles---a case almost certainly not realized in the hot CEGPs atmospheres. Hence our models show that photochemical hydrocarbon aerosols are insignificant in the atmospheres of CEGPs. In contrast, Jupiter and Saturn have a much higher abundance of hydrocarbon aerosols in their atmospheres which are responsible for strong absorption shortward of 600~nm. Thus the insignificance of photochemical hydrocarbon aerosols in the atmospheres of CEGPs rules out one class of models with low albedos and featureless spectra shortward of 600~nm. ",
        "introduction": "Hazes and clouds\\footnote{``Hazes\" refers to the diffuse and optically thin aerosol distribution, while ``clouds\" refers to the optically thick regions \\citep{Wetal86}.} in the atmospheres of jovian planets can strongly affect the ability to determine atmospheric composition at ultraviolet to infrared wavelengths. At wavelengths shorter than $\\sim$600~nm, the atmospheric line features in the jovian planets are ``washed out\" by the hazes/clouds in the atmospheres of planets (e.g., Karkoschka \\& Tomasko 1993; Karkoschka 1998). The main chemical compositions of the hazes/clouds on Jupiter are believed to be H$_2$O-NH$_3$, NH$_4$SH, NH$_3$, N$_2$H$_4$, and hydrocarbons from several bar to $\\sim$0.1~mbar \\citep{WL73,S83,Wetal86,PH91,Getal96,Betal98a,Betal98b,Wetal03}. Saturn may have a composition profile similar to Jupiter since they have similar 300-1000~nm spectra (e.g., Karkoschka 1998). Saturn's albedo has been successfully modelled by assuming a dichotomy in the aerosol distribution between the troposphere and stratosphere, where the number density of aerosols is much lower in the stratosphere \\citep{KT93}. It is found that the stratospheric aerosols are very dark at $\\sim$300~nm, implying the presence of hydrocarbon aerosols.  Since the recent increase in sample size of extrasolar planets (e.g., Udry et al. 2002; Butler et al. 2003), the planetary formation environment has been statistically analyzed, although not conclusively \\citep{Fetal02,SIMRU03}. The close-in extrasolar giant planets (CEGPs, with semi-major axes $\\lesssim 0.05$ AU, also known as ``hot Jupiters\") are of particular interest since they have more active chemical processes in their atmospheres (e.g., Liang et al. 2003) and the evolution of the atmospheres can currently be studied observationally (e.g., Vidal-Madjar et al. 2003, 2004). A number of simulations in the atmospheres of CEGPs have been performed to study the albedos and reflection spectra by including the formation of high temperature condensates, such as silicates (e.g., Sudarsky et al. 2000; Seager et al. 2000).  The importance and existence of the atmospheric aerosols have been addressed and discussed widely in recent years (e.g., Baraffe et al. 2003) and it is generally believed that more UV flux will result in more aerosols. The photochemistry in jovian atmospheres results in photochemical aerosols which significantly affect the ultraviolet-visible spectra and albedos; hence we were motivated to simulate the formation of various molecules, e.g., hydrocarbons, ammonia, and sulfuric acid, which are the possible sources of aerosols, in the atmospheres of CEGPs. In this letter, we focus on hydrocarbons and hydrocarbon aerosol formation. ",
        "conclusions": "} Using a simplified version of the Caltech/JPL KINETICS model, we have shown that the concentrations of the C$_2$H$_{2n}$ species (see Table~\\ref{column}) are insignificant in the atmospheres of CEGPs. These C$_2$H$_{2n}$ compounds are important sources for forming more complex C$_x$H$_y$ species, such as benzene and PAHs, which will lead to the formation of hydrocarbon aerosols (e.g., Richter \\& Howard 2000, 2002). Although we have used a simplified photochemical model that captures the main reactions, we have tested Models A-E using the reference temperature profile (solid-line in Figure~\\ref{profile}) incorporating the full version of hydrocarbon model by \\citet{Getal96}. Even for this case, we find that the C$_6$H$_6$ abundance for Model A is seven orders of magnitudes less than that of Jupiter and is two orders of magnitudes less for Model~E. Sulfur and nitrogen containing compounds are other potential sources for aerosols and we plan to explore their photochemistry in a later paper.  The CEGPs are extremely close to the parent star; in such an extreme environment, the C$_x$H$_y$ compounds will be lost either primarily by reactions with atomic hydrogen or also by photolysis. The production of atomic hydrogen is a consequence both of the high temperatures that allow the presence of H$_2$O vapor and of the high UV flux that causes photolysis of H$_2$O.  Therefore, the lifetime of the C$_x$H$_y$ compounds in the atmospheres of CEGPs is predicted to be much shorter than that on Jupiter.  The lifetimes of the hydrocarbons are $\\lesssim 10^3$~s, which are significantly shorter than the simulated circulation timescale of $\\sim$day \\citep{SG02,Cetal03}. Hence the abundances of the hydrocarbons will be affected by a factor of `a few' through the relatively longer lifetime of the atomic hydrogen ($\\sim$1 day, Liang et al. 2003).  The condensation temperatures for hydrocarbons (e.g., C$_4$H$_2$ and C$_4$H$_{10}$) are below 200~K at $\\sim$1~mbar (Moses et al. 2000). These temperatures are far colder than expected in the atmospheres of CEGPs \\citep{SS00,Betal02,Fetal03}. Nevertheless, we verified this by considering the saturation profiles together with the the temperature profiles and found that the required saturation pressure for CEGPs is far more than that present in the atmospheres.   Using the measured Rayleigh scattering cross sections of He and H$_2$ \\citep{CD65,FB73}, the pressure level with optical depth unity is $\\sim$1~bar at 300~nm and increases rapidly at longer wavelengths (Rayleigh scattering cross section $\\propto \\lambda^{-4}$). Without the shielding from the atmospheric aerosols and in the absence of high-temperature condensate clouds, we may be able to observe the atmospheric composition at short wavelengths up to the Rayleigh scattering limit.  In this letter, we have emphasized photolytically driven processes involving neutral species. We have not considered the possibility of ion-neutral chemistry, such as that found in the polar region of Jupiter \\citep{Wetal03}. This may be important in the atmospheres of CEGPs if the planet possesses a magnetic field. If the hydrocarbon aerosols can be formed in the polar region, then global circulation will redistribute them to lower latitudes.  Stellar wind may be another source of energetic charged particles that could result in the formation of aerosols. Another subject not addressed in this work is the formation of aerosols by heterogeneous nucleation in the presence of pre-existing solid dust grains. In this case, the formation of aerosols would be sensitive to the amount of dust particles in the atmosphere.  Additionally, we find that the mixing ratios of C, O, S, and C$_2$H$_2$ (other than H) are high at the top of the atmosphere, implying that these particles can readily escape. The recent detection of C and O in the extended upper atmosphere of HD~209458b by \\citet{Vetal04} supports this assertion and we comment that hydrodynamically escaping atmospheric species will yield new information on the evolution of CEGPs."
    },
    "0402/astro-ph0402437_arXiv.txt": {
        "abstract": "{We report the detection of an X-ray flare on the Bp star $\\sigma$\\,Ori\\,E with the ROSAT high resolution imager (HRI). The flare is shown to have  likely occurred on the early-type star, rather than on an hypothesized  late-type companion. We derive flare parameters such as total energy release, coarse estimates of size and density, and also present arguments for a magnetic origin of the flare.  We place our observations in the context of a magnetic character of Bp-type stars and speculate on a common physical basis and connection between Bp and Be stars.} ",
        "introduction": "For cool stars with outer convective envelopes, X-ray emission is generally considered a proxy indicator of magnetic activity.   Systematic surveys of X-ray emission among cool stars in the solar neighborhood have shown the occurrence of X-ray emission for all stars of spectral type F through M on the main sequence (Schmitt et al. \\cite{Sch:al1}, Schmitt \\cite{Sch}, Schmitt \\& Liefke \\cite{Sch:Lie}).   The observed X-ray luminosities vary (from star to star) over almost four orders of magnitude, and the most important parameter governing the X-ray output level of a given cool star appears to be its rotation rate (Pallavicini et al. \\cite{Pal:al2}). The X-ray emission from cool stars is (often) time-variable.  Specifically, X-ray flaring is frequently observed among cool stars, and in particular, X-ray flaring has been observed for all types of cool stars.  Such X-ray flaring is readily explained by the sudden release of magnetic energy and its ultimate conversion into heat and radiation. The X-ray properties of cool stars have to be contrasted to the X-ray properties of hot stars without outer convective envelopes.  To avoid complications from the possibility of X-ray emission from interacting winds in O-star binary systems, we here restrict our attention to single stars. Similarly to cool stars, all single hot stars appear to be X-ray emitters, at least down to a spectral type of around B2.  In contrast to cool stars, the X-ray luminosity of hot stars appears to be independent of rotation; rather their X-ray emission is characterized by the relation L$_X$/L$_{bol}$ $\\approx$10$^{-7}$ (cf., Pallavicini et al. \\cite{Pal:al2}). Thus hot stars can be very strong X-ray sources in terms of their total X-ray luminosity L$_X$, however, in comparison to their total radiative output only small fractions of their total luminosity are radiated at X-ray energies.  The X-ray emission from hot stars is usually attributed to instabilities in their radiatively driven winds, leading to stochastic velocity fields and shocks (cf., Lucy \\cite{Luc}).  Surprisingly, despite the stochastic nature of the X-ray production process in hot stars, their X-ray emission level appears to be very stable (cf.,  Bergh\"ofer \\& Schmitt \\cite{Ber:Sc2}, \\cite{Ber:Sc3}) and convincing reports of variable X-ray emission from early type stars are rare. For example, Bergh\"ofer \\& Schmitt (\\cite{Ber:Sc4}) report long-term spectral variability in the hot supergiant $\\zeta$ Ori, which they interpret as a propagating shock wave, and Gagn\\'e et al. (\\cite{Gag:al1}) report a periodic variation of the X-ray flux of $\\theta ^1$\\,Ori\\,C, which appears to be modulated with the star's rotation period of 15.4 days.  An actual X-ray flare in the B2e star $\\lambda$~Eri has been reported by Smith et al. (\\cite{Smi:al1}); an increase in count rate by about a factor of 6 over quiescent values over about 24 hours was observed leading to an overall energy release of 2 $\\times$ 10$^{36}$ erg.  In general, magnetic fields are thought to be unimportant for the dynamics of the outer envelopes of hot stars.   While there are no reasons not to assume the presence of magnetic fields on such stars, the typically available magnetic field strength upper limits of a few hundred Gauss show that the magnetic field pressure is much smaller than the ram pressures of their winds (Linsky \\cite{Lin}). Successful measurements (rather than upper limits) of magnetic fields in early-type stars are rare. One of the few examples is $\\theta ^1$\\,Ori\\,C, where a magnetic field of about 1.1\\,kG has been detected to be modulated with the rotation period of 15.4 days (Donati et al. \\cite{Don:al1}). The prototypical He-strong He-variable star $\\sigma$\\,Ori\\,E (HD\\,37479; spectral type B2Vp) is known to have a strong magnetic field of about 10 kG polar field strength (Landstreet \\& Borra, \\cite{Lan:Bor}). A weak stellar wind of some $10^{-10}M_{\\sun}$\\,y$^{-1}$ is channeled along magnetic field lines. In closed magnetic loops matter is captured, forming torus-shaped clouds (or a ring) in the plane of the magnetic equator (Groote \\& Hunger \\cite{Gro:Hu3} (GH1), Shore \\& Brown \\cite{Sho:Bro}). These clouds are magnetically coupled to the star, thus corotate and occult the star twice during one rotational period leading to absorption in the Stroemgren u-band (Hesser et al., \\cite{Hes:al1}), the higher Balmer lines (Groote \\& Hunger, \\cite{Gro:Hu1}), and in the strong resonance UV lines of C\\ion{IV}, Si\\ion{IV}\\ (Smith \\& Groote, 2001).  When viewed face on, the clouds exhibit redshifted emission in H$_{\\alpha}$ and some of the above metal lines. The clouds extend to about 6 $R_{*}$ and are believed to release matter at their inner boundary back to the stellar surface, thus increasing the He abundance near the magnetic equator (Groote \\cite{Gro}). First results from 2D MHD simulations by ud-Doula \\& Owocki (\\cite{udD:Owo}) support this scenario.  Observations of $\\sigma$\\,Ori\\,E at radio wavelengths (6 cm, Drake et al., \\cite{Dra:al1}) are consistent with the assumption of synchrotron radiation by mildly relativistic electrons. X-rays from $\\sigma$\\,Ori\\,E have been observed by Bergh\"ofer \\& Schmitt (\\cite{Ber:Sc1}) at a constant level of about $8.2\\times 10^{-3}$ counts\\,sec$^{-1}$ using the ROSAT high resolution imager (HRI).  The origin of this X-ray emission is unclear.  Converting the observed X-ray count rate to an X-ray flux yields a spectrum-dependent estimate of approximately $2\\times 10^{-13}$ erg\\,cm$^{-2}$\\,sec$^{-1}$ (appropriate for $T\\approx 20$\\,MK and $N_{\\rm H}\\approx 10^{21}$\\,cm$^{-2}$), which corresponds to a fraction $\\approx 3.9\\ 10^{-7}$ of the total bolometric flux of $\\sigma$\\,Ori\\,E.  $\\sigma$\\,Ori\\,E thus lies in line with the typical values observed from O-type stars.  Alternatively, the magnetic geometry of $\\sigma$\\,Ori\\,E might be relevant and the X-ray emission might originate from frictional heating in the wind when helium and hydrogen decouple from the driving metals (Krti\\v{c}ka J.\\& Kub\\'at, \\cite{Kri:Kub}) and/or from the impact of wind particles into the ring/clouds (Babel \\& Montmerle, \\cite{Bab:Mon}). Since the above described ring is constantly filled with new material from the star, matter also has to be released and it is natural to assume magnetic field line reconnection with heating to some 10$^{6}$~K in analogy to the Earth's magnetosphere. This release of matter is not expected to occur continuously, but rather sporadically,  leading to variable X-ray emission.  However, in contrast to these expectations, first observations (duration about 15 ksec) with the ROSAT HRI (Bergh\"ofer \\& Schmitt \\cite{Ber:Sc2}, \\cite{Ber:Sc3}) did not show any significant temporal variations in the X-ray flux from $\\sigma$\\,Ori\\,E. Since then more data on $\\sigma$\\,Ori\\,E (about 75 ksec distributed over more than 30 days) were taken with the ROSAT HRI within the context of another program and it is the purpose of this paper to present (Sect.\\,2) and discuss these observations (Sect.\\,3). We determine basic flare properties and will in particular compare the $\\sigma$\\,Ori\\,E data with the flare found in $\\lambda$~Eri (Smith et al. \\cite{Smi:al1}).  Our conclusions will be drawn in Sect.\\,4, and finally we speculate on a possible connection between Bp and Be stars. ",
        "conclusions": "\\subsection{Unseen companions ?}  Pallavicini et al. (\\cite{Pal:al1}) attribute a flare observed with XMM-Newton apparently from $\\sigma$\\,Ori\\,E to an unseen companion. While the stellar cluster around $\\sigma$ Ori\\,AB is indeed very young and young stars are known to be capable of producing strong X-ray flares, $\\sigma$\\,Ori\\,E most probably does not belong to the $\\sigma$\\,Ori cluster. Hunger et al. (\\cite{Hun:al1}) derived fundamental parameters $T_{\\rm eff}$ and log~$g$ for the stars $\\sigma$\\,Ori\\,E and $\\sigma$\\,Ori\\,D, using high resolution CASPEC data. The latter star turned out to lie near the zero age main sequence (log~$g=4.3$) with a spectroscopic distance $d=370$~pc, which is in excellent agreement with the Hipparcos distance of $d = 352$\\,pc for $\\sigma$\\,Ori\\,AB, while for $\\sigma$\\,Ori\\,E log~$g=3.95$ and a distance of $d = 640$\\,pc was reported, clearly identifying $\\sigma$\\,Ori\\,E as a background star. Using evolutionary tracks from Schaller et al. (\\cite{Sch:al0}) with $T_{\\rm eff}=22500$~K and log~$g=3.95$, $\\sigma$\\,Ori\\,E is a 9~M$_{\\sun}$ star with an age of 17~Myr. Thus both, distance and age make a membership to the $\\sigma$\\,Ori cluster unlikely.  The angular resolution of our HRI data is much better compared to the XMM-data and the relative X-ray positions of $\\sigma$\\,Ori\\,AB and E agree very well with their relative optical positions ($\\rho_{\\rm opt}= 41.36 $\\,arcsec, $\\rho_{\\rm X}= 41.07 \\pm 0.27$\\,arcsec, and position angle $\\alpha_{\\rm opt} = 61.55$\\,deg, $\\alpha_{\\rm X} = 61.60 \\pm 0.27$\\,deg). Because of this excellent agreement of X-ray and optical positions, we consider a chance alignment of  $\\sigma$\\,Ori\\,E with a low mass member of the closer $\\sigma$\\,Ori cluster unlikely.  A positional agreement of only 0.3 arcsec translates into a projected distance of 300\\,au of a hypothesized companion of $\\sigma$\\,Ori\\,E. Thus $\\sigma$\\,Ori\\,E could be a spectroscopic binary itself. However, in the high-resolution, high-signal-to-noise FEROS spectra (SNR$\\approx$250) of $\\sigma$\\,Ori\\,E (c.f. Reiners et al. \\cite{Rei:al1}) no sign of any companion can be found. We estimate that any G-star hidden in the spectrum must be smaller than about $0.2\\,R_{\\sun}$ with correspondingly larger limits for M-type stars. If an hypothesized companion of $\\sigma$\\,Ori\\,E is indeed very young (say 2 Myrs), its radius should be far larger and therefore be visible in the spectrum, contrary to what is observed. Further, radial velocity measurements by Groote \\& Hunger (\\cite{Gro:Hu2}) show that any RV velocity variations would have to stay below 1 km\\,s$^{-1}$. And finally, absolute flux measurements from UV to IR, taken from a region encompassing a substantial volume around $\\sigma$\\,Ori\\,E, agree very well with the fluxes calculated from appropriate model atmospheres and provide no evidence whatsoever for a hidden companion or a foreground star possibly belonging to the $\\sigma$\\,Ori system. Nevertheless, a, say, $\\le 0.5 M_{\\sun}$ main sequence (i.e. ``old'') star at $\\approx$300\\,au might be hidden in the spectroscopic and photometric data, and could thus be held responsible for the  observed X-ray emission. Such a star, however, would have an age of $\\approx$\\,20\\,Myr rather than 2\\,Myr, and hence be comparable to clusters like IC\\,2602 (Randich et al.\\cite{Ran:al1}) or IC\\,2391 (Patten \\& Simon \\cite{Pat:Sim}). G and K-type stars in those clusters show X-ray luminosities about a decade lower than observed for $\\sigma$\\,Ori\\,E. Further, the X-ray flux of $\\sigma$\\,Ori\\,E is quite constant outside  the flare, which is rather atypical for a very young late type star. Also, the observed flare energetics are unusual (albeit not unthinkable) for a 20\\,Myr late type star. Attributing the observed quiescent and flaring X-ray emission from $\\sigma$\\,Ori\\,E to a physical companion would thus imply a rather unusual companion star; note that $\\sigma$\\,Ori\\,E is already quite unusual by itself. Therefore, we consider this possibility again unlikely. For the following we therefore assume instead that the observed X-rays and flares are produced in the circumstellar environment of $\\sigma$\\,Ori\\,E.  \\subsection{Wind-shock X-ray emission}  It seems reasonable to compare $\\sigma$\\,Ori\\,E to other similar stars. As far as the quiescent flux of $\\sigma$\\,Ori\\,E is concerned, it is comparable to that observed from other early type stars.  However, $\\sigma$\\,Ori\\,E is a chemically peculiar star with a magnetic dipole field of about 10 kG polar field strength. X-ray emission has been observed from two other magnetic stars, i.e.,  $\\theta^1$\\,Ori\\,C, which is much hotter ($T_{\\rm eff}\\approx 45\\,000\\,$K), and IQ Aur which is much cooler ($T_{\\rm eff}\\approx 15\\,000\\,$K) than $\\sigma$\\,Ori\\,E ($T_{\\rm eff} = 23\\,500\\,$K).   Their X-ray production can be successfully explained by the MCWS model (Babel \\& Montmerle \\cite{Bab:Mo2}, Donati et al. \\cite{Don:al1}); in particular, this model provides a natural explanation for the observed periodicity in the X-ray flux of $\\theta^1$\\,Ori\\,C. X-ray flaring for B-type stars  has been reported only for $\\lambda$\\,Eri (spectral type B2e; Smith et al. \\cite{Smi:al1}), however, $\\lambda$\\,Eri is so far not known to be a magnetic star.  The central feature of the MCWS model of Babel \\& Montmerle is the channeling of the stellar wind emanating at high magnetic latitudes by the magnetic field.  The magnetic field forces the high-latitude wind into the magnetic equator plane, where the the wind streams from both hemisphere collide. The shocked gas leads to the observed X-ray emission in the post-shock region; the gas eventually cools and forms a thin cooling disk in the magnetic equator plane.  This model would appear to provide an adequate description for $\\sigma$\\,Ori\\,E at first sight, too, however, the constant X-ray flux of $\\sigma$\\,Ori\\,E (cf, \\ref{fig4}) has to be contrasted with the periodic variation of $\\theta^1$\\,Ori\\,C. In the latter case the phase variation is interpreted by Donati et al. (\\cite{Don:al1}) as a partial eclipse of the X-ray emitting post-shock region near to the stellar surface at a distance of about $1.5\\,R_*$. The magnetic field of $\\theta^1$\\,Ori\\,C is much smaller ($\\approx 1.1$ kG) than that of $\\sigma$\\,Ori\\,E ($\\approx 10$ kG), while, on the other hand, the wind from $\\theta^1$\\,Ori\\,C is much stronger due to the intense O star radiation field. Consequently, one would expect a much stronger magnetic confinement in the case of $\\sigma$\\,Ori\\,E, the cooling disk should be denser and extend further out (up to 6 $R_*$ as observed). Therefore, X-ray absorption at magnetic pole phases with the disk viewed face-on might be expected. The magnetic geometry of $\\sigma$\\,Ori\\,E ($i = 54$ deg, $\\beta = 67$ deg) implies that twice during each rotation cycle the disk is viewed edge-on, when X-ray emission should be at maximum, while at least at one phase ($\\Phi = 0.73$) the disk is viewed face-on, where X-ray emission should be at a minimum. Such a variation is clearly not observed (cf., \\ref{fig4}), and the model successfully explaining the phase variation of the X-ray flux of $\\theta^1$\\,Ori\\,C does not appear to be applicable for the case of $\\sigma$\\,Ori\\,E.  What is the reason for the failure of this model describing well the X-ray production in a hotter as well as in a cooler magnetic star? The mass loss of the three stars considered, i.e., IQ Aur, $\\sigma$\\,Ori\\,E, and $\\theta^1$\\,Ori\\,C increases from about 10$^{-11} M_{\\sun}$\\,y$^{-1}$ over some 10$^{-10}M_{\\sun}$\\,y$^{-1}$ to about some 10$^{-7}\\,M_{\\sun}$\\,y$^{-1}$, respectively. The low mass loss rate for IQ Aur is the result of non-solar abundances in the wind.  Helium (because it is partly neutral) will decouple already in the photosphere, and thus never leaves the stellar surface. Hydrogen decouples near the stellar surface before reaching escape velocity, and will thus be re-accreted if its velocity is too small to reach the magnetic equator. The He-depletion at the magnetic pole (Babel \\&Montmerle \\cite{Bab:Mon}) seems more likely to be an H-enrichment due to the re-accretion of hydrogen falling back to the surface (see Groote \\cite{Gro}). The metals instead receive sufficient momentum from the radiation field to be accelerated against gravity to velocities of more than $\\approx 800$ km\\,sec$^{-1}$, and have therefore sufficient kinetic energy to produce the observed X-ray emission. In the case of the O star wind of $\\theta^1$\\,Ori\\,C, radiation pressure on the driving metals is sufficient to accelerate both metals and the passive plasma, which is coupled to the metals through Coulomb forces, to nearly twice the escape velocity ($\\approx 2500$ km\\,sec$^{-1}$). In the transition region between massive O star winds, where metals, hydrogen, and helium are fully coupled, and the pure metallic winds of the cool B stars with essentially no coupling between metals and hydrogen and helium, the mass loss rate decreases with decreasing effective temperature and decreasing wind velocities.  Hunger \\& Groote (\\cite{Gro:Hu4}) first proposed that the He-enrichment in the atmospheres of He-variable stars may be the result of He-decoupling in the winds of these stars.  After decoupling, helium will be re-accreted to form He-rich spots on the stellar surface. Later they extended the idea to He-weak stars (Hunger \\& Groote, \\cite{Hun:Gro}, see also Groote, \\cite{Gro}), where hydrogen is also decoupled in the wind and may be re-accreted at the magnetic poles. Such a decoupling of passive elements seems inevitable if wind velocities drop below escape velocity. The coupling becomes weaker with the decrease of density with increasing distance from the star. Once helium (or hydrogen) is decoupled, the remaining wind particles may be efficiently accelerated further because of the substantial decrease in mass left for acceleration. In the case of $\\sigma$\\,Ori\\,E the initial wind velocities are probably low, additionally depending on magnetic latitude, and are perhaps of the order of only $200 - 300$\\,km\\,sec$^{-1}$. Therefore, heating in a post-shock region may be insufficient to produce any significant X-ray emission, so that other heating processes are required. Since no phase variation is observed (Bergh\"ofer \\& Schmitt, \\cite{Ber:Sc1}, and this work), we conclude that the source of the observed X-rays is not occulted either by the star nor by the disk, and must therefore be located outside the closed loops of the magnetosphere. This assumption also fits the finding that wind absorption seen in the resonance lines of \\ion{C}{iv} and \\ion{Si}{iv} does not vary with phase either (Groote \\& Hunger \\cite{Gro:Hu4}, Smith \\& Groote \\cite{Smi:Gro}). Groote \\& Hunger considered an expanding corona outside a wind-fed magnetosphere, which releases matter more or less regularly from the corotating clouds.  \\subsection{Size and location of the flaring region}  In view of the above considerations it is likely that also the site of the X-ray flare is located rather far away from the stellar surface.  Unfortunately, the essential light curve parameters such as peak flux and decay time are only rather indirectly available from the ROSAT data (see above), and the flare temperature is unknown.  For the following we adopt an {\\it ad hoc} temperature of 10$^7$\\,K, which is a typical temperature for cool star flares. Fortunately, the cooling function of optically thin plasma is rather insensitive to the precise value of temperature in that range.   If we further adopt the observed peak X-ray flux of 6.1 $\\times$ 10$^{31}$ erg\\,sec$^{-1}$ as the actual peak flux,  we compute a minimal emission measure of $EM_{\\rm min} \\approx 10^{54}$\\,cm$^{-3}$.   As to the decay time, $\\tau_{\\rm decay}$, it could be as long as $\\tau_{\\rm decay} \\approx $ 40 ksec if indeed the same morphology as for the flare on $\\lambda$\\,Eri applies, but also much shorter. If we adopt $\\tau_{\\rm decay} = $ 40\\,ksec and further assume that the flare cools only through radiative cooling we can compute a characteristic electron density $N_{\\rm e}$ from the relation \\begin{equation} N_{\\rm e} = \\frac {3 k T} {P(T) \\tau_{\\rm decay,}} \\end{equation} where $k$ denotes Boltzmann's constant and $P(T)$ the plasma cooling function. For $T \\approx$ 10$^7$\\,K, we find $P(T)$ $\\approx$ 10$^{-23}$ \\,erg\\,cm$^3$\\,sec$^{-1}$and hence $N_{\\rm e} \\approx 10^{10}$\\,cm$^{-3}$. We consider this a minimal plasma density, since both higher temperatures (which are usually observed in cool star flares) and shorter decay times (40 ksec is close to the maximally possible) will lead to higher densities. Some questions remain about the correct cooling function $P(T)$; the values used refer to a fully ionized plasma with cosmic abundances in collisional equilibrium, and these assumptions may not fully apply. Adopting then $N_{\\rm e} = 10^{10}$\\,cm$^{-3}$ and  $EM_{\\rm min} = 10^{54}$\\,cm$^{-3}$, we compute the flaring volume $V_{flare} = 10^{34}$\\,cm$^{3}$, which corresponds to a sphere with radius $1.3\\times 10^{11}$\\,cm or $0.4\\,R_{*}$.  Thus the flaring region is indeed small compared to the star itself, and its presumed distance from the stellar surface. We emphasize that other cooling mechanisms will imply even larger densities and correspondingly smaller volumes.  Assuming a field scale size of $1 R_{*}$, the magnetic field strength at $6 R_{*}$ will be about 50\\,G yielding a pressure of somewhat below 100\\,dyn\\,cm$^{-2}$; the thermal pressure of plasma with the above assumed temperatures and densities is about 30\\,dyn\\,cm$^{-2}$, i.e., of the same order.   Calculating the speed $v_{\\rm wind}$, that a plasma with density of $N_{\\rm e} = 10^{10}$\\,cm$^{-3}$ must have to produce a ram pressure of the same magnitude, results in a value of  $v_{\\rm wind} \\approx 10^8$\\,cm\\,sec$^{-1}$, which would be a plausible wind speed.  Assuming a plasma ``cloud'' 6 $R_{*}$ away from the surface with a thermal pressure of 30 dyn\\,cm$^{-2}$ is not unreasonable.  Groote \\& Hunger (\\cite{Gro:Hu3}) derive an electron densitiy of $\\approx 10^{12}$\\,cm$^{-3}$ at temperatures $\\approx 10^{4}$\\,K, yielding thermal pressures of the same order as the hot X-ray emitting gas. Also, the assumption of magnetic reconnection is very plausible. Adopting the above derived values for $V_{\\rm flare}$ and $E_{\\rm tot}$ we can compute the total reconnected field component $B_{\\rm recon}$, which is found to be 60\\,G, i.e., again of the same order as the (vacuum) dipole field. Thus, a magnetic reconnection scenario requires field strengths and pressures which are entirely plausible.  \\subsection{Heating}  The most commonly adopted heating process for single O and B stars are shocks formed by instabilities in the radiatively driven winds of these stars. Porter \\& Skouza (\\cite{Por:Sk1}) used 1D hydrodynamic calculations to model the case of $\\sigma$\\,Ori\\,E. According to Porter \\& Skouza (\\cite{Por:Sk1}) hydrogen decouples to form  shell-like structures which are then shocked and become the source of the observed X-rays. Krti\\v{c}ka \\& Kub\\'at (\\cite{Kri:Kub}) used 2D hydrodynamic calculations (albeit not fully self-consistent, and neglecting magnetic fields and rotation), and found that decoupling of helium is possible in B3 stars when modeling the outflow as a four component gas. Similarly, their models showed decoupling in B5 stars in the context of a 3 component gas. In both cases heating processes are important for the remaining wind component, and temperatures of several $10^6$\\,K may be found at larger distances. In the case of magnetic stars the wind will be slowed down additionally by the magnetic latitude dependent channeling of the magnetic field, and their results may also be valid for stars with somewhat higher temperatures. For the case of $\\sigma$\\,Ori\\,E the heating of the ionic component would then occur outside the closed magnetic field lines in an approximately radial magnetic field line configuration, forming a corona as already assumed by Havnes \\& Goertz (\\cite{Hav:Goe}) and GH1. In that case only a small part of the spherical X-ray emitting region with a diameter in excess of $6 R_*$ will be occulted by the dense gas ring. This view is also supported by VLBI observations of $\\sigma$\\,Ori\\,E presented by Phillips \\& Lestrade (\\cite{Phi:Les}), who find a scale size of 6 to 10 $R_*$ for the radio emitting regions, once the new (larger) distance towards $\\sigma$\\,Ori\\,E is taken into account (see also Hunger et al. \\cite{Hun:al2}). Such heating processes will not be present in a star like $\\theta^1$\\,Ori\\,C whose wind is dense and whose wind components are all well coupled. Within the context of such a scenario a plausible explanation for the observed X-ray flare(s) can thus be % found. The permanently ongoing filling of the ring/clouds by the wind will increase the plasma ram pressure, while the magnetic pressure stays constant. Once the ram pressure exceeds the magnetic pressure, matter cannot remain confined by the field, and must be released from the magnetosphere.  The reconnection of magnetic field lines occurring during this release then provides the additional heating with ensuing X-ray emission. GH1 estimated the time required to refill the clouds using the mass loss and found a time scale of the order of about a month."
    },
    "0402/astro-ph0402089.txt": {
        "abstract": "We re-examine the outer gap size by taking  the geometry of dipole magnetic field into account. Furthermore, we also consider that instead of taking the gap size at half of the light cylinder radius to represent the entire outer gap it is more appropriate to average the entire outer gap size over the distance. When these two factors are considered, the derived outer gap size $f(P, B, <r>(\\alpha))$ is not only the function of period ($P$) and magnetic field ($B$) of the neutron star, but also the function of the average radial distance to the neutron star $<r>$; which depends on the magnetic inclination angle ($\\alpha$). We use this new outer gap model to study $\\gamma$-ray luminosity of pulsars, which is given by $L_{\\gamma} = f^3(P, B, <r>(\\alpha))L_{sd}$ and $L_{sd}$ is the pulsar spin-down power, as well as the death lines of $\\gamma$-ray emission of the pulsars. Our model can predict the $\\gamma$-ray luminosity of individual pulsar if its $P, B$ and $\\alpha$ are known. Since different pulsars have different $\\alpha$, this explains why some $\\gamma$-ray pulsars have very similar $P$ and $B$ but have very different $\\gamma$-ray luminosities. In determining the death line of $\\gamma$-ray pulsars, we have used a new criterion based on concrete physical reason, i.e. the fractional size of outer gap at the null charge surface for a given pulsar cannot be larger than unity. In estimate of the fractional size of the outer gap, two possible X-ray fields are considered: (i) X-rays are produced by the neutron star cooling and polar cap heating, and (ii)X-rays are produced by the bombardment of the relativistic particles from the outer gap on the stellar surface (the outer gap is called as a self-sustained outer gap). Since it is very difficult to measure $\\alpha$ in general, we use a Monte Carlo method to simulate the properties of $\\gamma$-ray pulsars in our galaxy. We find that this new outer gap model predicts many more weak $\\gamma$-ray pulsars, which have typical age between 0.3-3 million years old, than the old model. For all simulated $\\gamma$-ray pulsars with self-sustained outer gaps, $\\gamma$-ray luminosity $L_{\\gamma}$ satisfies $L_{\\gamma}\\propto L^{\\delta}_{sd}$; where the value of $\\delta$ depends on the sensitivity of the $\\gamma$-ray detector. For the EGRET, $\\delta$ is $\\sim 0.38$ whereas $\\delta$ is $\\sim 0.46$ for the GLAST. For $\\gamma$-ray pulsars with $L_{sd} \\lesssim L_{sd}^{crit}$, $\\delta$ is  $\\sim 1$. $L^{crit}_{sd} = 1.5 \\times 10^{34} P^{1/3} ergs^{-1}$ is determined by $f(<r>\\sim r_L) =1$. These results are roughly consistent with the observed luminosity of $\\gamma$-ray pulsars. These predictions are very different from those predicted by previous outer gap model, which predicts a very flat relation between $L_{\\gamma}$ and $L_{sd}$. ",
        "introduction": "High-energy emission models for rotation-powered pulsars are generally divided into polar gap and outer gap models. In polar gap models, charged particles are accelerated in charge-depleted zones near the pulsar's polar cap and $\\gamma$-rays are produced through curvature-radiation induced $\\gamma$-B pair cascade (e.g. Harding 1981; Daugherty \\& Harding 1996; Zhang \\& Harding 2000) or through Compton-induced pair cascades \\cite{dermer94}. In the outer gap models, it is generally accepted that a magnetosphere of charge density \\begin{equation} \\rho_0\\approx {{\\bf{\\Omega}}\\cdot{\\bf{B}}\\over 2\\pi c} \\end{equation} surrounds a rotating neutron star with magnetic field $B$ and angular velocity $\\Omega$ (Goldreich \\& Julian 1969). The magnetospheric plasma is corotating with the neutron star within the light cylinder, at which the corotating speed equals the velocity of light and the distance from the spin axis is  $R_L=c/\\Omega$. In the corotating magnetosphere, the electric field along the magnetic field, $E_{||}={\\bf{E}}\\cdot{\\bf{B}}/B$, is nearly zero. However, the flows of the plasma along open field lines will results in some plasma void regions (where the charge density is different significantly from $\\rho_0$) in the vicinity of null charge surfaces where ${\\bf{\\Omega}}\\cdot{\\bf{B}}=0$ (Holloway 1973). In such charge deficient regions, which are called outer gaps, $E_{||}\\neq 0$ is sustained, electrons/positrons can be accelerated to relativistic energies and the subsequent high-energy gamma-ray emission and photon-photon pair production can maintain the current flow in the magnetosphere (Cheng, Ruderman \\& Sutherland 1976; Cheng, Ho \\& Ruderman 1986a, 1986b, hereafters CHRI and CHR II; Romani, 1996; Zhang \\& Cheng 1997; Hirotani 2001 ).  Based on known $\\gamma$-ray pulsars, the luminosity and conversion efficiency of $\\gamma$-rays in various models have been studied (for example, Harding 1981; Dermer \\& Sturner 1994; Rudak \\& Dyks 1998; Yadigaroglu \\& Romani 1995; Zhang \\& Cheng 1998). Observations by Compton Gamma-ray Observatory (CGRO) show that $\\gamma$-ray luminosity of rotation-powered pulsars is proportional to square root of the spin-down power (Thompson et al. 2001). Using the recent new polar gap models (e.g. Zhang \\& Harding 2000; Harding \\& Muslimov 2001; Harding, Muslimov \\& Zhang, 2002), Harding et al. (2002) have studied the deadline of $\\gamma$-ray pulsars based on the predicted luminosity of $\\gamma$-ray pulsars $L_{\\gamma}\\propto L^{\\delta}_{sd}$, where $\\delta \\sim 0.5$ when $L_{sd} \\gtrsim L^{break}_{sd}$ and $\\delta \\sim 1.$ when $L_{sd} \\lesssim L^{break}_{sd}$ respectively and $L^{break}_{sd} = 5 \\times 10^{33} P^{-1/2}$ erg/s.  For a rapid rotating pulsar, it is believed that its spin-down power, $L_{sd}$, is converted into radiation energy. Because the outer gap occupies only a part of the open field line region, the gamma-ray luminosity produced in the outer gap is a fraction of the spin-down power. It has been shown that the gamma-ray luminosity in the outer gap is proportional to $f^3$, {\\bf{i.e $L_{\\gamma}\\approx f^3L_{sd}$}} (CHR II; Zhang \\& Cheng 1997). In previous works, the factional size of outer gap only depends on the period and magnetic field for a gamma-ray pulsar. For example, the fractional size of the outer gap for Crab-like pulsars is $f\\propto B_{12}^{-13/20}P^{33/20}$ (CHR II). In the outer gap model described by Zhang \\& Cheng (1997), $f\\propto B^{-4/7}P^{26/21}$, the observed gamma-ray luminosities with energies greater than 100 MeV from the known gamma-ray pulsars except for the Crab pulsar can be explained approximately (Zhang \\& Cheng 1998). It should be noted that the model predicts that the $\\gamma$-ray pulsars with the same values of $(B/P)^{0.3}$ have same $\\gamma$-ray luminosities. For example, the ratio of $(B/P)^{0.3}$ for PSR B1055-52 to that for Geminga is $\\sim 0.9$, it means that the ratio of both gamma-ray luminosities is $\\sim 0.9$. However, the observed ratio of the luminosities with energies greater than 100 MeV of these two pulsars are $\\sim 8$ (Kaspi et al. 2000). Briefly, 7 known $\\gamma$-ray pulsars have rather different $\\gamma$-ray luminosity even their spin-down powers and ages are so similar (e.g. Geminga and PSR 1055-52). Their pulse shapes also differ so much. It is clear that there are other intrinsic parameters to control these observed properties (gamma-ray luminosity, pulse shape , spectrum etc).  In this paper, we re-study the gamma-ray emission from the outer gaps of the rotation-powered pulsars by using a new outer gap model. We follow the idea of self-sustained outer gap given by Zhang \\& Cheng (1997). However, we take the magnetosphere geometry as well as the average properties of the entire outer gap into consideration and show that the fractional size of the outer gap is a function of period, magnetic field and magnetic inclination angle. In fact, the effect of the inclination angle on the $\\gamma$-ray emission have been considered in other versions of the outer gap models. For example, Romani \\& Yadigaroglu (1995) and Yadigaroglu \\& Romani (1995) took the magnetic inclination angle into account in their outer gap models, Hirotani and his colleagues also included the magnetic inclination angle in their calculation of the outer gap model (e.g. Hirotani, 2001; Hirotani \\& Shibata 2001; Hirotani \\& Shibata 2002: Hirotani, Harding, Shibata 2003). Different from the treatment of Romani \\& Yadigaroglu (1995) who considered it in a less analytic way, we give a explicit expression for the fractional size of the outer gap.  In section 2, we describe the revised outer gap model. We estimate the $\\gamma$-ray luminosity for rotation-powered pulsars and compare them with the observed data in section 3. In section 4, we derive the death lines of the pulsars with outer gaps. Finally, we give briefly our conclusion and discussion. ",
        "conclusions": "After taking the geometry of the dipole magnetic field, we have given a revised version of the outer gap model given by Zhang \\& Cheng (1997). In the revised outer gap model, the fractional size of the outer gap is not only the function of period and magnetic field of the neutron star, but also the function of the radial distance ($r$) to the neutron star and the magnetic inclination angle ($\\alpha$). In other words, the fractional size of the outer gap has a form of $f(r,\\alpha)=f_i(P, B_{12})G_i(r,\\alpha)$, where $f_0(P, B_{12})$ is only the function of pulsar period and surface magnetic field, $G_i(r,\\alpha)$ changes with radial distance to the neutron star and the magnetic inclination angle and subscript $i$ represents the X-ray field considered. The fractional size of the outer gap is given by Eq. (\\ref{fsizepc}) in the X-ray field which is produced by the neutron star cooling and polar cap heating and by Eq. (\\ref{fsizeog}) in the X-ray field which is produced by outer gap heating. In this model, the fractional size of the outer gap has a minimum at the inner boundary for a given pulsar, increases with the radial distance along the last open field lines and then reaches its outer boundary where $f(r_b, \\alpha)=1$. In other words, the outer gaps of some pulsars do not start from the null charge surface to the light cylinder. We have shown the changes of the fractional size of the outer gap with the  inclination angle (see Fig.1). Further, we have found that the outer gaps of relative young pulsars such Vela can extend from the null charge surface to the light cylinder for any inclination angle , however, the outer gaps of some pulsars like as Geminga cannot extend to the light cylinder for a larger inclination angle (say $75^{\\circ}$).  In order to describe the average properties of high-energy radiation from a $\\gamma$-ray pulsar, we have defined an average radial distance $<r>$ (see Eq. (\\ref{rave})). Using Monte Carlo method described by Cheng \\& Zhang (1998) (also see Zhang, Zhang \\& Cheng 2000), we simulated two populations of $\\gamma$-ray pulsars whose energy fluxes are greater than EGRET threshold and GLAST threshold. In these simulations, we only consider the case of pulsars with self-sustained outer gap and uniform distribution of the inclination angles. We have plotted the change of $L_{\\gamma}$ with $L_{sd}$ (see Fig. 2). We also indicated the variation of $L_{\\gamma}/L_{\\gamma,0}$ with $L_{\\gamma}$ in Fig. 3, which show the importance of the inclination angle on $L_{\\gamma}$. In the model of Zhang \\& Cheng (1997), $L_{\\gamma}$ is independent on the inclination angle. Fitting the simulated results, we found that $L_{\\gamma}\\propto L^{0.38}_{sd}$ for EGRET threshold and $L_{\\gamma}\\propto L^{0.46}_{sd}$ for GLAST threshold. Compared with the observed data given by Thompson (2001), our simulated results are reasonable (see Fig. 4). In fact, current distance of the Vela pulsar (Cavareo,  et al. 2001) is less than that used by Thompson (2001) about a factor of two, which reduce the luminosity about a factor of four.  In Fig.4, we note that Geminga is not within the simulated population. Zhang \\& Cheng (2001) have used a three-dimension outer gap model to explain the phase-resolved spectra of Geminga. They discovered that in order to fit the observed $\\gamma$-ray data the solid angle $\\Delta \\Omega$ must be near 5$sr$. In Fig. 4, $\\Delta \\Omega = 1sr$ is used for all observed $\\gamma$-ray pulsars. If the larger solid angle is used, it brings Geminga within the simulated population. It should be pointed out that Yadigaroglu \\& Romani (1995) have studied the effect of $\\gamma$-ray beaming in their outer gap model (also see Romani 1996). Zhang et al (2000) also gave a approximate expression of the $\\gamma$-ray beaming fraction, which will be applied to estimate $\\gamma$-ray fluxes in our new model.  It is interesting to point out that both polar gap models and outer gap models predict $L_{\\gamma}=L_{sd}$ for low spin-down power pulsars. But the position of the break occurs at $5 \\times 10^{33} P^{-1/2} erg/s$ for the polar gap model (Harding et al. 2002) whereas the outer gap models predict a higher position at $1.5 \\times 10^{34} P^{1/3} ergs^{-1}$. For higher spin-down power pulsars, the polar gap models predict $L_{\\gamma}=L^{1/2}_{sd}$, whereas the outer gap cannot give precise prediction because the model $L_{\\gamma}$ also depends on a less well-known parameter $\\alpha$. The statistical predictions of the outer gap model give $L_{\\gamma}=L^{\\delta}_{sd}$, where $\\delta$ depends on the properties of $\\gamma$-ray detector. For example, $\\delta$ = 0.38 for the EGRET and $\\delta$ = 0.46 for GLAST.  According to our model, the fractional size of outer gap at the null charge surface for a given pulsar ($f(r_{in}, \\alpha)$) reaches a minimum. Therefore, the outer gap should exist only if $f(r_{in},\\alpha)\\le 1$. Averaging $f(r_{in}, \\alpha)$ on two possible (uniform and cosine) distributions of the magnetic inclination angles respectively, we have obtained the death lines of the pulsars with outer gaps in the two possible X-ray fields and the comparison them with the observed data in Figs.5 and 6.  Compared with the death line derived from the outer gap model of Zhang \\& Cheng (1997), the revised model predict that more pulsars will have their outer gaps and then emit high-energy photons.  We would like to make the conclusion as follows. Our results indicate that (i) the intrinsic parameters for explaining the observed $\\gamma$-ray properties of rotation-powered pulsars are the magnetic inclination angle, period and magnetic field. We have obtained a very concrete functional form of the prediction of gamma-ray luminosity which only depends on  these intrinsic parameters, the inclination angle could be known if the radio data is sufficiently good; (ii) the conversion efficiency for 7 known gamma-ray pulsars, that is rather scattered , can be explained in our revised model using the Monte Carlo method. Although our estimation of the outer gap size cannot precisely predict the thickness of individual pulsar when the magnetic inclination angle is poorly known, it is a reasonable estimation of the statistical properties of gamma-ray pulsars using the statistical method; (iii)Unlike those 7 observed gamma-ray pulsars, mature pulsars (ages~0.3-3 million years) can also be gamma-ray pulsars and their efficiency is insensitive to the inclination angle. Most importantly, their gamma-ray luminosity is proportional to the spin-down power, which can be tested by GLAST; (iv) The mean cut-off age of gamma-ray pulsars is increased by a factor of 3, therefore older gamma-ray pulsars (up to about 3 million years old) can move up to higher galactic latitude. Some unidentified EGRET gamma-ray sources could be mature pulsars with ages between 0.3-3 million years old pulsars predicted by this model. In fact, Parkes Observatory survey has discovered a large number of radio pulsars on the error boxes of EGRET unidentified gamma-ray point sources (Torres et al. 2003); and (v)  The previous works on death line based on the outer gap model (e.g. Chen and Ruderman 1993) did not give a detail physical reason rather they used a phenomenological approach. Here we proposed a new criterion based on concrete physical reason for the death line, which predicts more gamma-ray pulsars. (vi)Our model predicts many more weak $\\gamma$-ray pulsars with age between 0.3 \u0096 3 million years old than previous outer gap models and they satisfy $L_{\\gamma} \\propto L_{sd}$, which can be verified by GLAST.  Finally, we would like to remark that the gamma-ray luminosity formulae developed in this paper may not be able to explain the gamma-ray luminosity of individual pulsar. Two important factors, i.e. distance uncertainty and beaming fraction, which play crucial roles in determining the gamma-ray luminosity, have not been considered here. For example, the luminosity ratio between PSR B1055-52 and Geminga, which have same spin-down power of $3\\times 10^{34} erg~s^{-1}$, is about a factor of 6 (Thompson et al. 2001). But the model prediction (cf. Fig. 2) is no more than a factor of 3. In fact, the estimate of the distance of PSR B1055-52 is very uncertain, its distance can change from $\\sim$ 1.5 kpc (for example, Thompson et al. 2001) to a small value of $\\sim$500 pc (see, for example, \\\"{O}gelman \\& Finley 1993; Combi et al. 1997; Torres, Butt, Camilo 2001; Hirotani \\& Shibata 2002), which makes the ratio change from  $\\sim$ 10 to $\\sim$ 3. Recent distance estimate for PSR B1055-52 from the dispersion measure is $\\sim$ 0.72 kpc (see http: //rsd-www.nrl.navy.mil/7213/lazio/ne\\_model, also see Mignani, DeLuca \\& Caraveo 2003). However, it is well known that 25\\% error is common in determining the dispersion measure (MaLaughlin \\& Cordes 2000), which gives a factor of 2 uncertainties in the luminosity. Also the distance estimate of Geminga is known to have at least 25\\% uncertainties (Caraveo et al. 1996), which give another factor of 2 uncertainties in gamma-ray luminosity estimate. Other uncertainty results from the gamma-ray beaming fraction. In principle, both distance estimate and beaming fraction will be obtained more accurate. Then our model still predicts that the difference in the magnetic inclination angle can still cause a large difference in gamma-ray luminosity for pulsars with same spin-down power."
    },
    "0402/astro-ph0402571_arXiv.txt": {
        "abstract": "{ The Sunyaev-Zel'dovich (SZ) effect of galaxy clusters is a tool to measure three quantities: Compton parameter, electron temperature, and cluster peculiar velocity. However, a major problem is non-removed contamination by astrophysical sources that emit in the SZ frequencies. This includes interstellar dust emission, infra-red (IR) galaxies, and radio sources in addition to primary Cosmic Microwave Background (CMB) anisotropies. The three former contaminations induce systematic shifts in the three SZ parameters. In this study, we carefully estimated, both for a large beam experiment (namely Planck Surveyor) and a small beam experiment (ACT-like), the systematic errors that result if a fraction of the expected levels of emission from dust, IR galaxies, and radio sources remains non-removed. We found that the interstellar dust emission is not a major contaminant for the SZ measurement. Unfortunately, the IR and radio source-induced systematic errors may be extremely large. In particular the intra-cluster temperature and peculiar velocity will be determined inaccurately for Planck and ACT-like experiments, if only the frequency dependences are used for the cleaning. The Compton parameter is also affected by the astrophysical contaminations. The systematic errors in this case were a factor of 2 to 5 times larger than the expected statistical error-bar for Planck.  For the ACT-like experiment, the statistical error-bars were larger than in the case of Planck by a factor of about 5, and therefore the systematic shifts remain within about $50\\%$ of the statistical errors. We have thus shown that the systematic errors due to contaminating astrophysical emissions can be significantly larger than the statistical errors, which implies that future SZ surveys aiming at measuring cluster temperatures and peculiar velocities will not be able to do so on their own without including additional information like cluster shapes or follow-up observations. ",
        "introduction": "}  The well-studied Sunyaev-Zel'dovich (SZ) effect \\cite[]{sz72,sz80} has been observed towards a few tens of known galaxy clusters (see reviews by \\cite{carlstrom2002} and Birkinshaw 1999). It is a powerful tool for cosmological and cluster studies, and when combined with other observations (X-rays, optical, lensing) the SZ effect allows us to measure cosmological parameters such as the Hubble constant and matter density in the universe $\\Omega_{\\mathrm m}$ (e.g. \\cite{myers97}, \\cite{grego2001}, \\cite{reese2002}, \\cite{batt03}).  The SZ effect is independent of redshift and is therefore an excellent tracer of large scale structure formation and evolution. This property has inspired groups to propose blind SZ surveys and investigate the way they can probe the cosmological parameters (e.g. \\cite{bartlett94}, \\cite{barbosa96}, \\cite{holder2000}, \\cite{da_silva2000}, \\cite{kneissl2001}, \\cite{xue2001}).  As a consequence, several SZ experiments are either planned, under construction, or already observing. Some of these experiments are interferometric arrays like AMIBA \\cite[]{lo2001}, AMI \\cite[]{kneissl2001}, SZA~\\footnote{//astro.uchicago.edu/sza/} \\cite[]{mohr2002}. Others are single dish multi-frequency instruments like Planck surveyor\\footnote{//www.rssd.esa.int/Planck/}, SPT\\footnote{//astro.uchicago.edu/spt}, ACT\\footnote{//www.hep.upenn.edu/\\~{}angelica/act/act.html}, OLIMPO \\cite[]{masi2003}, SUZIE-II \\footnote{//www.stanford.edu/\\~{}schurch/suzie\\_instrument.html}, ACBAR\\footnote{//cosmology.berkeley.edu/group/swlh/acbar/}, MITO \\footnote{//oberon.roma1.infn.it/mito/}  The SZ effect can also be used to characterise the galaxy clusters themselves. \\cite{pointecouteau98} proposed  measuring the intra-cluster gas temperature from the relativistic corrections to the SZ. This is a particularly important issue in the context of future SZ surveys for which X-ray counterparts will not be easily available  and since X-ray temperatures have not been measured for all clusters. This is because X-ray temperature determination is expensive, and one cannot expect that all the many clusters to be observed in future SZ surveys will have X-ray temperature determination. The SZ effect can also be used to measure cluster radial peculiar velocities as suggested by \\cite{sz80}. This provides a new distance free measurement of the cluster velocities, which is very interesting in view of future SZ surveys. However this needs measurement of the intra-cluster temperature. So far, it is only for a handful of known rich clusters that attempts have been made to extract the intra-cluster gas temperature \\cite[]{hansen02} or to extract the peculiar radial velocities \\cite[]{holzapfel97,lamarre98,benson2003}.  It is fair to say that only upper limits have been obtained, and an actual temperature or peculiar velocity determination has still to be made.  The SZ effect is a potentially powerful cosmological probe; however, it is not free of contamination. As a matter of fact, SZ measurements are contaminated by other astrophysical sources that are mainly of two types: (i) due to our galaxy (free-free, synchrotron and dust emission from the milky way) or (ii) due to extra-galactic point sources (radio and infra-red (IR) galaxies), and even (iii) due to the Cosmic Microwave background (CMB) itself. The kinetic SZ and primary CMB fluctuations have the same frequency dependence, and therefore will a separation of the CMB signal and the kinetic effect only be possible through the relativistic corrections to the kinetic effect.  This would require both very high sensitivity (of the order $0.1 \\mu K$) and additional observing frequencies at the extrema of the relativistic corrections to the kinetic effect (near 200 and 500 GHz) and at its cross-over (near 300 GHz). The effect of the primary CMB anisotropies is well known (see for example \\cite{haehnelt96} and \\cite{aghanim97}). It induces a further uncertainty to the peculiar velocity and limits its accuracy. This effect intervenes mostly for large beam experiments like Planck, as small beam experiments are indeed less affected by the CMB which power is severely damped on a scale of about a few arcminutes.  These contaminations may be monitored in the context of pointed SZ observations (i.e. towards known clusters) or follow-up observations of likely clusters. However, systematic follow-up observations of all the clusters in the SZ survey will be very time-consuming. Ignoring contaminations modifies the effectiveness of a survey by the loss of certain clusters and the appearance of some other \"fake\" artificial clusters. For future surveys accurate knowledge of the completeness and reliability is thus crucial for extracting cosmological information. The question of survey selection function is being studied extensively for the Planck survey \\cite[]{white2003a,geisbusch2004,schafer2004}; however, the systematic effect from the assumed cluster structures still remains to be studied carefully \\cite[]{birkinshaw2004,hansen04a}.   Contamination of SZ measurements by extra-galactic sources, especially radio sources, is not a new problem (see for example reviews by \\cite{rephaeli95}, Birkinshaw (1999)). It can be due to radio emission of the galaxies in the cluster itself \\cite[]{ledlow96,cooray98,lin2002} or to foreground galaxies. More specifically, the emission of radio sources can dilute the SZ signal. \\cite{holder2002} estimated the expected dilution and its effect on the SZ power spectrum.  IR dusty galaxies whose emission dominates at high frequencies may also contaminate SZ measurements. In particular, gravitational lensing of dusty galaxies causes enhancement of the confusion noise, which is likely to affect SZ observations of nearby clusters \\cite[]{blain98}.  The effect of the point sources (both radio and dusty galaxies) have been recently revisited by \\cite{white2003}, who quantified it in terms of an equivalent noise and computed the associated power spectra. Also recently, \\cite{knox2003} investigated the effects of IR galaxy contamination on the statistical error-bars of the SZ parameters. Contamination of SZ clusters by radio sources at low frequencies that can be monitored by the interferometric arrays remains potentially a very important source of errors for single dish experiments. On the other hand, multi-frequency observations with single dish experiments should help in solving the problem of contamination.  In this study, we have focused on one category of SZ instruments planned for SZ surveying, namely the single-dish multi-frequency experiments. The SZ number counts that will be provided to us by these surveys will certainly probe and constrain the cosmological parameters. However, such instruments are theoretically able not only to measure the SZ effect amplitude  for each detected cluster (through the Compton parameter) but also its radial peculiar velocity and its gas temperature. We explored the capabilities of these future SZ surveys in terms of measuring the three cluster SZ parameters: Compton parameter, intra-cluster gas temperature, and radial peculiar velocity independently of other observations.  In a previous work \\cite[]{aghanim03}, we investigated the effects of the cluster parameter degeneracies on the of the above-mentioned cluster SZ parameters in terms of error-bars.  Here we were interested in the effects of the different sources of contamination on the measurement of these parameters. The contamination did not affect the statistical error-bars strongly.  The main question we are seeking to answer in this article is: {\\it How big are the systematic errors due to non-removable contamination?} We show that the systematic errors are very important. They are by far the dominant source of trouble for the future single-dish SZ blind surveys when these surveys are used alone to estimate cluster parameters of unresolved clusters. We start by presenting our models for the major astrophysical contaminants in Sect. 2, and we then discuss the SZ parameter extraction technique and the sample of galaxy clusters used in our study in Sect. 3.  In Sects. 4 and 5, we present the results for the three SZ experiments under consideration in our study, and discuss the results in Sect. 6.  We finally offer our conclusions in Sect. 7. ",
        "conclusions": "The future blind SZ surveys will be contaminated by astrophysical sources that can only be partly removed, such as interstellar dust emission, infra-red galaxies, and radio sources. Such non-removable contamination will induce a systematic shift in the derived SZ parameters.  We have shown that IR and radio source-induced systematic errors may be extremely large, virtually removing the possibility of measuring peculiar velocity and cluster temperature if using purely the SZ observations.  Also the systematic shift in the Compton parameter will be significantly larger than the expected statistical error-bars. Therefore these contaminations are potentially disastrous for future survey experiments, and must be considered very as very serious.  On the other hand, the situation may be less desperate if the contamination levels are reduced by complementary follow-up surveys.  As mentioned in Sect. \\ref{sec:req}, the radio follow-ups should be planned at, or close to, the SZ frequencies and the radio source variability requires simultaneous monitoring. CIB fluctuations need much better resolution than what is accessible now. One must therefore be realistic about future SZ surveys, and consider them as excellent tools for identifying high redshift clusters and not as tools for measuring peculiar velocity or cluster temperature."
    },
    "0402/astro-ph0402236_arXiv.txt": {
        "abstract": "We present \\chandra\\ observations of 17 optically-selected, X-ray weak narrow-line Seyfert 1 (NLS1) galaxies.  These objects were optically identified by Williams et al.~(2002) in the Sloan Digital Sky Survey Early Data Release, but were not found in the ROSAT All-Sky Survey (RASS) despite having optical properties similar to RASS--detected NLS1s.  All objects in this sample were detected by \\chandra\\ and exhibit a range of 0.5--2~keV photon indices $\\Gamma=1.1-3.4$.  One object was not detected in the soft band, but has a best--fit $\\Gamma=0.25$ over the full 0.5--8~keV range.  These photon indices extend to values far below what are normally observed in NLS1s. A composite X-ray spectrum of the hardest objects in this sample does not show any signs of absorption, although the data do not prohibit one or two of the objects from being highly absorbed.  We also find a strong correlation between $\\Gamma$ and $\\l1kev$; this may be due to differences in $L_{\\rm bol}/L_{\\rm Edd}$ among the NLS1s in this sample.  Such variations are seemingly in conflict with the current paradigm that NLS1s accrete near the Eddington limit.  Most importantly, this sample shows that strong, ultrasoft X-ray emission is not a universal characteristic of NLS1s; in fact, a substantial number may exhibit weak and/or low--$\\Gamma$ X-ray emission. ",
        "introduction": "\\citet{ostpog85} initially defined narrow-line Seyfert 1 galaxies (NLS1s) by their striking optical spectral characteristics: strong, narrow H$\\beta$ emission \\citep[later formally defined to be FWHM~$\\leq 2000 \\kms$ by][]{goodrich89}, weak [\\ion{O}{3}] relative to H$\\beta$, and strong \\ion{Fe}{2}.  These properties put NLS1s at one extreme end of the so--called \\citet{bg92} ``Eigenvector 1,'' thought to correspond to emission from lower-mass nuclear black holes coupled with near--Eddington accretion rates \\citep{boroson02}.  X-ray observations have revealed a strong soft X-ray ``excess'' in NLS1s \\citep[e.g.,][]{leighly99}, further bolstering the high $L_{\\rm bol}/L_{\\rm Edd}$---low--mass black hole hypothesis \\citep{pounds95,wang96}. Indeed, \\citet[][hereafter BBF96]{bbf96} found a possible anticorrelation between X-ray spectral slope and H$\\beta$ line width, with NLS1s generally having softer X-ray spectra than other AGN.  Ultrasoft X-ray selection has consequently proven to be an essential tool for the discovery of large numbers of NLS1s \\citep[e.g.][]{grupe00,grupe04}.  The disadvantage of selecting NLS1s solely upon their X-ray properties is that it can introduce into NLS1 samples a strong bias toward those exhibiting an ultrasoft excess \\citep{forster99}. Since NLS1s are primarily defined by their optical properties, the true nature of their X-ray emission is thus difficult to determine.  Though it is well known that \\emph{some} NLS1s are ultrasoft X-ray sources, previous samples of optically selected NLS1s were simply too sparse to determine how many, as well as whether or not a significant number of NLS1s have hard X-ray spectra.  With the advent of the Sloan Digital Sky Survey \\citep[SDSS;][]{york2000}, it is possible to build large catalogues of NLS1s with homogeneous selection criteria based on their optical spectra alone \\citep[see][hereafter WPM02]{wpm02}.  Of the 150 NLS1s listed in WPM02, 52 were detected in the ROSAT All-Sky Survey \\citep[RASS;][]{rass}.  Forty--five of these had sufficient counts in the 0.1--2.0~keV range to derive power-law photon indices ($\\Gamma$, where $N(E)\\propto E^{-\\Gamma}$) based on hardness ratios. Most of these objects were optically bright ($g \\la 18.5$), low--redshift ($z \\la 0.4$), and exhibited typical NLS1 photon indices of $\\Gamma \\ga 2.0$.  However, a substantial number of optically--bright, low--redshift NLS1s did \\emph{not} have RASS source identifications.  The optical spectra of these objects appear completely normal (within the limitations of the SDSS resolution and signal--to--noise) and only one or two are in regions of high Galactic \\ion{H}{1} column density, which could potentially obscure the X-ray flux.  It is thus possible that these objects represent a subset of NLS1s which are optically normal but X-ray weak.  Only a few such objects have previously been found; for example, RX J2217.9-5941 \\citep{grupe01b} and PHL 1811 \\citep{leighly01}. Since these SDSS NLS1s were not detected in the RASS, however, the nature and extent of their X-ray weakness could not be determined.  They could emit significant flux at higher energies not covered by the 0.1--2.4~keV ROSAT band, or they might also be ultrasoft X-ray sources but with substantially lower overall X-ray fluxes than typically seen in NLS1s (i.e., much higher $\\alpha_{\\rm{ox}}$).  Another possible explanation is variability, but it seems unlikely that such a large fraction ($\\sim 40\\%$) would be in an exceptionally low state during the RASS observations.  In reality, all three of these factors probably have some bearing on the X-ray weakness of these objects, but we cannot infer how many are affected by which factors, if any, from the existing data.  As a first step toward solving this puzzle, we have observed 17 of these optically--selected but RASS--undetected NLS1s with \\chandra.  Due to its excellent sensitivity and low background levels, \\chandra\\ is able to detect objects at far lower flux levels than ROSAT, and its large (0.5--8~keV) energy range allows detection of objects with harder X-ray spectra as well.  Our primary goals are (1) to detect these objects in X-rays, or set upper limits to their X-ray emission, and (2) to obtain rough estimates of $\\Gamma$ for the \\chandra--detected NLS1s.  Given this information, we can gain some insight as to which of the aforementioned scenarios (if any) sufficiently explain these ROSAT--unobserved NLS1s.  In this paper we present the results of these \\chandra\\ observations, and the possible implications for the NLS1s in our sample. ",
        "conclusions": "}  Through short--duration \\chandra\\ observations of RASS--undetected NLS1s, we have determined that six of the 17 objects have X-ray luminosities substantially lower than NLS1s with similar optical properties.  Of the brighter objects, at least two exhibit flux levels which should have been detectable by the RASS, indicating that their luminosities may have increased by a factor of two or more between the RASS and \\chandra\\ observations.  Many of the remaining bright objects were near or just below the RASS detection limit, and were most likely not seen due to smaller luminosity variations or Poisson noise. Across the entire sample, a strong correlation is seen between the X-ray spectral slope $\\Gamma$ and $\\l1kev$.  This is probably not entirely due to intrinsic absorption, since individual spectra of bright objects as well as a coadded spectrum of the faintest objects do not indicate high degrees of absorption (although one or two of the faintest hard--spectrum objects may be absorbed in X-rays but not at optical wavelengths). If $\\Gamma$ is indeed correlated with $L_{\\rm bol}/L_{\\rm Edd}$, then the $\\Gamma-\\l1kev$ relation suggests that variations in $\\l1kev$ are primarly due to differences in $L_{\\rm bol}/L_{\\rm Edd}$ among objects with comparatively similar black hole masses.  This interpretation is complementary to that of \\citet{dai04}, who find a similar relation but whose sample more likely includes objects with a large range of $\\mbh$ but $L_{\\rm bol}/L_{\\rm Edd}\\sim 1$.  These observations may hold important implications for the ``Eigenvector 1'' (Principal Component 1; PC1) paradigm posited by \\citet{bg92} and reinforced by \\citet{boroson02}.  In this picture, PC1 (which is primarily driven by an anticorrelation between [\\ion{O}{3}] and \\ion{Fe}{2}) is an indicator of $L_{\\rm bol}/L_{\\rm Edd}$. NLS1s typically lie at one extreme end of PC1---the end thought to correspond to the highest relative accretion rates.  Since $\\Gamma$ is thought to be related to $L_{\\rm bol}/L_{\\rm Edd}$, $\\Gamma$ and PC1 should be correlated; indeed, \\citet{brandt98} find such a correlation. However, the sample presented herein contains objects which from their optical spectra are at the supposed high--$L_{\\rm bol}/L_{\\rm Edd}$ end of PC1, yet also exhibit very low values of $\\Gamma$, as well as low inferred $\\l1kev/L_{\\rm Edd}$.  These extreme objects may indicate that while PC1 is usually correlated with $L_{\\rm bol}/L_{\\rm Edd}$, it may also be affected by orientation, black hole mass, or other physical drivers \\citep[as noted by][]{boroson04}.  This is not a completely new phenomenon; for example, BBF96 note that Mrk~507, with $\\Gamma=1.6\\pm 0.3$, has an unusually flat X-ray spectrum for a NLS1. This sample simply demonstrates that Mrk~507 is not an isolated case, and in fact a small but interesting subset of NLS1s do not appear to fit within the PC1 framework.  The apparent lack of strong absorption in some of these flat sources indicates that their X-ray spectra actually are intrinsically flat. Further studies of X-ray weak NLS1s, as well as much larger samples from surveys such as the SDSS, should offer greater insight into the mechanism(s) behind PC1.  Due to the short exposure times of these observations ($\\la 2$~ksec), we cannot infer much outside of $\\Gamma$ and luminosity estimates for individual objects; indeed, this program was intended to study the group properties of an X-ray weak NLS1 sample.  However, there are several objects in this sample with exceptionally low $\\Gamma$ which may be worthy of further study. These hard X-ray NLS1s may represent a new, rare subclass which are optically normal but highly absorbed in the X-rays, or which exhibit abnormally low $L_{\\rm bol}/L_{\\rm Edd}$, or both."
    },
    "0402/astro-ph0402146_arXiv.txt": {
        "abstract": "We first present a summary of our numerical work on accretion discs in close binary systems. Our recent studies on numerical simulations of the surface flow on the mass-losing star in a close binary star is then reviewed. ",
        "introduction": "An accretion disc around a compact star in a close binary star system is an ubiquitous and essential object. Accretion discs play, for example, important roles in cataclysmic variables, nova and X-ray sources. The standard theory to explain the physics in accretion discs is the $\\alpha$-disc model proposed by Shakura and Sunyaev (1973). In this theory, the accretion disc is in some kind of turbulent state, in which turbulent viscosity is parameterized by a phenomenological parameter $\\alpha$.  However, the $\\alpha$-disc model is rather crude approximation to an accretion disc in a close binary system: tidal effects due to the companion star are, for example, not taken into account. To better take these effects into account, one has to rely on numerical simulations.  \\bigskip  \\subsection{Numerical simulations of gas flow in a close binary system}  A pioneering numerical study of accretion discs in close binary systems was started by Prendergast (1960). At that time both computers and computational fluid dynamics were not well developed, so his work was only a preliminary one. It should be noted that Prendergast also started a pioneering work on barred galaxies at that time.  Prendergast \\& Taam (1974) made a first reliable calculation of gas flow in a close binary system using the beam scheme developed by Prendergast. The beam scheme can be considered as a forerunner of the lattice Boltzmann scheme. In order to solve for the fluid flow, the scheme uses the Boltzmann equation rather than the Euler equation as a basic equation. The velocity distribution function has values only at fixed points in velocity space. The original scheme had the drawback of too much artificial viscosity.  At that time, Sorensen Matsuda \\& Sakurai (1974, 1975) were working on a numerical study of gas flow in a close binary system, and they were surprised to find the paper by Prendergast and Taam. However, since the size of their mass-accreting star was very large, their model corresponded to the maybe less interesting Algol-type binaries rather than to cataclysmic variables or X-ray stars.  Sorensen et al. (1974, 1976) adopted a much smaller size of the mass-accreting star to simulate a compact star, although the size was still much larger than that of a realistic compact star, i.e. a white dwarf, a neutron star or a black hole. If the numerical size of the compact star is smaller than the so-called circularization radius, an accretion disc may be formed.  Sorensen et al. used the Fluid in Cell Method (FLIC) with first order accuracy, and computed the flow only in the orbital plane, using a Cartesian grid. Figure 1 shows the density distribution and velocity vectors of a Roche-lobe over-flow in a semi-detached binary system with a mass ratio of one. Gas flows out from a mass-losing star (left) through the L1 point towards a mass-accreting compact star (right). The L1 stream, similar to an elephant trunk is visible but the accretion disc is not well resolved.  Lin \\& Pringle (1976) investigated a similar problem using the sticky particle method, which utilizes both particles and cells. Particles entering a cell are assumed to collide, and velocities of particles after the collision are calculated assuming the conservation of momentum and angular momentum. This method may be thought as a forerunner of SPH scheme, which is a particle scheme frequently used in the astrophysics community, but it would be more appropriate to consider it as a forerunner of the Direct Simulation Monte Carlo method (DSMC) developed later by Bird (see Bird, 1994). In DSMC, the number of particles in a cell is generally much larger, typically 10-100, and collision pairs are selected randomly based on collision probability. The present authors investigated applications of DSMC to astrophysics.  \\begin{figure} \\begin{center} \\psfig{figure=Sorensen.eps,scale=0.35} \\caption{Two-dimensional hydrodynamic simulation of accretion disc using FLIC method: Density distribution and velocity vectors on the rotational plane are shown. The left oval shape is a mass-losing companion star, while the dot at the right shows the position of a mass-accreting compact object. Gas from the companion star flows through the L1 point towards the compact star due to the gravitational attraction to form a so-called elephant trunk (after Sorensen et al. 1975).} \\label{Sorensen} \\end{center} \\end{figure}  \\bigskip  \\subsection{Modern calculation of accretion flow}  Sawada et al. (1986, 1987) investigated again two-dimensional calculations of accretion discs using the Osher upwind scheme and Fujitsu VP200/400 vector supercomputers. Figure 2 shows the density distribution in the orbital plane in a semi-detached binary system with unit mass ratio. They first made their calculations using a first-order scheme. When they switched to a second-order scheme, they discovered a pair of spiral shaped shock waves, as seen in the figure. It is very suggestive that using higher order scheme reveals a new feature which could not be seen in a scheme with lower accuracy.  Spiral shocks in an accretion disc may represent an interesting possibility to solve a long-standing mystery in the theory of accretion discs, i.e. the problem of angular momentum transfer. In order for accretion to occur, the gas in the accretion disc has to lose its angular momentum. In conventional standard disc model, the disc is supposed to be in a turbulent state and the transfer of angular momentum is supposed to occur through the turbulent viscosity. However, in spite of many efforts to show the disc to be unstable, there has been no success.  In the spiral shock model, gas loses angular momentum at the shocks. Nevertheless, the spiral shock model had not attracted much attention from researchers, and there was even an opinion that spiral shocks did not exist in three-dimensional calculations. Sawada \\& Matsuda (1992) performed the first three-dimensional hydrodynamic calculation and obtained spiral shocks. Figure 3 shows our recent calculation by Fujiwara et al. (2001). The figure shows an iso-density surface of an accretion disc around a compact object. Flow-lines on the iso-density surface and on the orbital plane are visualized by the LIC method.  \\begin{figure} \\begin{center} \\psfig{figure=Sawada.eps,scale=0.5} \\caption{Calculation based on the second-order Osher scheme: Density distribution on the rotational plane is shown. A circle at the center represents a mass-accreting compact object. Gas from the mass-losing companion (at left) flows through the L1 point and forms an accretion disc. A pair of spiral shock in the accretion disc can be seen (after Sawada et al. 1986, 1987).} \\label{Figure 2} \\end{center} \\end{figure}   \\begin{figure} \\begin{center} \\psfig{figure=Nagae.eps,scale=0.3} \\caption{Recent three-dimensional calculation: Iso-density surface and flow-lines on the surface/rotational-plane are shown. Three-dimensional structure of spiral shocks is evident. It is remarkable that the flow from the L1 point penetrates into the disc (after Matsuda et al. 2000).} \\end{center} \\end{figure}  \\subsection{Discovery of spiral shocks by observation}  As was pointed out earlier, the spiral shock model may solve the long-standing angular momentum problem. Even if not so, if spiral shocks are present, they must have some observational implications. In 1997, they were apparently detected by Steeghs, Harlaftis \\& Horne (1997) in the cataclysmic variable, IP Pegasi, using the Doppler tomography technique.  Tomography is a technique used, for example, to visualize a cross section of the human body by measuring the absorption of irradiated X-rays. In Doppler tomography, emission lines of hydrogen or helium emitted from hot gas circulating around a compact star are observed and analyzed to give a Doppler map. In X-ray tomography, it is the illuminator that rotates about a human, but in the case of Doppler tomography, use is made of the rotation of the binary system. From temporal variation of the spectrum, one can construct a Doppler map, which is a distribution of emission in the velocity space. From a Doppler map only, it is however not possible to construct uniquely the density distribution in the configuration space. Nevertheless, we may draw useful information from Doppler maps. For example, if spiral structure of hot region emitting spectrum lines exists, it is reflected as a spiral structure in the Doppler map. Steeghs et al. found such a structure.  The ring-like structure observed in a Doppler map represents an accretion disc. If the disc is axi-symmetric around a compact star, as is assumed in the standard disc model, the emission structure should be also axi-symmetric. However, the emission structure shows a spiral feature.  Interestingly, the surface of the mass-losing companion star is also bright. This is because the surface of the companion is irradiated by a radiation from the hot central part of the accretion disc. Moreover this bright region on the companion start is shifted slightly from a symmetry axis. It may be due to a current on the surface of the companion star. ",
        "conclusions": ""
    },
    "0402/hep-th0402021_arXiv.txt": {
        "abstract": "This is a short overview of spatially flat (or open) four-dimensional accelerating cosmologies for some simple exponential potentials obtained by string or M theory compactification on some non-trivial curved spaces, which may lead to some striking results, e.g., the observed cosmic acceleration and the scale of the dark energy from first principles. ",
        "introduction": "The universe appears to be accelerating, but the reason why is yet to be explored. Cosmologists are trying to confirm the hypothesis of a dark universe, providing two alternative candidates for the dark energy -- a positive cosmological constant and a variable $\\Lambda$-term, like a slowly varying scalar potential; both exert a gravitationally self-repulsive force. One would hold out less hope of understanding the dark energy (or cosmic acceleration) unless or until there is a unified theory with a compelling particle-physics motivation that takes it closer to the bedrock of space and time.  The usual idea in superstring or M theory is that spacetime is a four dimensional nearly flat metric times a small six or seven dimensional internal manifold. The current interest in cosmology with the extra (spatial) dimensions is two fold : \\\\ (i) can inflation and/or observed cosmic acceleration naturally arise from string or M theory compactification on some non-trivial curved internal spaces of time-varying volume~\\cite{Townsend03a,Ohta03a,MNR1,IPN03b,IPN03c,IPN03d,Mohaupt}, explaining the scale of the dark energy~\\cite{GKL03a,IPN03e}.\\\\ (ii) can one derive a scalar potential from string theory compactification, with static and warped extra dimensions, that has at least one stationary point with $V> 0$~\\cite{KKLT}.  One of the obstacles for a warped de Sitter type compactification is a no-go theorem~\\cite{nogo1}. The strong energy condition holds for all $D=10$ or $11$ supergravities. If one takes the extra dimensions to be warped (and static), then for the compactified theory (with or without form-fields in the extra dimensions), one finds $R_{00}^{(4)} \\geq 0$, which does not allow the universe to accelerate. It seems less likely that a (false) de Sitter state one may obtain in warped string theory background, with static extra dimensions, is responsible for cosmic acceleration, because such a vacuum is too unstable for a significant period of inflation to occur. In certain cases, it may be possible to have late time accelerated expansion from a tachyonic potential~\\cite{padma02a}.  The models like $V\\sim \\lambda \\varphi^4$, $\\sim m^2\\varphi^2$, or their hybrids, do not deal with the basic puzzles such as that of the initial singularity, nor with the mystery behind dark energy indicated by recent observation. So it might be worthwhile to carefully consider the string or M theory compactification in time dependent backgrounds. I shall argue that the dark energy of the universe is possibly a gravitational scalar potential arising from slowly varying size of extra dimensions, and so it is dynamical. ",
        "conclusions": ""
    },
    "0402/astro-ph0402620_arXiv.txt": {
        "abstract": "We build on recent new evolutionary models of Jupiter and Saturn and here extend our calculations to investigate the evolution of extrasolar giant planets of mass 0.15 to 3.0 \\mj.  Our inhomogeneous thermal history models show that the possible phase separation of helium from liquid metallic hydrogen in the deep interiors of these planets can lead to luminosities $\\sim$~2 times greater than have been predicted by homogeneous models.  For our chosen phase diagram this phase separation will begin to affect the planets' evolution at $\\sim$~700 Myr for a 0.15 \\mj\\ object and $\\sim$~10 Gyr for a 3.0 \\mj\\ object.  We show how phase separation affects the luminosity, effective temperature, radii, and atmospheric helium mass fraction as a function of age for planets of various masses, with and without heavy element cores, and with and without the effect of modest stellar irradiation.  This phase separation process will likely not affect giant planets within a few AU of their parent star, as these planets will cool to their equilibrium temperatures, determined by stellar heating, before the onset of phase separation.  We discuss the detectability of these objects and the likelihood that the energy provided by helium phase separation can change the timescales for formation and settling of ammonia clouds by several Gyr.  We discuss how correctly incorporating stellar irradiation into giant planet atmosphere and albedo modeling may lead to a consistent evolutionary history for Jupiter and Saturn. ",
        "introduction": "Over the past 8 years, nearly 120 giant planets have been found in orbit around other stars.  These planets have added immensely to the regions of parameter space in which we find giant planets and we are just beginning to understand how these interesting (and often hot!) environments affect the evolution of these objects. (See \\citet{Hubbard02}, for a review.)  However, as we strive to understand these strange new worlds we must remind ourselves that our understanding of our closest two gas giants, Jupiter and Saturn, is far from complete.  As our examples of giant planets that will always be the most amenable for detailed study, it is of great importance to refine our understanding of these planets, and the physics that governs them, so that we will have confidence in our understanding of more distant giant planets.  As radial velocity studies reach longer time baselines, we are assured of finding planets similar to Jupiter and Saturn, and at similar orbital distances.  These wider orbital separation planets are also more likely to be directly imaged because they are farther from the glare of their parent stars.  This paper focuses on applying recent advances in our understanding of the evolution of Jupiter and Saturn to hypothetical extrasolar giant planets of various masses and orbital distances.  Our understanding of the evolution of Jupiter and Saturn is currently imperfect.  The most striking discrepancy between theory and reality is Saturn's luminosity. Saturn's current luminosity is over 50\\% greater than one predicts using a homogeneous evolution model, with the internally isentropic planet radiating over time both its internal energy and thermalized solar radiation.  This discrepancy has long been noted \\citep{Pollack77, Grossman80, Guillot95, Hubbard99}.  Homogeneous evolutionary models of Saturn tend to reach an effective temperature of 95.0 K (Saturn's current known $T_{eff}$) in only $\\sim$~2.0 to 2.7 Gyr, depending on the hydrogen-helium equation of state (EOS) and atmosphere models used.  However, purely homogeneous models appear to work well for Jupiter.  \\mbox{Figure~\\ref{figure:js}} shows homogeneous evolutionary models for both planets from \\citet{FH03} (hereafter \\pf).  It has also long been believed that the most promising route to resolving this discrepancy is the possible phase separation of neutral helium from liquid metallic hydrogen in the planet's interior, beginning when Saturn's effective temperature reached $\\sim$100 - 120 K \\citep{SS77a,SS77b}.  Immiscibilities in two-component systems are common, and are the byproduct of the interaction potentials of the types of atoms (or molecules) in the mixture.  Once the temperature of a system becomes low enough, the energy of mixing becomes small enough that the Gibbs free energy of the system can be minimized if the system separates into two distinct phases.   One phase contains slightly less solute (here, helium) in the solvent (here, liquid metallic hydrogen) than there was initially, and the other phase nearly pure solute.  This is often termed immiscibility, insolubility, or phase separation, and the mixture is said to have a miscibility gap or solubility gap.  In general, the smaller the percentage of atoms in a mixture that is solute, the lower the temperature the mixture must attain for immiscibility to occur.  A common approximation \\citep[see][]{Stevenson79, Pfaff} for the saturation value of $x$, the number fraction of the solute, is \\begin{equation} \\label{x} x=exp(B - A/k_bT), \\end{equation} where $B$ is a dimensionless constant, $k_B$ is Boltzmann's constant, $T$ is temperature, and $A$ is a positive, pressure dependent constant with units of energy.  As described in \\pf\\ $B$ should be close to zero, and $A$ is the increase in free energy upon addition of a helium (or whatever atom in general) to pure liquid metallic hydrogen.  In \\pf\\ we showed that whether $A$ increases or decreases with pressure has important effects on giant planet evolutionary models.  In a solar composition mixture $x$ for helium is about 0.085.  Oxygen, the 2nd most abundant element, is down by a factor of over 100.  Since helium is relatively abundant in the hydrogen mixture, the helium (which is predicted to be neutral) will perturb the structure of the proton-electron plasma.  This $A$ constant has been calculated by various methods in the papers we will mention below to be $\\sim$~1-2 eV, which for solar composition leads to temperatures on the order of 5000-10000 K for the onset of helium immiscibility.  (This is also dependent on pressure.)  \\mbox{Figure~\\ref{figure:phased}} shows in detail our current knowledge of the high pressure phase diagram of hydrogen and helium \\citep{Hubbard02}.  Labeled are the current interior adiabats of Jupiter, Saturn, and a hypothetical 0.15 \\mj\\ planet.  Relevant experimental and theoretical boundaries are also labeled, as are regions of calculated helium immiscibility.  \\citet{Salpeter73} was the first to note the effects of the immiscibility of helium in liquid metallic hydrogen on the structure and evolution of a hot, adiabatic, hydrogen-helium planet.  \\citet{Stevenson75}, using perturbation theory, performed the first detailed calculation of the hydrogen-helium phase diagram, in an effort to map the regions of pressure-temperature-composition space in which helium was likely to become immiscible.  His calculations roughly agreed with estimates of the current pressures and temperatures of liquid metallic hydrogen in Jupiter and Saturn's interiors.  These calculations indicated that as pressure increased, the saturation concentration of helium in liquid metallic hydrogen would increase, leading to constant composition curves that slant down and to the right in \\mbox{Figure~\\ref{figure:phased}}.  Soon after, \\citet{SS77a, SS77b} performed detailed calculations on the dynamics and distribution of helium in giant planets.  They found that when helium becomes immiscible in liquid metallic hydrogen, the composition that separates out is essentially pure helium, and this helium on fairly short timescales (relative to the convective timescale) will coalesce to form helium droplets.  These droplets, once they reach a size of $\\sim$~1 cm, will attain a Stokes velocity greater than the convective velocity and will then fall down through the planet's gravitational field.  If the droplets reach a region where helium is again miscible at higher concentration, they will redissolve, enriching the deeper regions of the planet in helium.  They found that this ``helium rain'' could be a substantial additional energy source for giant planets.  Helium would be lost from $all$ regions with pressures lower than the pressures in the immiscibility region, since the planet is fully convective (or nearly so) up to the visible atmosphere.  Excess helium would be mixed down to the immiscibility region and be lost to deeper layers.  This would leave all molecular regions up to the visible atmosphere depleted in helium.  Relatively few studies have been done since then on phase diagrams of hydrogen-helium mixtures.  \\citet{HDW}, using a Monte Carlo technique, but similar assumptions to that of \\citet{Stevenson75}, obtained essentially the same results.  The most recent calculations \\citep{Pfaff} utilized molecular dynamics to predict a helium-immiscibility region with a shape very different from that of \\citet{Stevenson75} and \\citet{HDW}.  They find that as pressure increases, the saturation concentration of helium in liquid metallic hydrogen decreases, which leads to constant composition lines that slant up and to the right in \\mbox{Figure~\\ref{figure:phased}}.  Interestingly, these constant-composition lines run nearly parallel to the giant planet adiabats.  Detailed inhomogeneous evolutionary models including helium phase separation were not performed until \\citet{Hubbard99}.  These authors investigated the cooling of Jupiter and Saturn when the mass of helium rained out linearly with time since Jupiter and Saturn's formation, or alternatively, rained out just before the planets reached their known effective temperatures.  These are two logical limiting cases.  There is observational evidence that helium phase separation has begun in both Jupiter and Saturn.  The protosolar mass fraction of helium is calculated to be near $Y$=0.27 (\\citet{Lodders03} puts the number at 0.2741).  The atmospheres of both Jupiter and Saturn are depleted relative to this value.  With the assumption that these planets globally contain the protosolar $Y$, the missing helium must be in deeper layers of the planet.  The case for Jupiter's depletion is clear cut, with a value of $Y=0.234 \\pm .005$ from the Helium Abundance Detector (HAD) on the Galileo Entry Probe \\citep{vonzahn98}.   The case for Saturn is much less clear.  Without a past or planned Saturn entry probe, Saturn's atmospheric helium abundance can only be obtained through indirect methods, using infrared spectra with or without radio occultation derived temperature-pressure \\emph{(T-P)} profiles.  Analysis of Voyager measurements indicated $Y=0.06 \\pm 0.05$ in Saturn \\citep{Conrath84}.  However, the mismatch between the Voyager derived value for Jupiter ($Y=0.18 \\pm 0.04$, \\citep{Gautier81}) and the accurate HAD measurements, along with \\citet{Hubbard99} evolutionary and \\citet{Guillot99} static models, led \\citet{CG00} to perform a reanalysis of the Voyager data.  The details of their investigation will not be described here, but by disregarding the occultation derived \\emph{T-P} profile, which may be in error, they obtain $Y=0.18-0.25$ for Saturn's atmosphere.  Noting the clear need to better understand Jupiter and Saturn in light of these atmospheric $Y$ values, in \\pf\\ we calculated the first evolutionary models that coupled high-pressure phase diagrams of hydrogen-helium mixtures and a grid of radiative atmosphere models for giant planets.  A variety of Saturn evolutionary models were calculated that included helium phase separation.  The main findings of \\pf\\ were as follows.  The phase diagram of \\citet{HDW}, which is essentially the same as that of \\citet{Stevenson75}, is inapplicable to the interiors of Jupiter and Saturn, if helium phase separation is Saturn's only additional energy source.  These phase diagrams predict that $A$ from equation (\\ref{x}) is a decreasing function of pressure.  As \\mbox{Figure~\\ref{figure:js2}} shows, this phase diagram prolongs Saturn's cooling 0.8 Gyr, even in the most favorable circumstance that all energy liberated is available to be radiated, and does not instead go into heating the planet's deep interior.  \\pf\\ found that if one were to match the $Y_{atmos}$ of \\citet{CG00} and prolong Saturn's evolution to a $T_{eff}$ of 95.0 K at 4.56 Gyr, the helium that becomes immiscible and rains down to deeper layers needs to rain far down into the planet, likely all the way to the core, in order for enough energy to be released and still match the relatively high ($Y_{atmos}$=0.18 - 0.25) helium abundance.  In \\pf\\ an ad-hoc phase diagram was created that was essentially a modification of the phase diagram calculated by \\citet{Pfaff}.  (\\citet{Pfaff} find a phase diagram in which $A$ from equation (\\ref{x}) is an increasing function of pressure.)  To simplify the evolution of the planets the ad-hoc phase diagram was constructed such that helium immiscibility region runs exactly parallel to the planets' internal adiabats.  Therefore, there is no region for the helium droplets to redissolve in the liquid metallic hydrogen.  This causes helium that phase separates to rain all the way down to the planet's heavy element core.  Consequently, all hydrogen/helium regions, molecular and metallic, become more helium poor as the helium layer on top of the core grows.  This ad-hoc phase diagram allows Saturn to reach an age of 4.56 Gyr and $T_{eff}$ 95 K while its $Y_{atmos}$ drops to 0.185.  \\mbox{Figure~\\ref{figure:js2}} shows the evolution of $T_{eff}$ vs.~time.  With this phase diagram, Jupiter evolves homogeneously to the present day and reaches $\\sim$~4.7 Gyr at 124.4 K without helium becoming immiscible.  Jupiter would then begin to evolve inhomogeneously at $T_{eff}$ below 123 K.  Still in need of explanation is Jupiter's depletion of helium (and neon, which may have been carried away in the helium \\citep{Roulston95}) in its atmosphere.  What we have from \\pf\\ is a high-pressure hydrogen-helium phase diagram that is calibrated to Jupiter and Saturn.  Specifically, both planets reach their known effective temperatures after $\\sim$~4.6 Gyr, with an improved but still imperfect understanding of helium phase separation.  The purpose of the present paper is to investigate the effects of helium phase separation on the evolution of theoretical giant planets in orbits around other stars using the phase diagram derived in \\pf.  Here we concentrate on planets ranging in mass from 0.15 \\mj\\ (half Saturn's mass) to 3.0 \\mj, and derive the luminosity, $T_{eff}$, radius, and $Y_{atmos}$ as a function of time during the planets' evolution.  Of the many extrasolar planets found to date, the planet with the smallest minimum mass is HD 49674b, at 0.12 \\mj\\ \\citep{Butler03}.  As we will discuss later, a planet with a mass this small likely contains no liquid metallic hydrogen, only dense molecular hydrogen.  Approximately 70\\% of all known planetary candidates have minimum masses of less than 3.0 \\mj, so planets of the masses we explore here are sure to be abundant.  Our standard models (to be discussed in Section \\ref{isolate}) incorporate a primordial 10 \\me\\ heavy element core for all planets, but later in Section \\ref{cores} we investigate the effects of 20 \\me\\ and coreless models.  (1 Jupiter mass, \\mj, is 317.7 Earth masses, \\me.)  We will later find that for the more massive planets, varying the core mass has little effect on $T_{eff}$ but a large effect on planetary radii.  In Section \\ref{irrad} we investigate the effects of modest stellar irradiation.  Stellar heating retards a planet's cooling, and we find that if a planet is within a few AU of its parent star, the planet will reach its equilibrium temperature before its interior reaches temperatures cool enough for helium to become immiscible.  We calculate the cooling of planets in isolation and at 10 and 5 AU from a constant luminosity 1.0 L$_{\\odot}$ star.  In Section \\ref{discuss} we discuss the atmospheric properties and detectability of these EGPs likely to be undergoing helium phase separation. ",
        "conclusions": "Using the phase diagram described in \\pf, which is calibrated to Jupiter and Saturn and allows both planets to reach their known ages and effective temperatures, we have explored the effects that helium phase separation will have on a variety of EGPs.  The additional energy liberated as helium rains to deeper layers of a planet will significantly delay the cooling and contraction of giant planets.  Once helium phase separation is underway, our inhomogeneous evolutionary models predict luminosities $\\sim$~2 times greater than predictions from homogeneous models.  This will make these giant planets in the 0.15 to 3.0 \\mj\\ mass range somewhat easier to detect that has been previously thought.  Improvements in understanding the evolution of these EGPs will come through a better understanding of our closest giant planets, Jupiter and Saturn."
    },
    "0402/astro-ph0402550_arXiv.txt": {
        "abstract": "We examine theoretically the behaviour of the inner accretion disk in GRS~1915+105 when soft X-ray dips are present in the X-ray light curve. We assume the presence of a radial shock in the accretion disk, as in some of the Two Component Advective Flow (TCAF) solutions. We discuss the behaviour of the flux tubes inside a TCAF (which we name Magnetized TCAF or MTCAF model for brevity) and compare various competing forces on the flux tubes. In this MTCAF model, we find that the magnetic tension is the strongest force in a hot plasma of temperature $\\gsim 10^{10}$K and as a result, magnetic flux tubes entering in this region collapse catastrophically, thereby occasionally evacuating the inner disk. We postulate that this magnetic `rubber-band' effect induced evacuated disk matter produces the blobby components of outflows and IR/radio jets. We derive the size of the post-shock region by equating the time scale of the Quasi-Periodic Oscillations to the infall time of accreting matter in the post-shock region and found the shock location to be $\\sim 45-66 r_g$. We calculate the transition radius $r_{tr}$, where the Keplerian disk deviates into a sub-Keplerian flow, to be $\\sim 320r_g$. Based on the derived X-ray spectral parameters, we calculate the mass of this region to be $\\sim$10$^{18}$g. We conclude that during the X-ray dips the matter in the post-shock region, which manifests itself as the thermal-Compton component in the X-ray spectrum, is ejected, along with some sub-Keplerian matter in the pre-shock region. ",
        "introduction": "GRS~1915+105 has proven to be an ideal source to study in detail many of the physical concepts regarding accretion onto black holes. Ever since its discovery (Castro-Tirado et al. 1992), it has been continuously bright in X-rays and it shows a variety of X-ray variability characteristics (Morgan, Remillard and  Greiner 1997; Muno et al. 1999; Yadav et al. 1999; Belloni et al. 2000). It has been monitored extensively in the radio band (Mirabel and Rodriguez 1994; Pooley and Fender 1997; Fender et al. 1999) and several episodes of high radio emissions, huge flares associated with superluminal motions, radio oscillations etc. are observed in this source. Several attempts have been made to associate the radio emission, presumably coming from jets, to the X-ray emission from the accretion disks (Fender et al. 1999; Naik et al. 2000; Naik and Rao 2000). The Spectral signature of winds from the Comptonising region has also been identified (Chakrabarti et al. 2001).  Chakrabarti and Manickam (2000, hereafter CM00) have applied the Two Component Advective Flow (TCAF) model of Chakrabarti and  Titarchuk (1995) to explain various X-ray variability characteristics of GRS~1915+105. Recently there has been overwhelming evidence that the TCAF model is valid for many black hole candidates (Smith, Heindl and Swank, 2001; Smith et al. 2001). CM00 invoked outflows from the inner accretion disk to explain a correlation between the QPO frequency and the duration of the burst-off states during a regular oscillations seen in the source. These outflows, however, are confined to the sonic sphere and fall back on the accretion disk after being cooled down by an inverse Compton effect. It was pointed out by Naik and Rao (2000) that the source does not show appreciable radio emission during such oscillations. A detailed inflow/outflow model has not been presented for this source to explain the radio emission, particularly to explain the superluminally moving radio blobs.  Recently Naik et al. (2000) have detected a series of soft X-ray dips during the declining phase of a huge radio flare and have postulated that such soft dips are responsible for the jet emission. Vadawale et al. (2001) made a detailed study of X-ray dips observed during the radio flare using the Rossi X-ray Timing Experiment (RXTE) data and have presented evidence for the disappearance of the inner accretion disk during the dips. Since the disappearance of the inner disk is seen to be correlated with intense radio activity, the role of the magnetic field must be studied in order to understand the system completely. Rodriguez and Mirabel (1999) estimated the field in radio blobs to be around tens of mG at 500-1000AU (in 1994 observation). Fender et al. (1997) requires the field to be around 8G at around 1AU (in their 1996 observations). From the similarity of $\\sim 30$min oscillations in IR and Radio, they concluded that the radio blobs are adiabatically expanding and are independently ejected from the disk every $30-40$ minutes. If the trapped field inside a radio blob is of roughly $1/r$ (for toroidal field) then its interpolated value close to a black hole is around $10^{7}$G at around $10r_g$ which is comparable to an equipartition  value. Thus, one needs to correlate fields ejected from the disk with those observed inside the radio blobs.  In this Paper, we examine the mass ejection based on the TCAF model in presence of a magnetic field (we call this as Magnetized TCAF or MTCAF model) amplified due to strong shear at the transition radius of the Keplerian and a sub-Keplerian flow. In the next Section, we discuss forces which govern the motions of the flux tubes and write equations of toroidal flux tubes inside an accretion disk with a constant angular momentum. We show that close to the black hole, where the flow could be very hot ($\\gsim 10^{10}$K)the flux tubes move at least with the Alf\\'ven speed and may catastrophically collapse like a stretched rubber band. We conjecture that such rapid collapse would assist evacuation of matter from the disk and cause X-ray `dips' seen in the light curves.  In \\S3 we estimate the mass of the ejecta which agrees with observations. Finally, in \\S4 we draw our conclusions. ",
        "conclusions": "In this Paper, we have given a physical basis for a sudden mass ejection in GRS~1915+105. We showed that if matter brings in a particularly strong magnetic field, this would be sheared and amplified to a value much above the equipartition value before it can be expelled by buoyancy. Magnetic tension collapses these toroidal flux tubes at a highly supersonic speed, much faster than the flow velocity. This has the effect of displacing matter from the disk in transverse direction (much like a fast boat causing spillage on a shore) and depositing it to outflowing winds. From the observed fits of Vadawale et al. (2001) we estimated the electron number density and the mass of the post-shock region and the sub-Keplerian flow to be around $10^{18}$g and $10^{20}$g respectively. Our estimate of the post-shock mass is a factor of ten less than the mass estimate ($10^{19}$g) of `baby-jets' (Mirabel et al 1998) associated with IR and radio flares and could therefore be due to ejection of some sub-Keplerian matter as well. These `baby-jets' are found to be associated with class $\\beta$ light curves which have soft X-ray dips. These dips are also seen in class $\\theta$ light curves. During a major portion of the huge radio flares associated with superluminal blob emission a series of soft dips are present (Naik et al. 2000). Mirabel and Rodriguez (1999) have pointed out that in each epoch of this type of outflow, the mass condensation is around $10^{23}$g. In order to achieve this, we require that matter is accumulated from disk evacuation at each `dip' and within each epoch, successive mini-blobs move faster than the earlier blob in order to `catch up'. This may indicate some other runaway process with a longer time scale of tens of days. Naik et al (2000) have observed such X-ray dips at a rate of once in a few hundred seconds during the peak or the radio flare. If there are $\\sim 1000$ evacuation events during an episode of superluminal blob ejection (in a few days), then total mass condensation would be $10^{23}$g. Hence in order to explain the observation of Mirabel and Rodriguez (1999), one must require that in each epoch, matter is accumulated from at least a thousand evacuation events. Future observation would tell if such is the case.  It is to be noted that the mass of the condensation as estimated by Mirabel and Rodriguez (1999) is based on the presence of one electron per proton, i.e., no pair production is assumed. With a pair density, say, ten times larger, the estimated mass would be ten times less. However, at the same time, estimated mass of the disk would also be reduced by the same factor. Hence, the number of ejection events is not affected."
    },
    "0402/astro-ph0402599_arXiv.txt": {
        "abstract": "{We present a study of the stellar parameters, distances and spectral energy distributions (SEDs) of \\object{HD 34282} and \\object{HD 141569}, two pre-main sequence Herbig AeBe stars.  Both objects have been reported to show `anomalous positions' in the HR diagram in the sense that they appear below the main sequence.  A significant result of this work is that both stars are metal-deficient.  The {\\it Hipparcos} distance of HD 34282 is very uncertain and the current study places the star at the expected evolutionary position in the HR diagram, i.e. as a PMS star. The distance for HD 141569 found in this work matches the {\\it Hipparcos} distance, and the problem of its anomalous position is solved as a result of the low metallicity of the object: using the right metallicity tracks, the star is in the PMS region.  The SEDs are constructed using data covering ultraviolet to millimetre wavelengths. Physical, non-parametric models, have been applied in order to extract some properties of the disks surrounding the stars. The disk around HD 34282 is accreting actively, it is massive and presents large grains in the mid-plane and small grains in the surface. HD 141569 has a very low mass disk, which is in an intermediate stage towards a debris-type disk. ",
        "introduction": "\\label{THEINTRO}  The study of protoplanetary disks is currently undergoing an exciting stage partly propelled by the discovery of extrasolar planetary systems following the detection of 51 Peg B by Mayor \\& Queloz (1995). The possible discovery of telluric planets in the near future, in addition to the jovian-like planets already detected, will pose interesting questions on the formation of extrasolar planetary systems. Knowledge of the properties of protoplanetary disks and how they evolve to debris disks around main-sequence (MS) stars is one of the tools required for modelling such a process.  Observed spectral energy distributions (SEDs) of pre-main sequence (PMS) stars are widely used to study the properties of protoplanetary disks and to classify T Tauri and Herbig AeBe stars (HAeBe hereafter) into an evolutionary scheme (Adams et al. 1987; Hillenbrand et al. 1992). Although the interpretation of a given SED is, to some extent, model dependent, the theoretical modelling of SEDs constitutes an invaluable tool for understanding the structure and properties of protoplanetary disks, e.g.  Chiang \\& Goldreich (1997, 1999), D'Alessio et al. (1998, 1999, 2001), Dullemond et al.  (2001).  The EXPORT consortium (Eiroa et al. 2000) observed a large sample of PMS and Vega-type stars during the 1998 International Time Programme of the Canary Islands' Observatories. One of the driving goals of this effort was the study of the evolution and properties of protoplanetary disks by analysing the SEDs of the young stars, taking advantage of the fact that the EXPORT optical and near-IR photometry were obtained simultaneously.  This observational approach is appropriate since T Tauri and HAeBe stars vary markedly in these spectral regimes and a significant part of the total luminosity of the object is radiated by the PMS stellar photosphere at these wavelengths.  Among the stars in the EXPORT sample with measured {\\it Hipparcos} parallaxes, the HAeBe stars HD 34282 and HD 141569 are the only ones whose positions fall below the zero-age main sequence in a $\\log L - \\log T_{\\rm eff}$ HR diagram or equivalent (e.g.  van den Ancker et al. 1998, Weinberger et al. 2000).  This result is difficult to reconcile with some observational results and with theoretical PMS evolutionary models (e.g.  Yi et al. 2001).  In this paper we present an analysis of the stellar properties of HD 34282 and HD 141569 and their circumstellar disks, based on EXPORT data complemented with new spectroscopic observations and data from the literature.  The structure of the paper is as follows. In Section \\ref{THESTARS} we briefly review what is currently known about the stars. Sections \\ref{THEOBS} and \\ref{THERESULTS} present the observations used in this work and their results. In Section \\ref{THEPARAMS} we do a comprehensive study of the stellar parameters and distances to the stars. In Section \\ref{THESEDS} the disk models used to reproduce the observed SEDs are presented. In Section \\ref{THECONC} we summarize the results of the work. ",
        "conclusions": "\\label{THECONC}  We have presented a study of the properties and spectral energy distributions of the HAeBe stars HD 34282 and HD 141569 and their disks, based mainly on observations made by the EXPORT consortium (Eiroa et al. 2000). The main conclusions can be summarized as follows:  \\begin{itemize}  \\item[--]The spectroscopic and photometric observations of both stars have been described and analysed in detail with the aim of throwing light on their evolutionary status and the dynamics and variability of their photospheres and disks.  \\item[--]Both stars are metal-deficient. Our estimates are [Fe/H]=$-0.8$ for HD 34282 and $-0.5$ for HD 141569.  \\item[--]The analysis has provided values for the stellar distances. The distance to HD 34282, namely 348 pc, is consistent with the value given by Pi\\'etu et al. (2003) and corrects the {\\it Hipparcos} distance that is uncertain due to the large error in the parallax. The new distance for HD 34282, together with its low metallicity, resolves the anomalous position of this star in the HR diagram previously reported under the assumption that the star has solar abundances. The distance to HD 141569, namely 108 pc, is consistent with the {\\it Hipparcos} data and in this case, the question of the anomalous position of this star in the HR diagram is resolved by using the correct set of tracks with the appropriate metallicity.  \\item[--]In addition to the above two parameters, our analysis has provided for both stars values for the effective temperatures, spectral types, masses, gravities, luminosities, ages, projected rotational velocities and extinctions.  These results are given in Table \\ref{STARS}.  \\item[--]Complete SEDs from the ultraviolet to the millimetre range have been constructed for both stars and their disks using the EXPORT data plus additional results gathered from the literature and on-line databases. Table \\ref{TableSED} gives a full account of this compilation and the data sources.  \\item[--]The self-consistent irradiated accretion disk models of D'Alessio et al. (1998, 1999 and 2001) have been used to fit the SEDs of both systems. A summary of the disk parameters obtained is given in Table \\ref{DISKS}.  \\item[--]The SED of the disk of HD 34282 has been reproduced using a three-component model. The {\\it mid-plane disk} is modelled with a maximum dust grain size of 1 cm and emits mainly at submillimetre and millimetre wavelengths, the {\\it surface disk} model has a maximum grain size of 1 $\\mu$m and is responsible for the mid-IR excess emission. A near-IR bump similar to those seen in other HAeBe stars (see Natta et al. 2001) is fitted with the emission of a frontally illuminated {\\it wall} placed at the dust destruction radius (Dullemond et al.  2001; Muzerolle et al. 2003b, D'Alessio et al. 2003). The geometry of the wall is constrained by the near-IR photometry.  \\item[--]In the process of fitting the SED of HD 34282, we have tried for the first time a new method to estimate the mass accretion rate towards a HAeBe star using ultraviolet Walraven photometry and the magnetospheric accretion shock models of Muzerolle et al. (2003a). We obtained un upper limit for the mass accretion rate in HD 34282 compatible with the parameters needed to fit the millimetre SED. The low value found does not support the idea by Herbst \\& Shevchenko (1999) about the FU Ori-like mechanism causing the photometric activity of UX Ori stars.  \\item[--]The SED of HD 141569 has been fitted with a disk model irradiated by the central star with negligible accretion and with large grains ($a_{\\rm max}=1$ mm). A slightly flared dust and gas distribution was obtained, which is consistent with a substantial amount of residual gas in the disk.  \\end{itemize}"
    },
    "0402/astro-ph0402507.txt": {
        "abstract": "We study the rotational distortions of the vacuum dipole magnetic field in the context of geometrical models of the radio emission from pulsars. We find that at low altitudes the rotation deflects the local direction of the magnetic field by at most an angle of the order of $\\rn^2$, where $\\rn = r/\\rlc$, $r$ is the radial distance of the radio emission and $\\rlc$ is the light cylinder radius. To the lowest (ie.~second) order in $\\rn$, this distortion is symmetrical with respect to the plane containing the dipole axis and the rotation axis ($\\om$ plane). The lowest order distortion which is asymmetrical with respect to the $\\om$ plane is third order in $\\rn$. These results confirm the common assumption that the rotational sweepback has negligible effect on the position angle (PA) curve. We show, however, that the influence of the sweepback on the outer boundary of the open field line region (open volume) is a much larger effect, of the order of $\\rn^{1/2}$. The open volume is shifted backwards with respect to the rotation direction by an angle $\\delta_{\\rm ov} \\sim 0.2\\sin\\alpha \\rn^{1/2}$ where $\\alpha$ is the dipole inclination with respect to the rotation axis. The associated phase shift of the pulse profile $\\Delta\\phi_{\\rm ov} \\sim 0.2\\rn^{1/2}$ can easily exceed the shift due to combined effects of aberration and propagation time delays ($\\approx 2\\rn$). This strongly affects the misalignment of the center of the PA curve and the center of the pulse profile, thereby modifying the delay-radius relation. Contrary to intuition, the effect of sweepback dominates over other effects when emission occurs at low altitudes. For $\\rn \\la 3\\cdot 10^{-3}$ the shift becomes negative, ie.~the center of the position angle curve precedes the profile center. With the sweepback effect included, the modified delay-radius relation predicts larger emission radii and is in much better agreement with the other methods of determining $\\rn$. ",
        "introduction": "There are independent observational arguments which imply that the pulsar radio emission occurs in a form of a narrow beam centered (or roughly centered) on the magnetic dipole axis. In many cases the position angle (PA) of the observed linearly polarized radiation changes its direction by nearly $180^\\circ$ when our line of sight crosses the radio beam (eg.~Lyne \\& Manchester 1988). If associated with the direction of $\\vec B$, this change of PA can be naturally interpreted as a result of our line of sight passing near the magnetic pole (Radhakrishnan \\& Cooke 1969). Moreover, the width $\\rho$ of the radio beam determined for different pulsars from the observed width of their radio pulse profiles scales with the rotation period $P$ as the opening angle of the open field line region, ie.~$\\rho \\propto P^{-1/2}$ (Rankin 1990; Rankin 1993).  It is commonly believed that the emission region associated with the beam does not extend beyond the region of open field lines (hereafter called ``open volume\") which cross the light cylinder of radius $\\rlc = c/\\Omega$ ($c$ is the speed of light and $\\vec \\Omega= \\Omega \\hat z$ is the angular velocity of pulsar rotation). The angular size of the open volume at (small) radial distance $r$ is equal to $\\thov \\simeq (r/\\rlc)^{1/2}$ and the cone formed by tangents to magnetic field lines at the rim of the open volume has angular radius of $\\thb \\simeq 1.5\\thov$. The radial distance of the emission region has not been established so far: both a high-altitude emission region extending over a small fraction of $\\thov$, as well as a low-altitude emission region which fills in a much larger fraction of $\\thov$ may be responsible for the same shape of the radio beam.  The sweepback effect was first investigated in detail by Shitov (1983) who considered it to explain the observed dependence of radio luminosity of pulsars as a function of period. He estimated the magnitude of the rotational distortions of the magnetic field from the torque responsible for the observed slowing down of pulsars. He found that at moderate altitudes within the open volume, ``near\" the dipole axis, the direction of the distorted magnetic field deflects from the direction of the pure (ie.~static shape) dipole barely by an angle \\begin{equation} \\delta_{\\rm sb} \\simeq 1.2\\left(\\frac{r}{\\rlc}\\right)^3 \\sin^2\\alpha, \\label{deflection} \\end{equation} where $\\alpha$ is the dipole inclination with respect to the rotation axis. Gil (1983) proposed the sweepback effect to explain why the separation between the main radio pulse and the interpulse observed in the profile of PSR B0950$+$08 is significantly different from $180^\\circ$.  In 1985 Shitov incorporated the sweepback effect into the model of pulsar position angle curves proposed by Radhakrishnan \\& Cooke (1969) and showed that the sweepback results in a lag of the profile center (measured as the midpoint between the outer edges of the pulse profile) with respect to the center, or the ``inflection point\" of the position angle curve. Shitov emphasized that the lag of the profile center was a sum of \\emph{two} effects: not only the center of the PA curve is shifted toward earlier phases (with respect to the nondistorted case) according to the eq.~(\\ref{deflection}), but also the center of the open volume is displaced backwards, which contributes to the total effect.  In most of subsequent investigations, however, the sweepback has been neglected, mainly on the basis of eq.~(\\ref{deflection}). Blaskiewicz et al.~(1991, hereafter BCW91) proposed a relativistic model of pulsar polarization which took into account two important effects overlooked by Shitov: the presence of the corotational acceleration and the aberration effect. An excellent result of their work was the ``delay-radius\" relation, according to which the center of the PA curve \\emph{lags} the profile center by \\begin{equation} \\Delta\\Phi_{\\sss BCW} \\approx 4\\frac{r}{\\rlc}\\ {\\rm rad,} \\label{bcw} \\end{equation} where $r$ is the radial distance of the radio emission. With no dependence on viewing geometry parameters (like the dipole inclination $\\alpha$, or the viewing angle $\\zobs$ between the rotation axis and the observer's line of sight), their relation appears to provide a powerful method of determining $r$. Equally important, the delay-radius relation depends neither on the observed width of the pulse profile $\\wi$ nor on the separation between the conal components in the pulse profile. Therefore, the altitudes of radio emission provided by eq.~(\\ref{bcw}) may serve to determine which magnetic field lines are associated with the outer edge of the profiles and which field lines correspond to the maxima of conal components (Mitra \\& Rankin 2002; Dyks et al.~2004a). Von Hoensbroech \\& Xilouris (1997) used the delay-radius relation to probe the radius-to-frequency mapping at high radio frequencies. Given that the method is based on a measurement of \\emph{tiny} shifts (of magnitude usually being a small fraction of one degree) between the centers of the PA curve and the profile, it is extremely sensitive to the assumed geometry of the magnetic field. The latter was taken to be a dipole of static shape, with no rotational distortions.  Gangadhara \\& Gupta (2001) proposed another relativistic method of estimating radio emission altitudes for pulsars with both core and conal components. By considering the effects of the aberration and the propagation time delays they showed that the core component lags in phase the midpoint between the maxima of conal components, if the core originates from lower altitudes than the cones, and if the cones are axially symmetric around the core in the reference frame corotating with the star (CF). Dyks et al.~(2004a) revised their method and showed that the phase shift between the core component and the pairs of conal components is equal to \\begin{equation} \\Delta\\phi_{\\sss DRH} \\approx 2\\frac{r}{\\rlc}\\ {\\rm rad} \\label{drh} \\end{equation} which provides another method for determining $r$ without information about viewing geometry (nor $\\wi$). Dyks et al.~(2004a) used the above relation, and the results of work by Gupta \\& Gangadhara (2003) to calculate $r$ for 6 pulsars with well defined core-cone systems. As in the case of the delay-radius relation, the above formula holds only for the magnetic field which is symmetrical about the $\\om$ plane (where $\\vec \\mu$ is the magnetic moment of the pulsar magnetic field), at least as long as one associates the assumed symmetry of the core-cone system with the geometry of the underlying magnetic field. Any asymmetrical (with respect to the $\\om$ plane) distortions of the magnetic field would pose a serious problems for the framework of the model leading to eq.~(\\ref{drh}). Again,  based on eq.~(\\ref{deflection}), any influence of the sweepback was neglected.  Given that the above-mentioned methods of determining $r$ are so sensitive to the assumed symmetry of the magnetic field around $\\vec \\mu$, it is important to have the symmetry hypothesis well justified. It is also important to revise this assumption in view of the unacceptably low values of $r$ which are often being derived with the BCW91 method: as found in BCW91, the ``delay radii\" $\\rdel$ implied by their method (eq.~\\ref{bcw}) are often smaller (in some cases by an order of magnitude --- see fig.~29 in BCW91) than the geometrical radii $\\rgeo$ determined with the traditional geometrical method based on the measurement of profile widths (Cordes 1978; Gil \\& Kijak 1993; Kijak \\& Gil 2002). This poses a real problem for the BCW91 method, because the geometrical radii, in the absence of strong refraction effects (Lyubarski \\& Petrova 1998), should be considered as lower limits of $r$ (Dyks et al.~2004a). Although one could explain this disagreement in many different ways (eg.~underestimated theoretical width of the open volume, systematically overestimated impact angles and dipole inclinations $\\alpha$ etc.), we show below that the rotational distortions of the static shape dipole may account for a large part of the discrepancies between $\\rdel$ and $\\rgeo$.  Recently Kapoor \\& Shukre (2003) considered the aberration effect \\emph{and} the rotational sweepback to investigate the relative locations of core and cone components in the pulsar magnetosphere. Although included in the model, the sweepback is again estimated with the help of eq.~(\\ref{deflection}). Being aware of the limitations of Shitov's estimate, the authors emphasized the need for derivation of a more advanced formula describing the rotational distortions of the magnetosphere. They noted that a proper derivation ``should make use of at least the magnetic field given by the full Deutsch solution (Deutsch 1955)\".  Such an estimate based on the Deutsch solution was done by Arendt \\& Eilek (1998), who concluded that the rotation distorts the magnetic field by a magnitude of the order of $r/\\rlc$. Being much larger than the Shitov's estimate, this distortion would strongly affect results in BCW91, GG2001, Hibschman \\& Arons (2001, hereafter HA2001), and Dyks et al.~(2004a). On the contrary, HA2001 noted that the leading terms in the difference between the Deutsch field and the rigidly rotating static-shape dipole are of the order of $(r/\\rlc)^2$. Recently, Mitra \\& Li (2004) emphasized that on the theoretical side there is a great need to develop and understand the details of the sweepback effect.  In this paper we investigate the rotational distortions of the pulsar magnetic field assuming the approximation of the vacuum magnetosphere. The twofold nature of the sweepback, first noticed by Shitov (1983) will be highlighted, and limitations in applicability of eq.~(\\ref{deflection}) will be clarified (Section 2). The significance of the sweepback for the relativistic model of pulsar polarization will appear to be much larger than previously thought, which will have serious consequencies for the delay-radius relation (eg.~modification of eq.~\\ref{bcw}, Section 3). ",
        "conclusions": "The rotational distortions of the vaccum dipole have twofold effect: in addition to the small changes of the local direction of the magnetic field, the region of the open magnetic field lines undergoes a strong distortion.  The change of the local direction of $\\vec B$ is second order in $\\rn$. We find, however, that it is symmetrical with respect to the $\\om$ plane. The largest asymmetrical change of $\\vec B$ direction is much smaller -- third order in $\\rn$, in agreement with Shitov's estimate.  The displacement of the open volume shifts the pulse window toward later phases with respect to the center of the position angle curve. The shift has the magnitude of the order of $\\rn^{1/2}$. The open volume shift modifies the delay-radius relation if the center of the pulse profile is determined as the midpoint between the outer edges of the pulse profile, which are assumed to lie close to the outer boundary of the open volume. At low altitudes, where effects of aberration and propagation time delays are small, the open volume shift dominates and may result in the center of the PA curve \\emph{preceding} the center of the profile (negative $\\dphnew$). A majority of pulsars exhibits positive $\\dphnew$ which means that the radio emission altitudes typically exceed $\\sim 10^{-2}\\rlc$.  The radii derived with the misalignment formula exceed those derived with the original delay-radius relation by a factor which increases quickly with decreasing altitude (and $\\dphnew$). This explains the trend of the delay-radius relation to predict emission radii smaller than the geometrical radii. The underestimate may also be produced/enhanced by the low-altitude emission within the central parts of the pulse profile. When both these effects work together, the standard delay-radius relation may underestimate $r$ by an order of magnitude even for the relatively large phase shifts ($\\dphnew \\sim 0.3^\\circ$).  The influence of the open volume shift on the method based on the core-cone shift (eq.~\\ref{drh}) is difficult to assess, because the locations of the conal maxima do not need to follow the open boundary of the open volume. If they did, eq.~(\\ref{drh}) would have to be replaced with $\\Delta\\phi_{\\sss DRH} \\approx 2\\rn - F\\rn^{1/2}$. The influence of the open volume shift on the geometrical method is small because the rotation increases the transverse size of the open volume insignificantly (by a factor smaller than $\\sim 1.2$, cf.~Fig.~\\ref{caps}).  The discussed distortion of the open volume is generated by a high-order effect which was being neglected for years on the basis that its ``local magnitude\" at low altitudes is small ($\\sim \\rn^2$). The actual importance of this effect appears to be much larger than that. This suggests that other ``high-order\" effects, eg.~the longitudinal polar cap currents, of magnitude $\\sim \\rn^{3/2}$, or the inertia of electrons, may be a lot more significant than their low altitude magnitude suggests. The toroidal currents due to the corotation of the charge-filled magnetosphere ($\\sim \\rn^2$) have been shown to notably modify the shape of the open volume, however, in a way which is symmetrical with respect to the $\\om$ plane (cf.~fig.~4.11 in Beskin et al.~1993)."
    },
    "0402/astro-ph0402285_arXiv.txt": {
        "abstract": "Some phenomenological properties of the unidentified EGRET detections suggest that there are two distinct groups of galactic gamma-ray sources that might be associated with compact objects endowed with relativistic jets. We discuss different models for gamma-ray production in both microquasars with low- and high-mass stellar companions. We conclude that the parent population of low-latitude and halo variable sources might be formed by yet undetected microquasars and microblazars. ",
        "introduction": "Population studies indicate that there are at least three different groups of galactic gamma-ray sources (Gehrels et al. 2000; Grenier 2001, 2004). One is a group of bright sources distributed along the galactic plane, with a concentration toward the inner spiral arms. These sources are well-correlated with young stellar objects and star-forming regions (Romero et al. 1999, Romero 2001). We will call these sources Gamma-Ray Population I (GRP I). There is another group of sources distributed around the Galactic Center, forming a halo with a radius of $\\sim 60^{\\circ}$. These sources are softer and more variable than GRP I sources. We will call this group Gamma-Ray Population II (GRP II). Finally, there is a group of sources correlated with the Gould Belt (Grenier 1995, 2000). These are nearby, relatively young and weak sources. We will name them the Local Gamma-Ray Population (LGRP).  Recent variability studies (Nolan et al. 2003; Torres et al. 2001a,b) show that GRP II sources are very variable. There is also a subgroup of variable sources among those in the GRP I. This variability, with timescales from weeks to months, indicates that the counterparts are compact or small objects. Supernova remnants, OB star forming regions, and other extended sources are ruled out. Isolated pulsars, which are known to be steady sources, are also discarded.  Recently, Kaufman Bernad\\'o et al. (2002) have suggested that the parent population of GRP I sources are microquasars with high-mass stellar companions and jets that form a small angle with the line of sight, i.e. the ``microblazar\" population suggested by Mirabel \\& Rodriguez (1999). In the Kaufman Bernad\\'o et al. (2002) model, the gamma ray emission is produced by inverse Compton (IC) up-scattering of UV photons from the companion stars, as suggested by Paredes et al. (2000) for the microquasar LS 5039 (see also Bosch-Ramon \\& Paredes 2004). More recent models (Romero et al. 2003) incorporate the effect of hadrons in the interaction of the jet with the wind of the star. In the present paper, we extend the Kaufman Bernad\\'o et al. (2002) proposal to encompass both GRP I and GRP II variable sources. We suggest that old low-mass microquasars, either ejected from the galactic plane or from globular clusters, might be the counterparts of the GRP II. For GRP I sources, we suggest that a subgroup of them (the most variable ones) might be high-mass microquasars. Some of them might have soft TeV counterparts. In the next sections we present some models for these objects. ",
        "conclusions": "There are variable gamma-ray sources in the Galaxy. There are reasons to think that at least two different populations of such sources exist: one population of young sources in the galactic plane, and one population of old sources at high latitudes, forming a halo around the galactic center. We suggest that these two groups of sources correspond to high-mass microquasars and low-mass microblazars, respectively. Our calculations indicate that the required high-energy spectra and luminosities can be produced by these objects. In a separate communication (Kaufman Bernad\\'o et al. in preparation), a detailed discussion of the source statistics will be presented. Some of these sources, especially those at low latitudes with a peak at MeV energies (see Fig. \\ref{fig2}), might be detectable by INTEGRAL. The recent findings by Combi et al. (2004) seem to be a promising step in this direction.   \\subsection*"
    },
    "0402/astro-ph0402291_arXiv.txt": {
        "abstract": "We have obtained high-resolution ({\\it R} $\\simeq$ 50,000 or 90,000), high-quality (S/N $\\ga$ 100) spectra of 22 very metal-poor stars ([Fe/H] $\\la$ --2.5) with the High Dispersion Spectrograph fabricated for the 8.2m Subaru Telescope. The spectra cover the wavelength range from 3500 to 5100~{\\AA}; equivalent widths are measured for isolated lines of numerous elemental species, including the $\\alpha$ elements, the iron-peak elements, and the light and heavy neutron-capture elements. Errors in the measurements and comparisons with previous studies are discussed. These data will be used to perform detailed abundance analyses in the following papers of this series. Radial velocities are also reported, and are compared with previous studies. At least one moderately r-process-enhanced metal-poor star, HD 186478, exhibits evidence of a small-amplitude radial velocity variation, confirming the binary status noted previously. During the course of this initial program, we have discovered a new moderately r-process-enhanced, very metal-poor star, CS~30306--132 ([Fe/H] $= -2.4$; [Eu/Fe] $= +0.85$), which is discussed in detail in the companion paper. ",
        "introduction": "Very metal-poor stars ([Fe/H] $\\la -2.5$)\\footnote{We use the usual notation [A/B]$\\equiv \\log_{10} (N_{\\rm A}/N_{\\rm B})_* - \\log_{10}(N_{\\rm A}/N_{\\rm B})_\\odot$ and log$\\epsilon({\\rm A})\\equiv \\log_{10}(N_{\\rm A}/N_{\\rm H})+12.0$, for elements A and B. Also, the term ``metallicity'' will be assumed here to be equivalent to the stellar [Fe/H] value.} are believed to have been born in the early Galaxy; their chemical compositions are living records of the nucleosynthesis processes that preceded their formation. As a result of considerable efforts by many astronomers, a large list of candidate stars with [Fe/H] $< -2.5$ have been provided by wide-field objective-prism surveys in the past two decades (e.g., the HK survey: Beers et al. 1985, 1992; Beers 1999, and the Hamburg/ESO Survey: Christlieb \\& Beers 2000; Christlieb et al. 2001; Christlieb 2003). Over the past several years, high-resolution spectroscopic studies have enabled the measurement of elemental abundances for many of the metal-poor stars found by these surveys (e.g., McWilliam et al. 1995b; Ryan, Norris, $\\&$ Beers 1996; Burris et al. 2000; Carretta et al. 2002; Cayrel et al. 2004), including detailed studies of the lowest metallicity stars yet identified (e.g., Norris, Ryan, \\& Beers 2001; Christlieb et al. 2002). These observational studies, which continue at present, are providing strong constraints on models of the dominant nucleosynthesis processes in the earliest epochs of star formation in our Galaxy, in particular those associated with massive stars and Type II supernovae.  Remarkable progress has been made, in particular, through studies of the neutron-capture elements in very metal-poor stars. High-resolution spectroscopic studies of very metal-poor stars have revealed, for example, that a small fraction (presently estimated to be on the order of 2\\%--3\\%, Beers, private communication) of giants with [Fe/H] $< -2.5$ exhibit large overabundances (e.g., [r-process/Fe] $ > +1.0$) of neutron-capture elements associated with the r-process (e.g., [r-process/Fe] $ > +1.0$ ;McWilliam et al. 1995b; Sneden et al. 2000, 2003; Cayrel et al. 2001; Hill et al. 2002). These, along with a handful of other metal-poor stars with moderately enhanced r-process elements ($+0.5 \\leq$ [r-process/Fe] $\\leq +1.0$, e.g., Westin et al. 2000; Johnson \\& Bolte 2001; Cowan et al. 2002), display remarkably similar abundance patterns in the range 56 $\\leq Z <$ 76, all apparently in good agreement with the solar-system r-process component. In addition, some of the neutron-capture-enhanced, metal-poor stars exhibit abundance patterns associated with s-process nucleosynthesis (e.g., Norris et al. 1997; Van Eck et al. 2001; Aoki et al. 2002; Lucatello et al. 2003).  These efforts are having a large collective impact on studies on the origin of the neutron-capture elements in the Galaxy \\citep[e.g.,][]{ishimaru99,fields02,qian02}, and on the underlying physics and astrophysical sites of the r- and s- processes \\citep[e.g.,][]{gallino98,wanajo02,wanajo03,schats02,truran02}. Furthermore, detailed studies of the r-process-enhanced, very metal-poor stars have provided new, potentially quite powerful, methods for obtaining hard lower limits on the age of the Galaxy and the universe, from the application of cosmo-chronometry based on the observed (present-day) abundance ratios of radioactive nuclei (Th and U), as compared with one another, and with stable elements originating in the r-process, \\citep[e.g., Eu,][]{sneden96,wes00,cayrel01,schats02,wanajo02,wanajo03,sneden03}.  In order to develop a more clear understanding of the individual nucleosynthetic processes that were operating in the early Galaxy, further abundance studies are required, based on high-quality spectra, for much larger samples of very metal-poor stars than have been examined to date. We have initiated such a set of investigations with the Subaru Telescope High Dispersion Spectrograph (HDS, Noguchi et al. 2002). In this paper we present observations of 22 very metal-poor stars observed during the commissioning phase of this instrument. In \\S 2 we discuss the selection of targets and details of the observations that have been carried out. Our spectra cover the wavelength range from 3500 to 5100~{\\AA} with high spectral resolution (a resolving power of $R = 50,000$ or $R = 90,000$) and high signal-to-noise (S/N$ \\gtrsim$ 100 per resolution element). We report the equivalent widths measured for the spectra in \\S 3, where we also discuss the random errors of our measurements, and make comparisons with previous studies of stars in common. Radial velocity measurements for our program stars are presented in \\S 4, along with a comparison with previous measurements for a number of stars. These data will be used in the detailed abundance analyses that will follow in additional papers of this series. ",
        "conclusions": "We have obtained high-resolution, high-S/N ratio, spectra for 22 very metal-poor stars with Subaru/HDS, taken during the commissioning phase of this instrument. These stars were selected so as to include as many objects with known (or suspected) enhancement of r-process elements as possible. In this paper we have reported the measurements of equivalent widths for isolated absorption lines in the reduced spectra, and also precision radial velocities (in some cases, at several epochs), for each star. Comparisons of our measured equivalent widths with previous work demonstrates that there exists no systematic differences, except in the cases of stronger lines in a few objects for which the S/N ratios of the previous work was rather low. In the following papers of this series (Paper II, others in preparation), the results of the detailed abundance analysis for these data will be presented."
    },
    "0402/astro-ph0402258_arXiv.txt": {
        "abstract": "{ The giant radio galaxy M87 is usually classified as a Fanaroff-Riley class I source, suggesting that M87 is a mis-aligned BL Lac object. Its unresolved nuclear region emits strong non-thermal emission from radio to X-rays which has been interpreted as synchrotron radiation. In an earlier paper we predicted M87 as a source of detectable gamma ray emission in the context of the hadronic Synchrotron-Proton Blazar (SPB) model. The subsequent tentative detection of TeV energy photons by the HEGRA-telescope array would, if confirmed, make it the first radio galaxy to be detected at TeV-energies.  We discuss the emission from the unresolved nuclear region of M87 in the context of the SPB model, and give examples of possible model representations of its non-simultaneous spectral energy distribution.  The low-energy component can be explained as synchrotron radiation by a primary relativistic electron population that is injected together with energetic protons into a highly magnetized emission region.  We find that the $\\gamma$-ray power output is dominated either by $\\mu^\\pm$/$\\pi^\\pm$ synchrotron or proton synchrotron radiation depending on whether the primary electron synchrotron component peaks at low or high energies, respectively.  The predicted $\\gamma$-ray luminosity peaks at $\\sim$100 GeV at a level comparable to that of the low-energy hump, and this makes M87 a promising candidate source for the newly-commissioned high-sensitivity low-threshold Cherenkov telescopes H.E.S.S., VERITAS, MAGIC and CANGAROO III. Because of its proximity, the high-energy spectrum of M87 is unaffected by absorption in the cosmic infrared (IR) background radiation field, and could therefore serve as a template spectrum for the corresponding class of blazar if corrected for mis-alignment effects. This could significantly push efforts to constrain the cosmic IR radiation field through observation of more distant TeV-blazars, and could have a strong impact on blazar emission models.  If M87 is a mis-aligned BL-Lac object and produces TeV-photons as recently detected by the HEGRA-array, in the context of the SPB model it must also be an efficient proton accelerator. ",
        "introduction": "The Fanaroff-Riley (FR) class I giant radio galaxy M87, situated nearly at the center of the Virgo cluster, was the first extragalactic jet to be discovered (\\cite{old}), and has since then been intensively observed at all wavelengths. Its proximity ($\\sim 16.3$ Mpc; \\cite{Cohen2000}) makes it an interesting laboratory for testing and understanding extragalactic jets of radio-loud Active Galactic Nuclei (AGN) and their powering engines.  Because M87 is sufficiently near for ultra-high-energy cosmic rays (UHECRs) to be little affected by the GZK-cutoff at $\\sim 5\\times 10^{19}$ eV (\\cite{Gre66}, \\cite{Zat66}), and because its size scales could allow magnetic confinement of the most energetic cosmic rays (\\cite{Hillas}), M87 has long been considered as one of the prime candidate sources of high energy cosmic rays. This idea has recently received some support by the suggestion that possible clustering observed in the arrival directions of the UHECRs can be understood in terms of deflection of UHECRs from M87 by our Galaxy's magnetized wind assuming a Parker-spiral magnetic structure (\\cite{Ahn2000}; \\cite{Biermanetal2001}). It appears, however, that such deflection in the Galactic wind may be insensitive to the direction of the cosmic ray sources (\\cite{Billoir2000}). Nevertheless, \\cite{Ray2003} found by using the cosmic ray output predicted in the SPB model that M87 could explain the observed UHECR flux if the magnetic field topology between M87 and our Galaxy were favourable.  They found that UHECR with energies above $10^{20}$ eV could easily be produced because neutrons produced in the pion photoproduction process would be relativistically beamed along the jet direction and Doppler boosted in energy.  Even though M87's jet is mis-aligned with respect to our line-of-sight, these Doppler boosted neutrons escape from the jet and decay into UHECRs which maintain their Doppler boosted energies and may propagate in all directions, including towards our Galaxy if the magnetic field topology were favourable. Of course we note that M87's nuclear region is not the only possible source of the UHECRs observed at Earth; see e.g. \\cite{ProtheroeClay2004} for a recent review.   According to the unification model of AGN (e.g. \\cite{UrryPadovani95}) FR~I radio galaxies, with their jet axis at a large angle to our line-of-sight, are the parent population of BL Lac objects whose jets are closely aligned to our line-of-sight. This motivates us to consider M87's nuclear region as a mis-aligned blazar of BL Lac type.  The spectral energy distribution (SED) of BL Lac objects can usually be explained satisfactorily by either leptonic or hadronic blazar emission models. \\cite{BaiLee2001} have discussed M87 on the basis of the leptonic Synchrotron-Self Compton (SSC) model where synchrotron photons produced by interactions of relativistic electrons with the ambient magnetic field serve as the target photons for inverse Compton scattering by the same electrons. By interpreting the non-thermal radiation from the radio to the X-ray band as synchrotron emission with luminosity peaking in the far-ultraviolet, the authors considered M87 to be a mis-aligned high-frequency peaked BL Lac (HBL), and predicted $\\gamma$-ray emission with an inverse Compton peak at $\\sim$100 GeV. The predicted inverse Compton flux is consistent with the recent HEGRA detection of M87 (\\cite{HEGRAdet}, see Sect.~2).  Detectable TeV-emission from Comptonization of galactic photon fields has recently been suggested by \\cite{Stawarz2003}.  In contrast to former models, they consider, however, the large scale jet to be the site of $\\gamma$-ray production.  While in leptonic models a relativistic electron-positron plasma is usually assumed to be responsible for the non-thermal jet radiation, in hadronic models a relativistic proton-electron (p e$^-$) plasma is assumed to be the main constituent of the jet material.  In the hadronic Synchrotron-Proton Blazar (SPB) model, proposed recently by, e.g., \\cite{MP2001}, accelerated protons interact with the synchrotron radiation field produced by the co-accelerated electrons via meson photoproduction and Bethe-Heitler pair production and, more importantly, with the strong ambient magnetic field emitting synchrotron radiation (mesons and muons also emit synchrotron radiation). The SPB model neglects external photon field components, and this seems appropriate for BL~Lac objects and their parent population which possess only weak accretion disks.  \\cite{MP2001} have shown that this model can reproduce the commonly observed double-humped blazar SED.  Hadronic models require high proton energies that can only be achieved in a highly magnetized environment where synchrotron losses can become severe. Magnetic field values around $10^3$~G are thought to exist near the horizon of a supermassive black hole (\\cite{BZ77}) with a mass of $\\sim$$10^9$~M$_{\\sun}$ as estimated for M87 (\\cite{Marconi1997}).  However, with M87's rather low accretion rate if the equipartition value of $B$ scales with $\\dot M$, and assuming magnetic energy flux conservation, magnetic field strengths of order 10-100~Gauss are expected within 30 Schwarzschild radii $r_g$ where the jet is probably formed (\\cite{Nature}). In the present work, we discuss in more detail than our earlier work (\\cite{Ray2003}), and in the context of the recent HEGRA detection (\\cite{HEGRAdet}), the nuclear (core) emission, i.e.\\ from the M87 jet, in the framework of the SPB-model.  In Section 2 we summarize the data on M87's core emission.  In section 3 we give a brief model description, and calculate the steady-state synchrotron component as described in the appendix. The modeling procedure is described in section 4, and we conclude with a summary and discussions in section 5. ",
        "conclusions": "We have made SPB model fits to the non-simultaneous SED of M87's nuclear emission, and find that all parameter sets which satisfactorily represent the data predict the main contribution to the high energy luminosity at about 100 GeV to be due to either $\\mu^\\pm$/$\\pi^\\pm$ synchrotron or proton synchrotron radiation depending on whether the primary electron synchrotron component peaks at low (Model L) or high (Model H) energies, respectively.  In the EGRET energy range, the lower synchrotron peak energy model (Model L) predicts a softer spectrum than the fit with a higher synchrotron peak energy (Model H).  While it is obvious that EGRET's sensitivity was more than an order of magnitude above the expected flux level from M87, we find that the satellite-based $\\gamma$-ray instrument GLAST might possibly detect a weak signal from this radio galaxy (see Fig.~\\ref{predictfig}).  In all the fits we have presented, the high energy radiation cuts off with a strong steepening in the TeV range in agreement with the spectral limits from the HEGRA observation. We also find that the HEGRA detection at $>$730 GeV can only be explained if the proton acceleration rate is extremely high ($\\xi(E'_{p\\rm{,max}}) \\approx 1$). We therefore expect M87, if it is indeed a mis-aligned SPB, could be an important source of UHECRs (see also \\cite{Ray2003}).   For almost all proposed models of particle acceleration in different astrophysical environments, $\\xi(E)$ remains a rather uncertain model parameter.  On the other hand, any postulation of acceleration of high energy protons in compact $\\gamma$-ray production regions actually implies that $\\xi(E'_{p\\rm{,max}})$ at these energies should be close to unity, which corresponds to the maximum theoretically possible acceleration rate based on simple geometrical consideration (e.g. \\cite{Hillas}). An interesting possibility could be particle acceleration at the annihilation of magnetic fields in the fronts of poynting flux dominated jets (\\cite{Blandford76}; \\cite{Lovelace76}).  It has been argued that this mechanism could provide effective acceleration of extremely energetic protons with $\\xi(E'_{p\\rm{,max}}) \\sim 1$ (Haswell et al. 1992).  A quantitative investigation of particle acceleration mechanisms is beyong the scope of this paper.   Both model fits presented seem to under-predict the emission in the radio domain as compared to the observations. Our modeling, however, assumes the same size for the emission region at all energies, while the data indicate a smaller width of the optical than the radio jet (\\cite{Sparks1996}) though with {\\emph{roughly}} the same morphology.  In addition, the inter-knot region is observed to be weaker in the optical than in the radio band (\\cite{Sparks1996}). It appears therefore reasonable to attribute the missing flux in our model to the inter-knot region, and to a larger blob size in the radio band as compared to higher frequencies.   As previously noted, the Chandra data lie above the extrapolation of the optical spectrum to higher energies (\\cite{WilsonYang2002}). We point out that the modelled SED of M87 is based on non-simultaneous data, and that the X-ray flux could have been much lower than the Chandra data suggest at the time of the optical observations.  Another critical point is that different resolutions of the images at the various energies have been used in the literature for the flux determinations.  We have used for our compilation of the SED the highest resolution data available at each frequency, ranging from arcmin (HEGRA) to sub-arcsecond (Chandra, Hubble, VLBI) scales, and could introduce additional non-negligible uncertainties into the flux measurements.  However, a flatter X-ray spectrum consistent with the Chandra data might be achieved if either the magnetic field is highly inhomogenous, or a secondary e$^\\pm$ population is responsible for the X-ray flux (\\cite{WilsonYang2002}).  While, in principle, the SPB model provides secondary synchrotron emitting e$^\\pm$s from the various cascades which may extend even into the X-ray domain, the fits presented can not easily explain the high flux level observed at these energies as electron synchrotron radiation or as part of the high energy hump cascade component.  Nevertheless, for magnetic fields of order a few Gauss SSC radiation might become detectable also in hadronic models.  Fig.~\\ref{LBLfig} shows that the SSC component peaks at X-ray energies with a spectral signature that is in agreement with the Chandra observations, but its flux is roughly an order of magnitude too low to explain the Chandra data.  If the magnetic field or the size of the emission region were a factor $\\sim 3$ lower model L would predict the Chandra X-rays as SSC photons. The latter change moves model L slightly in the direction of model H. Thus, X-ray variability might be related to relatively small changes in the model paramaters, raising or lowering the importance of inverse Compton losses with respect to synchrotron losses.   Another possible source for the observed emission in the Chandra band might be a contribution, either directly or reprocessed through cascading, from photon fields other than the primary electrons' synchrotron radiation. The spectral continuum data do not show any thermal bump from the putative accretion disk. \\cite{dimatteo} have shown that an advection dominated accretion disk could account for a large fraction of the observed X-ray nuclear flux. The radiative efficiency is extremely low, but the accretion rate is found to be large enough for Comptonization of the synchrotron emission of the disk or the thermal bremsstrahlung emission to dominate the X-ray emission (\\cite{dimatteo}). A sudden drop of one order of magnitude of the accretion rate could lower the X-ray disk output by $\\sim 2$ orders of magnitude.  For an advection dominated disk with a X-ray luminosity $L_{\\rm disk}\\approx 7\\times 10^{40}$~erg/s (\\cite{dimatteo}) the total energy density in the jet frame $u'_{\\rm disk}$ can be derived through the transformation (see e.g.\\ \\cite{DermerSchlicki2002}) \\[ {u'_{\\rm disk}} \\approx {1.7 \\cdot 10^{10}\\mbox{eV cm}^{-3}{[\\Gamma_j(1-\\beta_j} \\mu(r))]^{2}} \\left(\\frac{L_{\\rm disk}}{10^{40}\\mbox{erg s$^{-1}$}}\\right) \\left(\\frac{r/\\mu(r)}{10^{15}\\mbox{cm}}\\right)^{-2} \\] where $r$ is the distance of the jet plasma blob from the black hole, $\\mu(r)=r/(r^2+r_{\\rm in}^2)^{1/2}$ and we have approximated the disk's radiation field as a luminous ring of radius $r_{\\rm in}$ that illuminates the moving blob.  Assuming a inner accretion disk radius of $r_{\\rm in}=1.23 r_g$ (for an extreme Kerr black hole of mass $M=10^9M_\\odot$) and taking $\\Gamma_j=1.5$ one obtains ${u'_{\\rm disk}} \\sim 2\\cdot 10^{10}$eV cm$^{-3}$ at $r=10^{15}$cm and $\\sim 2\\cdot 10^8$eV cm$^{-3}$ at $r=10^{16}$cm.  Hence, for $r\\ga 10^{16}$cm the accretion disk radiation proves to be unimportant as a target field for cascading and photon-particle interactions in M87 compared to the primary electron synchrotron emission.   Recently \\cite{DoneaProtheroe2003} have constrained the torus temperature of the torus to $ <250$~K using existing data from the literature.  On the other hand, during an extreme flaring state, related to the accretion rate changes or to a spin flip of the central black hole, the torus could undergo enough heating to become 'visible'. This alters the high energy part of the spectrum above several hundreds of GeV, as is discussed in \\cite{DoneaProtheroe2003}.  Therefore, a visibility-state of the torus (if present) could be achieved at the cost of not being able to observe very high energy gamma rays from the nucleus of M87. Regular monitoring of M87 at VHE gamma-rays and IR frequencies could be important to elucidate the problem of existence or non-existence of a dusty torus in M87.  For a temperature of the torus radiation of $< 250$K the co-moving frame energy density is $2\\cdot 10^{7} \\Gamma_j^2$ eV cm$^{-3} \\ll u'_{\\rm{phot}}$ for M87, and is therefore negligible as target photon field. The star and dust contribution of M87's host galaxy has been estimated to $630 \\Gamma_j^2$ eV cm$^{-3}$ and $6.3 \\Gamma_j^2$ eV cm$^{-3}$ in the jet frame, respectively (\\cite{Stawarz2003}), and can obviously also be neglected regarding M87's synchrotron radiation density of order $10^{10}$ eV cm$^{-3}$.    Fig.~\\ref{predictfig} shows that the recently commissioned Cherenkov telescope array VERITAS, the southern arrays H.E.S.S.\\ and CANGAROO III (though at large zenith angles $> 45\\degr$), and MAGIC may be able to detect M87.  The predicted integral fluxes $> 100$ GeV for Models H and L are $\\sim 4\\times 10^{-11}$~cm$^{-2}$~s$^{-1}$ and $\\sim 4\\times 10^{-12}$~cm$^{-2}$~s$^{-1}$, respectively.  We have used A.~Konopelko's simulator for the H.E.S.S. response (http://pluto.mpi-hd.mpg.de/~konopelk/WEB/simulator.html) to estimate the necessary observation time for statistically significant detection.  A 10~h on-source observation with the full phase I (four telescopes) H.E.S.S.\\ array would result in a $8-9 \\sigma$ detection (expected cosmic ray rate is $\\sim$ 0.7~s$^{-1}$, $\\gamma$-ray rate is 0.055~s$^{-1}$) in the case of a high-energy peaked photon target, and $4-5 \\sigma$ detection in the case of the low-energy peaked photon target for 300~h of usable data assuming the source at zenith (expected cosmic ray rate is here $\\sim$ 0.7~s$^{-1}$, $\\gamma$-ray rate is 0.006~s$^{-1}$).  Since the sensitivity of VERITAS (\\cite{VERITASsens}) is similar to that of H.E.S.S., similar numbers can be expected for VERITAS observations.  In Fig.~\\ref{predictfig} we summarize the minimum flux for a 50~h observation (with statistics exceeding 10 photons and a signal detection at a level of at least 5$\\sigma$) using the phase I H.E.S.S.\\ array, the VERITAS array, CANGAROO III and MAGIC (assuming the source at zenith) and GLAST, in comparison the the predicted high energy fluxes.  Note, however, that these predictions are based on a non-simultaneous observed SED and, depending on the actual activity state of M87, the predicted fluxes and spectra may change significantly. In addition, absorption of $\\gamma$-rays in infrared radiation from a putative torus could affect the spectrum above 1~TeV if the torus temperature $T_{\\rm{torus}}$ were higher than 250~K, and above 200~GeV if $T_{\\rm{torus}}\\geq 1000$~K (\\cite{DoneaProtheroe2003}). Work is in progress to make SPB model fits to other nearby FR~I radio galaxies.  In both models presented here, the power output in the high energy hump is roughly equal to the power output in the low energy hump of the SED. Because of M87's proximity, absorption of sub-GeV/TeV-photons in the cosmic infrared background radiation field is not expected to affect the spectrum below $\\sim$ 50~TeV. {\\emph{The observed spectral behaviour at high energies should be intrinsic to the source.}} Tracing the spectrum at GeV-TeV-energies would give a $\\gamma$-ray spectrum that for the first time includes an unabsorbed (by radiation fields external to the source) cutoff.  These data could serve as a typical template BL~Lac spectrum at source after correcting for M87's jet mis-alignment. By comparing this template with BL~Lac spectra at high redshifts, meaningful constraints for the extragalactic background radiation field around IR wavelengths can be derived."
    },
    "0402/astro-ph0402544_arXiv.txt": {
        "abstract": "Evolutionary models taking into account radiative accelerations, thermal diffusion, and gravitational settling for 28 elements, including all those contributing to OPAL stellar opacities, have been calculated for solar metallicity stars of 0.5 to 1.4 \\Msol{}.  The Sun has been used to calibrate the models.  Isochrones are fitted to the observed color-magnitude diagrams (CMDs) of M$\\,$67 and NGC$\\,$188, and ages of 3.7 and 6.4 Gyr are respectively determined.  Convective core overshooting is {\\it not} required to match the turnoff morphology of either cluster, including the luminosity of the gap in M$\\,$67, because central convective cores are larger when diffusive processes are treated.  This is due mainly to the enhanced helium and metal abundances in the central regions of such models.  The observation of solar metallicity open clusters with ages in the range 4.8--5.7 Gyr would further test the calculations of atomic diffusion in central stellar regions: according to non-diffusive isochrones, clusters should not have gaps near their main-sequence turnoffs if they are older than $\\approx 4.8$ Gyr, whereas diffusive isochrones predict that gaps should persist up to ages of $\\approx 5.7$ Gyr.  Surface abundance isochrones are also calculated.  In the case of M$\\,$67 and NGC$\\,$188, surface abundance variations are expected to be small. Abundance differences between stars of very similar \\teff{} are expected close to the turnoff, especially for elements between P and Ca. Moreover, in comparison with the results obtained for giants, small generalized underabundances are expected in \\MS{} stars.  The lithium to beryllium ratio is discussed briefly and compared to observations.  The inclusion of a turbulent transport parametrization that reduces surface abundance variations does not significantly modify computed isochrones. ",
        "introduction": "\\label{sec:intro} The open clusters M$\\,$67 and NGC$\\,$188 have about the solar metallicity, bracket the solar age, and have turnoff stars only a few hundred degrees hotter than the Sun.  As such, they are interesting testing grounds for the effects of atomic diffusion on age determinations and surface abundances, since, in the case of the Sun, there is now ample evidence from heliosismology that atomic diffusion has reduced the surface He abundance (\\citealt{GuzikCo92}; \\citealt{GuzikCo93}; \\citealt{ChristensenDalsgaardPrTh93}; \\citealt{Proffitt94}; \\citealt{BahcallPiWa95}; \\citealt{GuentherKiDe96}; \\citealt{RichardVaChetal96}; \\citealt{BrunTuZa99}).  Indeed, diffusive processes presumably also cause small underabundances of metals in the Sun: these are caused mainly by gravitational settling, but are also modified by radiative accelerations (\\gr{}), which are predicted to be especially important at the end of the \\MS{} phases of solar-type stars \\citep{TurcotteRiMietal98}.  What abundance anomalies are then to be expected in the turnoff stars of M$\\,$67, which are $\\sim 400$ K hotter than the Sun (\\citealt{HobbsTh91}), and in those of NGC$\\,$188, which are $\\sim 100$ K hotter (\\citealt{HobbsThRo90})?  As the cluster turnoff stars are expected to have smaller surface convection zones, they may show larger effects of atomic diffusion than the Sun.   On the other hand, since the radius and age of the Sun are used to calibrate the mixing length and assumed initial He abundance, this normalization may eliminate the effects of diffusion on age determinations.  Whereas an $\\approx 10$\\% reduction in age at a given turnoff luminosity --- compared with the predictions of models that neglect diffusion --- was derived by \\citet{VandenBergRiMietal2002} from the diffusive models for Population II stars computed by \\citet{RichardMiRietal2002}, it is not clear that a similar reduction should be expected in the case of Pop.~I stars.  In addition, there may be some important differences in the morphologies of the diffusive and non-diffusive isochrones in the age range where a transition is made between isochrones that have a gap near the main-sequence turnoff and those which do not.  One naively expects that both the sizes of convective cores in models for main-sequence stars of a given mass, and the predicted mass marking the transition between stars that have convective and radiative cores on the main sequence, will depend (to some extent) on whether or not diffusive processes are treated.  (For instance, the concomitant increase in opacity with the settling of Fe in the cores of stars would tend to enhance convective instability.)  In this regard, we note that the first studies of NGC$\\,$188 (\\citealt{Sandage62}; \\citealt{EggenSa69}) concluded that it has a gap near the top of its main-sequence on the $(V,\\;B-V)$-diagram reminiscent of that seen in M$\\,$67.  \\citet{McClureTw77} carried out a statistical test of the photographic photometry that they obtained for the same cluster, and confirmed the existence of the gap in the color-magnitude diagram (CMD) that was constructed for stars within Ring I on Sandage's original finder chart.  Curiously, no statistically significant evidence for a gap was found if their CMD included stars in Sandage's Ring II; but McClure \\& Twarog concluded that contamination by field stars was almost certainly much more severe in the outer ring and that, if it were possible to remove the field stars, ``the gap would be obvious\".  The proper-motion membership study of \\citet{DinescuGiAletal96} does not shed any light on this problem (because of the very large scatter in their CMD fainter than $V=15$: the gap is located at $V\\approx 15.5$ according to McClure \\& Twarog).  However \\citet{PlataisKoMaetal2003} provide a well defined CMD down to $V= 20$, with no indication of a turnoff gap.  Because there is no obvious indication of a gap in subsequent CMDs for NGC$\\,$188 (e.g., \\citealt{Kaluzny90}; \\citealt{CaputoChCaetal90}; \\citealt{SarajediniHiKoetal99}), and because the best-fitting isochrones for current best estimates of the cluster distance and reddening do not predict a gap at the turnoff $M_V$ (see the aforementioned papers), its existence is considered by many to be quite doubtful. However, the models that have been compared with the cluster CMD have not taken gravitational settling and radiative accelerations into account. If it were shown that diffusive isochrones do, in fact, predict a main-sequence gap, not necessarily for NGC$\\,$188 but for any range in age where a gap is not predicted by models that neglect diffusion, this would be an important development that would motivate a search for open clusters within the requisite age range to further test our understanding of stellar physics.  There are many other questions that need to be addressed.  In particular, does the same turbulence parametrization that Richard et al.~(2002) used, in conjuction with diffusion physics, to explain the Li abundances in field halo stars, also lead to good agreement between the predicted and observed Li abundances in solar-type stars, as well as in those that were in the Li gap at the age of the Hyades \\citep{Balachandran95}?  The surface abundances of Li in M$\\,$67 solar-type stars (see Fig.~7 of \\citealt{MartinBaPaetal2002}) vary from star to star at a given \\teff{}, which suggests that mixing processes below the surface convection zone vary from star to star at a given mass.  Did the turbulence differ only in the early stellar histories of the cluster stars or is it still different between one solar twin and another?  To what extent is this confirmed by abundance anomalies of other species and do such anomalies affect theoretical isochrones?  Furthermore, what abundance anomalies of Fe, Li, C, and O could be caused by diffusion and are they observed \\citep{BarrettBoKietal2001}?  In NGC$\\,$188, the surface Li abundance in solar-type stars appears to be consistent with a single-valued function of $\\teff$ just as in the Hyades (see \\citealt{RandichSePa2003}).  Furthermore, even though NGC$\\,$188 is older than M$\\,$67 by $\\sim 2$--$3\\times 10^9$ years (e.g., Sarajedini et al.~1999), the Li abundance in its G-type stars is comparable with the largest abundances measured in M$\\,$67 stars having similar colors/temperatures.  Could a simpler model account for the Li observations in the Hyades and NGC\\,188 than in M\\,67? Pre--\\MS{} evolution could be largely responsible for the Li destruction in the Hyades and NGC\\,188 (see, for instance, \\citealt{ProffittMi89, PiauTu2002}). Finally, we note that the relative abundances of Li and Be along the subgiant branch of M$\\,$67 has been evaluated in models including gravitational settling by \\cite{SillsDe2000}.  However, what are the ratios of the Li and Be abundances to be expected from diffusion if \\gr{} are also taken into account?  In this paper, after a very brief description of the calculations in \\S \\ref{sec:calcul}, the chemical composition expected on the surfaces of stars of M\\,67 and NGC\\,188 will be discussed in \\S \\ref{sec:composition} and Li/Be ratios in \\S \\ref{sec:li_be}. The effect of atomic diffusion on central convective cores is analyzed in detail in \\S \\ref{sec:CZ}. The effect of diffusion on isochrones is discussed in  \\S \\ref{sec:iso} and   these results are applied to M\\,67 and NGC\\,188 in \\S \\ref{sec:opencl}. The main conclusions are summarized in \\S \\ref{sec:conclusion}.  Throughout this paper the emphasis is on calculations in the presence of atomic diffusion and, in some cases, of turbulent transport with the same parametrization as used for Pop II stars by \\citet{RichardMiRietal2002}. The discussion of potential star--to--star variations of turbulent transport to explain Li abundance spread at a given \\teff{} is left to a paper in preparation. ",
        "conclusions": "\\label{sec:conclusion} Since the Sun is used to normalize convection parameters and initial abundances, one could have imagined that, in solar metallicity clusters having ages similar to that of the Sun, models with diffusion would lead to the same age and the same CMD properties as models without diffusion. The variations of $\\alpha$ and initial abundances required to fit the Sun in the diffusion model reproduce the same 1.0 \\Msol{} star at the same age and the two sets of models could be expected to do the same for star clusters. Reality turns out to be more complex.  For age determinations, partial cancellation effectively occurs, but the shapes of isochrones turn out to be quite different near the turnoff.  Both the normalization to solar abundances and the additional gravitational settling in the central regions of stars work together to cause an 18\\,\\% increase in the central metallicity. This increases the size of the convective core in stars of 1.09 to 1.3 \\Msol{} (see \\S \\ref{sec:CZ}) which, in turn, modifies the morphologies of isochrones (see \\S \\ref{sec:iso}) around the solar age.  An important consequence of the changes in the shapes of isochrones that arise when diffusive processes are treated is that it is possible to match the CMD of M$\\,$67 (including the luminosity of the gap near the turnoff) without having to assume an {\\it ad hoc} amount of convective core overshooting: a diffusive isochrone for 3.7 Gyr does a remarkably good job of matching the cluster observations.  The other significant result of this investigation, as far as isochrones are concerned, is that a gap near the turnoff is predicted to persist in open clusters up to an age of $\\approx 5.7$ Gyr by the diffusive models, whereas the limiting age is closer to 4.8 Gyr if diffusion is not treated.  It would be important to have detailed observations of the fiducial sequences for such clusters as those identified by \\citet{FrielJaTaetal2002} to test this prediction.  Unfortunately, NGC$\\,$188 appears to be too old to do this, given that our best estimate of its age is 6.4 Gyr based on the diffusive isochrones (which, incidently, provide a superb match to the observed CMD).  The predicted surface abundance variations among near turnoff stars turn out to be limited to approximately 0.1 dex in M\\,67 and 0.07 dex in NGC\\,188 (see \\S \\ref{sec:composition}). Most elements heavier than Si have their surface abundances modified by \\gr{} but no large overabundances are expected.  While not negligible, such variations are not easy to detect at the present time.  The existing Li/Be measurements in a few stars of M\\,67  (see \\S \\ref{sec:li_be}) suggest that another process may be required to reduce the Li abundance (see \\citealt{SillsDe2000}) though the Be/Li trend obtained with models including all aspects of atomic diffusion is very different from the trend obtained by these authors. With improved observations this  becomes a test of various turbulent models and will be further discussed in a paper in preparation on LiBeB abundances in cluster and field \\MS{} stars."
    },
    "0402/astro-ph0402634_arXiv.txt": {
        "abstract": "We present a method from an X-ray observation of a galaxy cluster to measure the radial profile of the dark matter velocity dispersion, $\\sigma_{\\rm DM}$, and to compare the dark matter ``temperature'' defined as $\\mu m_{\\rm p} \\sigma_{\\rm DM}^2 / k_{\\rm B}$ with the gas temperature. The method is applied to the XMM-Newton observation of Abell 1795. The ratio between the specific energy of the dark matter and that of the intracluster medium (ICM), which can be defined as $\\beta_{\\rm DM}$ in analogy with $\\beta_{\\rm spec}$, is found to be less than unity everywhere ranging $\\sim 0.3-0.8$. In other words, the ICM temperature is higher than the dark matter ``temperature'', even in the central region where the radiative cooling time is short. A $\\beta_{\\rm DM}$ value smaller than unity can most naturally be explained by heating of the ICM. The excess energy of ICM is estimated to be $\\sim1-3$~keV per particle. ",
        "introduction": "Early X-ray imaging observations with the {\\it Einstein} observatory and {\\it ROSAT} showed that in the central regions of clusters of galaxies the radiative cooling time is shorter than the age of the universe (e.g Canizares, Stewart, \\& Fabian 1983). As a result, the intracluster medium (ICM) should cool down to form a cold ($T<10^6$K) gas phase inducing a global inflow of gas. This ``cooling flow'' (see Fabian 1994 for a review) picture has been extensively discussed and formed a basic assumption in many arguments. The low resolution spectroscopy in 0.5-2 keV by {\\it ROSAT} showed that in some clusters the ICM temperature actually decreases towards the center (e.g. B\\\"{o}hringer et al. 1994; David et al. 1994; Allen \\& Fabian 1994). Higher resolution spectroscopy in 0.5-10 keV with {\\it ASCA}, however, can not be fully understood with the conventional cooling flow model. {\\it ASCA} spectra of cooling flow clusters can be well explained by a two (hot and cool) phase plasma without significant excess absorption features (e.g. Ikebe et al. 1999; Makishima et al. 2001). A naive cooling flow model predicting a range of temperatures with intrinsic absorption could also fit the {\\it ASCA} data but generally produce worse chi-square results (e.g. Allen et al. 2001). Most recently, very high resolution spectroscopy with XMM-Newton/RGS unambiguously show that there is very little X-ray emission from gas cooler than certain lower cut-off temperatures of $\\sim 1-3$~keV (Tamura et al. 2001; Kaastra et al. 2001; Peterson et al. 2001). Unless a large amount of cooled gas or the metals in the cold gas are hidden (Fabian et al. 2001a), there must exist a heating mechanism that prevents the ICM from radiative cooling.   The necessity of global heat input into the ICM in addition to gravitational heating has also been pointed out from the break of the self-similarity between dark matter and ICM, which is most clearly demonstrated in the X-ray luminosity-temperature relation ($L-T$ relation). A simple scaling analysis (Kaiser 1986) suggests a relation of $L\\propto T^2$, while observation shows $L\\propto T^{3}$ (e.g. Edge \\& Stewart 1991; David et al. 1993; White et al. 1997; Wu et al. 1999). The break of the self-similarity is also seen in the entropy vs temperature relation. Ponman, Cannon, \\& Navarro (1999) showed that cooler systems ($T<4$ keV) have entropies higher than achievable through gravitational collapse alone. Simulation studies show that the observed relations can be reproduced, if there is enough non-gravitational heat input into the ICM by feedback from galaxies (e.g. Metzler \\& Evrard 1994; Bower et al. 2001) or preheating before the cluster formation (e.g. Navarro et al. 1995; Tozzi \\& Norman 2000).   In order to shed some new light onto these ``cooling flow phenomena'' and ``break of self similarity'', we, in the present paper, perform a comparison of the temperature distribution of the ICM to the distribution of the velocity dispersion of the dark matter. A parameter, $\\beta_{\\rm spec}\\equiv\\sigma_{gal}^2/(k_{\\rm B}T/\\mu m_p)$, is often used as a measure of the average kinetic energy per unit mass in galaxies relative to that in the ICM. From observations of many clusters, the mean $\\beta_{\\rm spec}$ is $\\sim 1$ with large scatter (e.g. Wu et al. 1999), indicating that the energy equipartition between galaxies and ICM is roughly achieved on average. In analogy with $\\beta_{\\rm spec}$, we can introduce $\\beta_{\\rm DM}\\equiv\\sigma_{DM}^2/(k_{\\rm B}T/\\mu m_p)$ for comparison between the mean kinetic energy of the dark matter and that of the ICM, and define the dark matter ``temperature'' as $T_{\\rm DM}\\equiv\\mu m_{\\rm p} \\sigma_{\\rm DM}^2 / k_{\\rm B}$. Note that the above definition of ``temperature'', by means of proton mass instead of the actual mass of dark matter particles, is only for the sake of comparison with the gas temperature. Therefore we put ``temperature'' in quotes. We obtain, in this paper, the radial profile of the $\\beta_{\\rm DM}$ value observationally for the first time.     In Sect. 2, we describe the method of measuring the dark matter velocity dispersion in a cluster of galaxies from an X-ray observation. We applied the method to the XMM-Newton data of a prototypical cooling flow cluster, Abell~1795 (hereafter A1795), which is located at z=0.0616 (Struble \\& Rood 1987). The X-ray data analysis and results are presented in Sect. 3. Discussion and summary are found in Sect. 4 and Sect. 5, respectively. Throughout the paper, the Hubble constant is given as 70 $h_{70}$ km s$^{-1}$ Mpc$^{-1}$, and a flat universe ($\\Omega_{\\rm m,0}=0.3$, $\\Omega_{\\rm \\Lambda,0}=0.7$) is assumed. At the redshift of A1795, 1 arcsec corresponds to 1.19 kpc. ",
        "conclusions": "\\subsection{Anisotropy of $\\sigma_{\\rm DM}$}  Equation~(\\ref{eq:Jeans}) that we have used to calculate the dark matter velocity dispersion profile is based on the assumption of isotropic motion of dark matter particles. However, simulation studies (e.g.  Eke, Navarro, and Frenk 1998; Col\\'{i}n, Klypin, \\& Kravtsov 2000) indicate that the radial velocity dispersion, $\\sigma_{\\rm r}^2$, should be rather larger than the tangential velocity dispersion, $\\sigma_{\\rm t}^2 \\equiv \\frac{1}{2}(\\sigma_{\\theta}^2+\\sigma_{\\phi}^2)$. The degree of anisotropy is often measured with \\begin{equation} A \\equiv 1 - \\frac{\\sigma_{\\rm t}^2}{\\sigma_{\\rm r}^2} \\ , \\label{eq:anisotropy} \\end{equation} and the Jeans equation is modified as \\begin{equation} \\frac{GM}{R} = - \\sigma_{\\rm DM}^2 \\left( \\frac{d \\ln{\\rho_{\\rm DM}}}{d \\ln{R}} + \\frac{d \\ln{\\sigma_{\\rm DM}^2}}{d \\ln{R}} + 2 A \\right) \\ . \\label{eq:Jeans2} \\end{equation}   Employing $A = 0.65 \\frac{4 R/R_{\\rm vir}}{(R/R_{\\rm vir})^2 + 4}$ derived by Col\\'{i}n, Klypin, \\& Kravtsov (2000), which is as large as 0.5 at the virial radius, $R_{\\rm vir}$, and converging to 0 at the center, we solved Eq. \\ref{eq:Jeans2} as done in Sect. 4.6. The $\\sigma_{\\rm DM}$ profile thus derived for the case of the best fit NFW mass profile is shown in Fig.~\\ref{sigma_dm2}, overlaid with the $A=0$ solution given in Fig.~\\ref{sigma_dm}. The $\\sigma^2$ value with the anisotropy is greater than that without the anisotropy by only $\\sim$6\\% at 100kpc and at most $\\sim$40\\% around 1Mpc.   \\placefigure{concern}  \\placefigure{sigma_dm2}   \\subsection{Heating source}  The $\\beta_{\\rm DM}$ profile determined here from observations should provide information on the thermal history of the ICM. From numerical simulation studies, $\\beta_{\\rm DM}\\sim 1-1.4$ is expected, if there is no cooling or additional heating (e.g. Metzler and Evrard 1994; Navarro et al. 1995; Bryan \\& Norman 1998). An obvious way to explain the $\\beta_{\\rm DM}$ value smaller than unity in A1795 is heating of the ICM.   As suggested from the break of the self-similarity between dark matter and ICM, there should have been non-gravitational heating acting globally. We, from our results, estimated the excess energy of the ICM over that of the dark matter particles as \\begin{equation} \\Delta E (<R) = \\frac{3}{2} \\left( <kT> - \\mu m_{\\rm p} <\\sigma_{\\rm DM}^2> \\right) \\ , \\end{equation} where $<>$ denotes mass weighted mean within radius, $R$. As shown in fig. \\ref{mean_dE_pro}, the excess energy, $\\Delta E$, thus derived is found to be $\\sim 1-3$ keV per particle, which may be compared with theoretical model predictions. The amount of energy injection to the gas phase that explains e.g. the observed $L-T$ relation depends on the period when the heating occurred. Heating prior to cluster collapse, ``preheating'', needs $0.1-0.3$~keV per particle (e.g. Navarro et al. 1995; Tozzi \\& Norman 2001), while heating after a cluster formation requires higher values of $1-3$~keV per particle (e.g. Metzler \\& Evrard 1994; Loewenstein 2000; Wu, Fabian, \\& Nulsen 2000; Bower et al. 2001). Our results given above may indicate that the global non-gravitational heating that may cause the break of self similarity has happened mainly within a collapsed cluster. However, even if such non-gravitational heating that explains the global X-ray feature of the cluster is provided, the central region of A1795 has a short radiative cooling time (Fig.~\\ref{coolt}), and there must be another significant energy input to the ICM in the central region at the present epoch to prevent the ICM from cooling.   \\placefigure{mean_beta_pro}   \\placefigure{mean_dE_pro}     As a possible energy source in the central region, we first consider gravitational energy of the member galaxies and stellar components therein. The kinetic energy of the random motion of stars can be partially transferred to the ICM by stellar mass loss. Gas supplied by stellar mass loss has velocities of the bulk motions of stars relative to the ICM, which is the sum of a galaxy motion and the motions of the stars in the galaxy, and is likely to be thermalized by interactions with the ambient gas. If the stellar component moving in the same gravitational potential has a similar velocity dispersion profile as that of the dark matter (Fig.~\\ref{dm_tpro}), the gas temperature achieved from this process is expected to be comparable to the velocity dispersion of the dark matter. This process nicely accounts for X-ray emission from isolated X-ray compact elliptical galaxies (Matsushita 2001). The input rate of the kinetic energy of the gas from stellar mass loss may be simply estimated as $\\dot{E}=1/2 M_{\\rm star} \\dot{m} \\sigma^2$ = $10^{42}$ ergs/s, where $M_{\\rm star}$ is the total stellar mass of $1 \\times 10^{12}$ M$_{\\odot}$, $\\dot{m}$=$3\\times10^{11}$ M$_{\\odot}$ yr$^{-1}$ ($10^{11}$M$_{\\odot}$)$^{-1}$ is the stellar mass loss rate in unit time, and $\\sigma$(=580$\\pm$80 km/s) is the velocity dispersion of the galaxies from den Hartog \\& Katgert (1996; Fig. \\ref{dm_tpro}). (The stellar velocity dispersions in the galaxies are smaller and neglected here.) This is much smaller than the output energy in the central region by X-ray radiation in galaxy clusters like A1795, however. Thus for the case of galaxy clusters, we have to look for additional heat sources.   Kinetic energy of the stellar component might be more efficiently transferred to ICM via magnetic fields. As pointed out by e.g. Makishima et al. (2001), the motion of stars may amplify interstellar magnetic fields and reconnections of the fields may heat up the ICM rather efficiently. The galaxies must have lost their kinetic energies through interactions with the ICM and have gradually fallen inwards accumulating onto the central galaxy to form the cD galaxy. A deep optical image of the cD galaxy in A1795 derived by Johnstone et al. (1991) shows a concentration of elliptical galaxies of various sizes and stars forming a largely extended envelope with 131 $h_{70}^{-1}$~kpc effective radius, which strongly suggests the on-going formation process of the cD galaxy. Quantitatively, the total amount of dynamical energy of the stellar component in the member galaxies that has been lost in the past is estimated. The stellar component in the galaxies is assumed to have a negligible small potential ($U$) and kinetic energy ($K$) before the formation of the cluster, and the current energy of the stars is estimated to be $U+K\\sim - 10^{62}$ ergs. If the energy has been released over the last 10 Gyr and has been spent in ICM heating, the heating luminosity is expected to be $\\sim3\\times10^{44}$ ergs s$^{-1}$. This amounts to the bolometric luminosity of the ICM within 60$h_{70}^{-1}$~kpc (Fig. \\ref{coolt}), and may be sufficient to sustain the thermal energy of the ICM against radiative cooling. This model predicts that the stellar velocity dispersion became smaller than that of ICM, i.e. $\\beta_{\\rm spec} < 1$, which is consistent with the actual observed value in the central region (Fig. \\ref{dm_tpro}).   Alternatively, there may be sufficient non-gravitational heat input from an active galactic nucleus (AGN) powered by an accretion of low entropy gas at a cluster center (Churazov et al. 2002; B\\\"{o}hringer et al. 2002). Numerical simulations show that an outflow from an AGN may form hot bubbles of relativistic plasma rising with buoyancy, and the bubbles may uplift cold gas mixing with ambient ICM (Churazov et al. 2001; Quilis, Bower, \\& Balogh 2001; Br\\\"{u}ggen \\& Kaiser 2002; Basson \\& Alexander 2003). High resolution imaging observations revealed ripples and shock features in the central region of Perseus cluster (Fabian et al. 2003) and M87 (Forman et al. 2004), suggesting that the bubble energy may also be transfered by sound wave to larger distances, and that shocks may be the major contribution of the energy dissipation. The heating mechanism is self regulated: the lower the entropy, the higher the accretion rate onto the central engine. A portion of the accretion power is dissipated back into the ICM to make its entropy high and regulate the accretion rate to achieve an equilibrium state. This process automatically prevents the persistence of cold and hence dense clouds. The cD galaxy of A1795 has a radio source, 4C26.42, and the existence of an AGN is evident. Using the physical state of the ICM in the center we can actually estimate the energy provided by the AGN by applying the Bondi accretion model. According to the well-known correlation of the black hole mass with the mass of the bulge component (Magorrian et al. 1998), the black hole mass is expected to be $\\sim6\\times 10^9$ M$_{\\odot}$. Assuming that the gas profile is flat in the center, we can use the measured values of $n_{\\rm g}$=0.1 cm$^{-3}$ and $T_0$=2.8 keV to obtain the Bondi mass accretion rate \\begin{eqnarray} \\dot{M} & = & 4\\pi 0.25 \\rho_{\\infty} c_{s,\\infty}^{-3} (G M_{\\rm BH})^2 \\\\ & = & 0.23 M_{\\odot}/{\\rm yr} \\left( \\frac{n}{1 {\\rm cm}^{-3}} \\right) \\left( \\frac{T}{1 {\\rm keV}} \\right)^{-3/2} \\left( \\frac{M_{\\rm BH}}{6\\times 10^9 M_{\\odot}} \\right)^2 \\ , \\end{eqnarray} where $\\rho_{\\infty}$ and $c_{s,\\infty}$ are the density and sound velocity outside the Bondi accretion radius. We find $\\dot{M} \\sim 0.4 \\times 10^{-2} M_{\\odot}/{\\rm yr}$. Under the standard assumption of 10\\% of the accretion energy to be dissipated, the output energy is found to be $E = 0.1 \\dot{M} c^2 \\sim 3 \\times 10^{43}$ ergs/s. This amounts to the X-ray luminosity within the inner $\\sim20$ kpc region only. The Bondi accretion radius is estimated to be $R_{\\rm B} \\sim G M_{\\rm BH} / c_{\\rm s}^2 \\sim$ 30 pc, much smaller than the resolution of the temperature and density structure that can be measured with XMM-Newton. If the ICM density is not uniform but is clumpy, the Bondi accretion rate should be significantly larger and the heating rate could also be larger than the above estimation. We can note that the ICM conditions might temporally vary and that we currently see a relatively low state.   Note that there are other heating mechanisms discussed, which include e.g. a classical idea of thermal conduction from the hot outer regions (e.g. Takahara \\& Takahara 1979; Tucker \\& Rosner 1983; Bregman, \\& David 1988; Gaetz 1989), and the acoustic wave heating recently proposed by Fujita, Suzuki, \\& Wada (2004). Comparison of those model predictions with our results presented in the current paper would be very interesting.   \\placefigure{coolt}"
    },
    "0402/astro-ph0402402_arXiv.txt": {
        "abstract": "We discuss the polarization properties and first-order diffraction efficiencies of volume phase holographic (VPH) transmission gratings, which can be exploited to improve the throughput of modern spectrographs.  The wavelength of peak efficiency can be tuned by adjustment of the incidence angle. We show that the variation of the Kogelnik efficiency versus Bragg angle depends only on one parameter, given by $P_{\\rm tune} = (\\Delta n \\, d)/(n \\, \\Lambda)$, where: $\\Delta n$ is semi-amplitude of the refractive index modulation; $n$ is the average index; $d$ is the thickness of the active layer; and $\\Lambda$ is the grating period.  The efficiency has a well defined dependence on polarization. In particular, it is possible to obtain theoretical 100\\% diffraction efficiency with one linear polarization at any angle or to obtain 100\\% efficiency with unpolarized light at specific angles.  In the latter case, high efficiency is the result of aligning the peaks of the $s$- and $p$-polarization efficiency-versus-thickness curves.  The first of these `\\spphased\\ gratings' for astronomy is in use with the 6dF spectrograph.  Consideration of polarization is particularly important for high spectral resolution, which requires large incidence angles.  We also discuss the possibility of separating polarization states for improved throughput along the entire optical train of a spectrograph. ",
        "introduction": "\\label{sec:intro}  Astronomical spectrographs have undergone a major revolution during the past few decades \\citep{vBBH00,IM00,IM03,LM02,AD03}.  The revolution has concentrated on the multiplex advantage in order to allow large numbers of objects or contiguous spatial elements to be observed simultaneously.  This is possible because large detectors are now available, which can also lead to wide angle fields and/or wide spectral coverage.  Even though modern spectrographs can achieve up to 40\\% throughput (optics$+$disperser$+$detector), instrumental throughput remains a key issue for spectrograph design. Moreover, some fraction of the light lost along the optical train is received as stray light at the detector, and usually provides a major source of systematic error in the detected signal.  Now that detectors are widely available with 90\\% quantum efficiency in the visible wavelength region, the remaining gains must come from more efficient designs of the optics and dispersing element or elements, which we discuss. We concentrate specifically on volume phase holographic (VPH) gratings used in transmission \\citep{Arns95}. However, we note that similar consideration could apply to a much wider class of dispersing elements, e.g., reflection gratings, prisms.  What has not been discussed widely is the advantages of the polarization properties of VPH gratings, in particular, achieving the ideal of 100\\% throughput at any diffraction angle in one linear polarization. In addition, a theoretical diffraction efficiency of 100\\%, in both polarizations, can be achieved at specific angles with particular instrument configurations. An instrument that exploits the advantages of VPH gratings can, in principle, greatly reduce systematic error in the detected signal.  VPH gratings are already in use or are being brought into use in a number of spectrographs, including: LDSS++ and Taurus at the Anglo-Australian Telescope \\citep{glazebrook98jan}; OSIRIS at the Gran Telescopio Canarias \\citep{cepa00}; Goodman spectrograph at the SOAR Telescope \\citep*{CES00}; M2HES at the Magellan~II Telescope \\citep{bernstein02}; FORS at the Very Large Telescope \\citep*{MDR02}; FOCAS at the Subaru Telescope \\citep{ebizuka03}; and LRS at the Hobby-Eberly Telescope \\citep{hill03}. The potential of VPH gratings for astronomical applications has been investigated and discussed by Barden and others (\\citealt*{BAC98}; \\citealt{barden00,barden00proc,barden02}; \\citealt*{BCY03}; \\citealt{robertson00}; \\citealt*{RRD03}; \\citealt{tamura04}).  Here, we elucidate the physics of VPH gratings with emphasis on their polarization and tuning properties, and we discuss how these properties might be exploited to improve the performance of spectrographs. In \\S~\\ref{sec:intro-vph}, we describe the physics of VPH transmission gratings; in \\S~\\ref{sec:example-cases}, we describe some applications taking advantage of the well-defined polarization properties; in \\S~\\ref{sec:summary}, we summarize; and in the Appendix, we give equations for calculating the resolving powers of transmission gratings immersed between prisms. ",
        "conclusions": "\\label{sec:example-cases}  \\subsection{\\spphased\\ gratings}  In most spectrographs, the polarization states are not separated and therefore the efficiency with unpolarized light is important. For low resolution spectrographs ($\\alpha_{2b}\\la 20\\degr$), the theoretical diffraction efficiency can be above 95\\% with standard VPH gratings. At higher resolution, the efficiency of standard gratings can be significantly lower. Instead, \\spphased\\ gratings can be used to obtain higher efficiency at specific angles (\\S\\S~\\ref{sec:eff-1st-ord}--\\ref{sec:tuning}).  The first \\spphased\\ (Dickson) grating for astronomy was manufactured by Ralcon\\footnote{Ralcon Development Lab., P.O.~Box~142, Paradise, Utah 84328 ({\\tt http://www.xmission.com/$\\sim$ralcon/}), founded by R.~D.\\ Rallison.} for the 6dF multi-object spectrograph at the UK Schmidt Telescope \\citep{saunders01}.  It was specially designed to observe the Calcium triplet around 0.85\\micr\\ with a resolving power of about 8000 (first order, 1700 \\lmm).  The central wavelength is diffracted with a total beam deviation in air of 94\\degr, which makes use of the 35\\degr\\ special Bragg angle.  The theoretical diffraction efficiency is above 95\\% in the range 0.835--0.865\\micr\\ and the performance is near to that.  In addition, the camera and collimator are close to the grating.  Figure~\\ref{fig:6dF} shows the 6dF spectrograph in this configuration. This has significantly improved the PSF at the detector in comparison with using reflection gratings that have similar resolving powers (W.\\ Saunders 2003, private communication).  \\begin{figure} \\epsscale{1.15} \\plotone{f7.ps} \\caption{The 6dF bench-mounted spectrograph with the 1700 \\lmm\\ \\spphased\\ grating in place. There is a 90--100\\degr\\ beam deviation between the last element of the collimator (left) and the camera (right).} \\label{fig:6dF} \\end{figure}  If we wish to go to higher spectral resolution but are limited to a certain maximum deviation between collimator and camera beams, then prisms can be attached (Figs.~\\ref{fig:res-beam-dev}--\\ref{fig:prism-model}, Appendix).  For example, a grating with 2400 \\lmm\\ and 20\\degr\\ prisms ($n_2=1.25$, $\\Delta n_2 \\, d = 0.73$, $n_1=1.5$) that operates at $\\lambda=0.85$\\micr\\ with a 111\\degr\\ total beam deviation in air ($\\alpha_0=35.5\\degr$, $\\alpha_1=43\\degr$) can make use of the 55\\degr\\ special Bragg angle for high efficiency (\\ptune=1.4). This type of grating is challenging to produce because of the difficulty in testing the efficiency prior to attaching prisms (consider total internal reflection) but if high resolution and high efficiency are important over a narrow wavelength range then it could be useful.  Testing is needed to determine the laser exposure levels for the DCG. One solution would be to design the grating to work at the 35\\degr\\ Bragg angle, which has the same \\ptune\\ value. This would require testing at a wavelength related by 0.708 [$=\\sin(35.3) / \\sin(54.7)$] times the design wavelength, subject to variations in $n_2$ and $\\Delta n_2$ with wavelength (Eqns.~\\ref{eqn:bragg} and~\\ref{eqn:p-tune}).  Note that in order for such gratings to reduce systematic noise in the detected signal, it may be necessary to use a filter to block scattered light from outside the desired wavelength range.  \\begin{figure} \\epsscale{1.15} \\plotone{f8.ps} \\caption{Resolving powers of VPH transmission gratings versus total beam deviation. The top line represents a grating immersed between two 40\\degr\\ prisms (with $n_1=1.5$), the dashed line between two 20\\degr\\ prisms and the lower line represents a grating with no prisms attached.  The crosses are set at 10\\degr\\ intervals in Bragg angle (with $n_2=1.3$).  The resolving powers are normalized to unity for the zero deviation 40\\degr\\ prism model. See Fig.~\\ref{fig:prism-model} for the prism model and the Appendix for the calculation of resolving power. Note that the dispersion caused by differential refraction is not included.} \\label{fig:res-beam-dev} \\end{figure}  \\begin{figure} \\epsscale{1.15} \\plotone{f9.eps} \\caption{Diagram of a prism model for an immersed transmission grating. The dependence of the spectral resolving power on the angles and indexes is given in the Appendix. In Littrow configuration, with both prism angles equal to $\\gamma$, the resolving power is approximately proportional to $n_1 \\tan\\alpha_1 \\cos(\\alpha_1-\\gamma) / \\cos\\alpha_0$.} \\label{fig:prism-model} \\end{figure}  \\subsection{Separating polarization states in a spectrograph}  If we could envisage an ideal spectrograph, what properties would it have? The primary problems are scattered light at refractive index boundaries, and the difficulties of dispersing \\s- and \\p-polarization states without compromise of one or the other. We note that both of these problems arise from the geometry of the wave front with respect to the optical element/grating.  A substantial increase in efficiency, perhaps approaching the ideal, could be achieved by allowing the two polarization states to be handled separately in a spectrograph. One approach would be to separate the two polarizations at a polarizing beam splitter \\citep{Goodrich91}. They would then propagate along separate paths where all the optics would be oriented to minimize light loss for that polarization. In particular, a VPH grating can be optimized for almost any angle to obtain near 100\\% efficiency at blaze wavelength in one polarization.  An alternative to the use of a beam splitter would be to use a VPH grating itself.  The collimated beam would encounter a VPH grating in the normal way, with diffracted light going to a camera. But the grating would be designed so that first-order diffraction was optimized for a single linear polarization while zeroth order was optimized for the other polarization.\\footnote{At certain Bragg angles, it is theoretically possible to achieve 100\\% diffraction efficiency in one linear polarization (\\s\\ or \\p), with 0\\% diffraction efficiency in the other \\citep{DRY94,Huang94}. This is analogous to the special angles for \\spphased\\ gratings.} The undiffracted transmitted beam could then go on to a second VPH grating optimized for the other polarization, feeding into a second camera.  Whatever method is used to separate the polarizations, note that it is not necessary to achieve complete separation, and hence the requirement for the polarizing element is less demanding than for a polarimeter.  This is because the only effect of mixing in a small amount of light with the `wrong' polarization is that it will be less efficiently processed downstream. But it will still make some positive contribution to the signal, and the decrease in efficiency will be a second-order effect.  The cases where separate polarizations would most improve efficiency is with high spectral resolution. Large beam deviations, and large air-to-glass incidence angles, are needed to obtain the highest resolutions (Fig.~\\ref{fig:res-beam-dev}). At these large angles, the boundary-transmission efficiency of \\s-polarization is low but the efficiency in \\p-polarization remains high (Fig.~\\ref{fig:air-glass}).  For example, one could design a grating that operates with light incident near Brewster's angle on the air-glass boundary and with the DCG layer optimized for $p$-polarization efficiency.  This would provide a high-resolution grating with near 100\\% efficiency in one linear polarization. Naturally, polarimetry measurements could also take advantage of these one polarization optimized gratings.  Note also that separating polarization states could be used to optimize efficiency for reflection grating spectrographs \\citep{LA00} and for background-limited observing during bright Moon phases \\citep{BB01}.  We are moving toward precision measurements in many areas of astrophysics. The major limitation continues to be systematic sources of noise.  In order to combat this, many experiments are cast as differential measurements, e.g., alternating observations of source and background, in order to beat down the systematic errors. This is a highly effective strategy for dealing with external noise sources and some internal sources (e.g., apparatus instability).  However, there are internal sources of noise, which continue to haunt most spectrographs today, in particular, scattered light. This `ghost light' is not usually suppressed in differential experiments because it depends on the distribution of light sources over the field of view. Even with mitigation strategies based on light baffles and optimized AR coatings, there is always residual stray light, not least from the optical/IR detector because of its large refractive index compared to a vacuum.  But this situation is slowly improving as detectors and matched coatings approach their theoretical maximum.  The only way to guard against stray light is to consider the role of every element in the optical train very carefully and to orient the optical elements accordingly, particularly the choice of AR coating, and orientation of the interface to the incoming wave front. This is easier to do if the wavefront has been divided into its \\s- and \\p-states and each polarization is considered separately.  One only has to consider the AR coatings in each of two arms, one for \\s- and the other for \\p-states, which can be optimized for throughput.  This is not true for a skew ray in natural light, which would require a birefringent coating in order to optimize throughput in both polarization states.    VPH gratings are used in an increasing number of spectrographs because of their high diffraction efficiency. In this paper, we have outlined the basic physics necessary to design VPH gratings. In particular: we have defined a parameter, \\ptune, that determines how the efficiency of a grating varies with Bragg angle; we have described the possibility of creating \\spphased\\ gratings that can have 100\\% efficiency with unpolarized light at specific angles; and we have discussed the importance of considering the separate polarization states. The main points concerning tuning and efficiency are given below. \\begin{enumerate} \\item The grating period ($\\Lambda$) and the average refractive index of the DCG layer ($n_2$) determine the wavelength as a function of Bragg angle (Eqn.~\\ref{eqn:bragg} with $m = 1$). \\item The parameter \\ptune\\ of a grating (Eqn.~\\ref{eqn:p-tune}) determines how the efficiency varies with Bragg angle (Eqn.~\\ref{eqn:kogelnik2}). Standard grating designs have $P_{\\rm tune} \\la 0.5$ (Fig.~\\ref{fig:eff-bragg}). \\item \\spphased\\ gratings can be created by aligning the peaks of the $s$- and $p$-efficiency curves versus DCG thickness at particular Bragg angles (Fig.~\\ref{fig:eff-thick}).  For example, a grating with $P_{\\rm tune} \\approx 1.4$ has high efficiency with unpolarized light at a Bragg angle of 35\\degr. Here, the second peak of the $s$-curve is aligned with the first peak of the $p$-curve. \\item Bragg-condition diffraction efficiencies are lower than predicted by Eqn.~\\ref{eqn:kogelnik} or~\\ref{eqn:kogelnik2} if Kogelnik's condition is not satisfied (Eqn.~\\ref{eqn:kog-approx} or~\\ref{eqn:kog-approx2}). However, the efficiency can still be above 90\\% as long as $\\rho\\ga 3$ (the lost power is approximately $1/\\rho^2$). \\item The FWHM of the efficiency curve is approximately inversely proportional to the thickness of the grating (Eqn.~\\ref{eqn:bandwidth}, Fig.~\\ref{fig:bandwidth}). Therefore, it is generally optimal to have the thinnest possible DCG layer subject to manufacturing limitations and lost power from first-order diffraction. \\end{enumerate}  We have shown how VPH gratings can be manufactured and exploited to ensure higher transmission and better suppression of stray light. This will necessarily force instrument designers into a smaller parameter space, but we feel that there is sufficient freedom within that space to account for most design issues. In any event, we deem these considerations to be paramount if systematic sources of noise are ever to be effectively removed from the apparatus."
    },
    "0402/astro-ph0402128_arXiv.txt": {
        "abstract": "We present new observations of the gravitational lens system CLASS B0128+437 made in the optical, infrared and radio regimes. $Hubble Space Telescope$ observations detect only a very faint, extended object in $I$-band with no obvious emission from the lensed images visible; no detection at all is made in $V$-band. The lens system is detected with much higher signal to noise with UKIRT in $K$-band and, although resolved, the resolution is not sufficient to allow the lensed images and the lens galaxy to be separated. A careful astrometric calibration, however, suggests that the peak of the infrared emission corresponds to the two merging images A and B and therefore that the lensed images dominate at infrared wavelengths. The new radio data consist of high resolution VLBI radio images at three frequencies, 2.3, 5 and 8.4~GHz, made with the VLBA and the 100-m Effelsberg telescope. These reveal that the lensed source consists of three well-defined sub-components that are embedded in a more extended jet. Due to the fact that the sub-components have different spectral indices it is possible to determine, unambiguously, which part of each image corresponds to the same source sub-component. Our main finding is that one of the images, B, looks very different to the others, there being no obvious division into separate sub-components and the image being apparently both broader and smoother. This is a consequence we believe of scatter-broadening in the ISM of the lensing galaxy. The large number of multiply-imaged source sub-components also provide an abundance of modelling constraints and we have attempted to fit an SIE+external shear model to the data, as well as utilising the novel method of \\citeauthor{evans03}. It proves difficult in both cases, however, to obtain a satisfactory fit which strongly suggests the presence of sub-structure in the mass distribution of the lensing galaxy, perhaps of the kind that is predicted by CDM theories of structure formation. ",
        "introduction": "The Cosmic Lens All Sky Survey (CLASS), a radio survey of the northern sky, has been exceptionally successful in discovering new arcsec-scale gravitational lens systems, with a final tally of 22 \\citep{myers03,browne03}. These systems can, using the method of \\citet{refsdal64}, be used individually to determine the Hubble parameter \\citep[e.g][]{kochanek03,burud02,treu02,fassnacht02} and together to place constraints on the cosmological parameters $\\Omega_0$ and $\\lambda_0$ \\citep[e.g][]{helbig99,chae02}. As well as this, gravitational lenses are used to explore various aspects of high-redshift galaxies i.e. the lens galaxies themselves \\citep[e.g.][]{kochanek00,lehar00}, including their distribution of ionised and magnetised plasma, their radial mass profiles \\citep[e.g][]{koopmans03b,winn03a,biggs03,rusin02,rusin01,munoz01} and the presence of massive substructures such as those predicted by CDM theories of large-scale structure formation \\citep[e.g][]{mao98,bradac02,metcalf01,metcalf02a,metcalf02b,dalal02}.  B0128+437 \\citep{phillips00} was discovered by CLASS and consists of four images with a maximum separation of 540~mas (Fig.~\\ref{optical}). Initial modelling, in the absence of any information regarding the lensing galaxy, demonstrated that a singular isothermal ellipsoidal mass distribution plus a large external shear produces a good fit to the observed image positions and flux ratios. Subsequent optical spectroscopy with the W.~M.~Keck telescope \\citep{mckean04} measured a redshift for the lensed source of $z_s=3.12$ and identified a single line at a different redshift that is presumably that of the lensing galaxy. The most likely identification of this line (\\mbox{[O\\,{\\sc ii}]}) gives a lens redshift of $z_d=1.145$ although additional observations will be required to confirm this. The radio spectrum of the lensed source has a pronounced peak at around 1~GHz, thus making it a member of the Giga-hertz Peaked Spectrum (GPS) class.  In this paper we present new optical, infrared and radio VLBI observations of B0128+437 which demonstrate that this system is far more interesting than the original 5-GHz MERLIN data suggested. In the next section we show images made using the $Hubble Space Telescope$ ($HST$) and the United Kingdom Infrared Telescope (UKIRT). We then go on to describe new multi-frequency observations (2.3, 5 and 8.4~GHz) made with the VLBA together with the 100-m Effelsberg telescope and present the resulting maps of each of the lensed images. The remainder of the paper seeks to explain the strange appearance of image B and the difficulties found in fitting a lens model to the observed VLBI substructure. ",
        "conclusions": "We have presented VLBI observations of the gravitational lens system CLASS B0128+437 that reveal extensive structure within each of the four images. These data illustrate the role that high-resolution VLBI observations play in the study of lens systems, such as providing a large number of model constraints and identifying inconsistencies between the observations and the lens model. Interesting astrophysical conclusions regarding the lens galaxy can be drawn that would otherwise be missed using data from smaller arrays such as MERLIN and the VLA. Our observations demonstrate that despite the large number of model constraints made available by the VLBI imaging, we find it difficult to get smooth mass models to fit the observed image positions, signifying additional unmodelled structure in the deflector. This may be another instance of lensing providing evidence to support the existence of mass substructure in CDM haloes. Improved infrared $HST$ observations will be taken which may improve our understanding of the deflecting mass, leading to more rigorous modelling.  The most striking feature of our new VLBI maps is the anomalous appearance of image B compared to the other images. The most likely explanation for this is that this image is being scatter-broadened in the ISM of the lensing galaxy, a process which may also be affecting the other images to a lesser extent. An important consequence of this effect concerns the identification of lens systems in CLASS. The CLASS methodology was to use VLBA 5-GHz observations, once candidate systems had been identified with the VLA and MERLIN, to see whether any could be rejected on the basis of surface brightness arguments i.e. weaker components more extended, or through inconsistent structures in the different images. In a two-image system particularly, were one of the images to be distorted in a manner similar to that seen in B0128+437, that system may have been rejected. We are currently re-observing 13 CLASS two-image lens candidates at 15~GHz with the VLBA in order to check whether their initial rejection was sound."
    },
    "0402/astro-ph0402364_arXiv.txt": {
        "abstract": "I calculate the specific angular momentum of mass accreted by a binary system embedded in the dense wind of a mass losing asymptotic giant branch star. The accretion flow is of the Bondi-Hoyle-Lyttleton type. For most of the relevant parameters space the flow is basically an isothermal high Mach number accretion flow. I find that when the orbital plane of the accreting binary system and the orbital plane of the triple system are not parallel to each other, the accreted mass onto one or two of the binary system components has high specific angular momentum. For a large fraction of triple-star systems, accretion disks will be formed around one or two of the stars in the binary system, provided that the mass ratio of the two stars in the accreting binary system is $\\gtrsim 0.5$. Such disks may blow jets which shape the descendant planetary nebula. The jets' axis will be almost parallel to the orbital plane of the triple-star system. One jet is blown outward relative to the wind, while the other jet pass near the mass losing star, and is more likely to be slowed down or deflected. I find that during the final asymptotic giant branch phase, when mass loss rate is very high, accretion disk may form for orbital separation between the accreting binary systems and the mass losing star of up to $\\sim 400-800 \\AU$. I discuss the implications for the shape of the descendant planetary nebula, and list several planetary nebulae which may have been shaped by an accreting binary star system, i.e., by a triple star system. ",
        "introduction": "Both theory (Soker 1990) and observations (Sahai \\& Trauger 1998) suggest that many, but not all, planetary nebulae (PNs) are shaped by two oppositelly ejected jets from the progenitor system (for a recent review see Soker 2004). If not well collimated, these jets are termed CFW, for collimated fast wind. In principle, the jets (or a CFW) can be blown by the post-asymptotic giant branch (AGB) progenitor itself, or by an accreting companion. Theoretical and observational considerations, e.g., the detection of collimated outflows emanating from AGB stars vicinities (Imai et al.\\ 2002, 2003; Hirano et al.\\ 2004; Sahai et al.\\ 2003; Vinkovic et al.\\ 2004), suggest that in most cases, or even all, a binary companion blows these jets (see Soker 2004). The jets, as with most other astrophysical objects, are thought to be blown when an accretion disk is formed around a compact object, and the accretion rate is high enough.  For an accretion disk to form, the specific angular momentum of the accreted mass $j_a$, must obey the condition $j_a > j_b$, where $j_b=(G M_b R_b)^{1/2}$ is the specific angular momentum of a particle in a Keplerian orbit at the equator of the accreting star of radius $R_b$ and mass $M_b$. The accretion flow is of the Bondi-Hoyle-Lyttleton type (Hoyle \\& Lyttleton 1939; Bondi \\& Hoyle 1944), where gas with an impact parameter $b \\lesssim R_{\\rm acc}=2 G M_b /v_r^2$ is accreted. The impact parameter is the distance of the accreted mass from the symmetry line of the flow, termed the accretion line, at infinity before reaching the shock wave, \\begin{eqnarray} R_{\\rm acc} = 10.6 \\AU \\left( \\frac {M_b}{0.6 M_\\odot} \\right) \\left( \\frac {v_r}{10 \\km \\s^{-1}} \\right)^{-2} , \\end{eqnarray} is the Bond-Hoyle accretion radius, and $v_r$ is the relative velocity between the gas and the accreting body. For the rest of the paper I will be dealing mostly with large orbital separations to the mass losing star $a_0$, such that the orbital velocity is low, and the relative velocity can be taken as the wind velocity $v_r \\simeq v_s$. The condition for a companion in a circular orbit accreting from the AGB wind (but not via Roche lobe overflow [RLOF]), can be written in the following form (Soker 2001) \\begin{eqnarray} 1< \\frac {j_a}{j_b} = 0.25 \\left( \\frac {\\eta}{0.2} \\right) \\left( \\frac {M_a+M_b}{1.2 M_\\odot} \\right)^{1/2} \\left( \\frac {M_b}{0.6 M_\\odot} \\right)^{3/2} \\left( \\frac {R_b}{1 R_\\odot} \\right)^{-1/2} \\left( \\frac {a_0}{100 \\AU} \\right)^{-3/2} \\left( \\frac {v_s}{10 \\km \\s^{-1}} \\right)^{-4} , \\end{eqnarray} where $M_a$ is the mass of the mass-losing star. Here $\\eta$ is a parameter indicating the reduction in the specific angular momentum of the accreted gas because of the increase in the cross section for accretion from the low-density side. Namely, gas parcels with larger impact parameter are accreted from the low-density side. Livio et al.\\ (1986; see also Ruffert 1999) find that for high Mach number flows  $\\eta \\sim 0.1$ and $\\eta \\sim 0.3$, for isothermal and adiabatic flows, respectively. The scaling of the masses are for a star about to leave the AGB, most relevant stars do it with a mass of $M_a \\simeq 0.6 M_\\odot$, and taking a companion of equal mass at that point. For the mass loss rate and velocity of the AGB wind, a short period of very high mass loss rate and low velocity is assumed (sometimes referred to as the superwind or final intensive wind [FIW]).  Another plausible condition for the formation of jets (or a CFW) is that the accretion rate should be above a certain limit $\\dot M_{\\rm crit}$, which I take as $10^{-7} M_\\odot \\yr^{-1}$ for accretion onto a main-sequence star and $10^{-8} M_\\odot \\yr^{-1}$ for accretion onto a WD (Soker \\& Rappaport 2000). The Bondi-Hoyle-Lyttleton accretion rate is $\\dot M_b \\simeq \\pi R_{\\rm acc} ^2 v_r \\rho$, where the density at the location of the accretor is $\\rho = \\vert \\dot M_a \\vert / (4 \\pi a_0^2 v_s)$. Substituting the relevant parameters during the final intensive wind, and taking $v_r \\simeq v_s$, the mass accretion rate is \\begin{eqnarray} \\dot M_b \\simeq 3 \\times 10^{-7} \\left( \\frac {M_b}{0.6 M_\\odot} \\right)^{2} \\left( \\frac {v_s}{10 \\km \\s^{-1}} \\right)^{-4} \\left( \\frac {a_0}{100 \\AU} \\right)^{-2} \\left( \\frac {\\vert \\dot M_a \\vert }{10^{-4} M_\\odot \\yr^{-1}} \\right) M_\\odot \\yr^{-1}. \\end{eqnarray} It is commonly assumed that $\\sim 10 \\%$ of the accreted mass is blown into the CFW (or jets; Soker \\& Rappaport 2000), having speeds of the order of the escape velocity from the companion (Livio 2000), $100-10^3 \\km \\s^{-1}$ and $10^3-10^4 \\km \\s^{-1}$, for a main-sequence and a WD companion, respectively.  The basics of the binary model for shaping PNs and related objects, e.g., $\\eta$ Carinae, are reviewed in Soker (2004, where more references are given), while the ejection of two jets by a wide binary companion, which is one of several routes for the shaping of PNs and which is a process directly relevant to the present paper, is studied in Soker (2001; see also Soker \\& Rappaport 2000). Here I notice the following. (1) By equation (2), accreting WD companions may form accretion disks at orbital separation of $a_0 \\lesssim 150 \\AU$, while main sequence stars must be much closer at $a_0 \\lesssim 40 \\AU$. (2) When the mass losing upper-AGB star lose mass at a very high rate of $\\dot M \\sim 10^{\\-4} M_\\odot \\yr^{-1}$, the condition on the angular momentum (equation 2) is more difficult to fulfill than the condition on the mass accretion rate. (3) A single companion star will blow jets perpendicular to the binary equatorial plane. (or the jets will precess around this direction). In the following sections I'll show that these three properties not necessarily hold for a binary system accreting from the wind of AGB stars. Readers interested in results and observational implications only, may skip to section 4. ",
        "conclusions": "\\subsection{Summary of Theoretical Calculations}  In the previous sections I calculated the specific angular momentum of mass accreted onto a binary system, with the flow structure as follows (section 2). The binary system accretes from the dense wind of a mass losing giant star, most relevant are AGB stars. The flow is such that it is of the Bondi-Hoyle-Lyttleton type, i.e., the orbital separation of the accreting binary system, $a_{12}$, is much smaller than the Bondi-Hoyle accretion radius of the system $R_{\\rm acc}$ (eq. 1; fig. 1). I also showed that for relevant parameters, the post-shock wind cools fast relative to the flow time, leading to the formation of a dense accretion flow behind the binary system: the accretion column. I made the assumptions given in equation (8) regarding the relations between the relevant length scales in the problem. Typical length-scales are (see figures 1 and 2): the radius of the accreting star (a WD or a main sequence star): $0.01 R_\\odot \\lesssim R_1 \\lesssim 1 R_\\odot$; the orbital separation of the stars in the accreting binary system: $10 R_\\odot \\simeq 0.05 \\AU \\lesssim a_{12} \\lesssim 1 AU$; the Bondi-Hoyle accretion radius of the binary system (eq. 1): $10 \\AU \\lesssim R_{\\rm acc} \\lesssim 30 AU$; the orbital separation with the mass losing star: $30 \\AU  \\lesssim a_0 \\lesssim 300 \\AU$. I examine the formation of an accretion disk by demanding that the specific angular momentum of the accreted mass be larger than that of a test particle orbiting the equator of the accreting star.  The main results of these calculations can be summarized as follows. \\newline (1) The accreted mass acquires angular momentum as a result of the orbital motion of the accreting star around the center of mass of the accreting binary system. When the orbital plane of the accreting binary system is parallel to the orbital plane of the triple star system ($\\theta=0$ in fig. 1), the specific angular momentum is too small to form an accretion disk for the assumed parameters, i.e., the condition for an accretion disk formation is almost impossible to fulfill (eq. 14). \\newline (2) When the orbital planes are perpendicular to each other, the condition for disk formation is given by equation (10). The accretion in this particular case is in a steady state, because the distance of the stars from the accretion column does not change. The more massive star is expected to accrete most, or even all, of the mass. The condition can be easily met by many systems for which the mass ratio is $q=M_{b2}/M_{b1} \\gtrsim 0.3$. \\newline (3) For an inclined orbit, $0^\\circ < \\theta < 90^\\circ$, the accretion occurs mainly when the accreting star is at its closest approach to the accretion column. The more massive star will accreted more mass, but because of the periodic variation in distance from the accretion column of both stars, the less massive star accretes mass as well. Assuming most of the mass is accreted indeed when the star is close to the accretion column, the condition for disk formation is given by equation (15). This can be met by many systems if $\\sin \\theta$ is not too small and $q \\gtrsim 0.5$. \\newline (4) For the calculations here assumption (8) was used. This assumption is that the radius of the accreting star is much smaller than the width of the accretion column at the binary location, which is smaller than orbital separation of the accreting binary system, which is much smaller than the accretion radius of the binary system (eq. 1), which is smaller than the orbital separation with the mass losing star. However, the basic physics can hold for a binary system very close to the mass losing star. In that case, the mass in the accretion column will have some angular momentum relative to the center of mass of the accreting binary system, as single stars do (the term $j_a$ in eq. 2). The binary system may even have some tidal effect on the mass losing giant star. The flow becomes very complicated in such a case.  With these and earlier results, the following conclusions can be drawn. \\newline (5) For accreting binary systems which fulfill condition (15), the constraint on disk formation becomes the mass accretion rate as given by equation (3). Because the mass accretion rate depends on total mass square, the accretion rate is higher than in a single-star case. Consider two equal mass main sequence stars of $M_{b1}=M_{b2}=1 M_\\odot$, and demanding that each star accretes at a rate of  $> 10^{-7} M_\\odot \\yr^{-1}$. For the mass loss rate and wind velocity as in equation (3), the constraint from the mass accretion rate becomes $a_0 \\lesssim 400 \\AU$. For accreting WD stars, i.e., at least one of the stars in the binary system is a white dwarf which accretes half of the mass, I take $M_{b1}=M_{b2}=0.6 M_\\odot$, and as in section 1 requires the accretion rate to be $> 10^{-8} M_\\odot \\yr^{-1}$. The constraint from the mass accretion rate becomes $a_0 \\lesssim 800 \\AU$. These numbers should be compared with the constraint on accreting single stellar companion (first section) of $a_0 \\lesssim 40 \\AU$ and $a_0 \\lesssim 150 \\AU$, for main sequence and WD single companions, respectively. In the later case the constraint is the accreted specific angular momentum. \\newline (6) The constraint on the mass accretion rate can be eased if we consider the nature of the Bondi-Hoyle-Lyttleton accretion flow. The flow along the accretion column is unstable, with large variation in density, hence in the mass accretion rate on short time scales (Cowie 1977; Soker 1991). This implies that even when the average accretion rate is lower than the critical value, on short time scales it can be higher, leading possibly to sporadic jet formation. \\newline (7) The formation of jets at large distance from the mass losing star will lead to departure from axisymmetry, even when jets are blown perpendicular to the triple-star orbital plane (Soker 2001). Here there is another effect. Because the accretion disk is almost perpendicular to the orbital plane, i.e., its axis, hence the jets if are formed, is close to the triple-star orbital plane and pointing from the the accreting binary system to the mass-losing star. This means that one jet expands toward lower density medium and expands almost undisturbed. The opposite jet, on the other hand, expands toward the mass losing star and encounters dense medium. This jet can be deflected and/or slowed down by the dense wind from the mass losing star. \\newline (8) In addition to the accreting binary system which can reside at a large orbital separation $a_0$, the mass losing giant star may have a close companion. That closer companion will cause axisymmetrical mass loss as well, but most likely with a different axis of symmetry. \\newline (9) When $q \\sim 1$, i.e., almost equal mass companions, both stars can blow jets, each stars on its turn. This may further complicate the structure of the nebula. \\newline (10) The allowed large orbital separation for jets formation may result in delayed jets (Soker 2001). When the orbital separation to the mass losing star is $a_0 \\sim 200-500 \\AU$, and the wind is slow $v_0 \\simeq 5-10 \\km \\s^{-1}$, the time required for the wind to flow from the AGB mass losing star to the accreting system is $\\sim 100-500 \\yr$. This is a non-negligible fraction of the evolution time during the super-wind phase. The jets (or a CFW) may be blown after the mass loss rate from the AGB star has been substantially reduced. Namely, while the post-AGB star and its wind may show a post-AGB age of 100-500 years, the jets may still be active! \\newline (11) The triple-star scenario adds to the many routes through which binary systems can shape PNs (Soker 2002). This strengthens my earlier request not to use phrases like `peculiar', `unique', `extraordinary', and `unusual', in describing the kinematics and structure of PNs. All known PNs with large departures from axi-symemtrical structures can be fitted into the binary, including triple, stellar model for the shaping of PNs.  \\subsection{Observational Consequences}  The discussion in the preceding subsection shows that PNs descendant from AGB stars which have two companions in a binary system that blow jets, will most likely have a significant departure from pure axisymmetrical structure. These systems may not have even a mirror symmetry, because the two jets expands into different media. The exact PN structure which result from jets blown by an accreting binary system must be derived from 3D numerical simulations. Here I only list some PNs which potentially were shaped by jets blown by an accreting binary-system companion to their AGB progenitor. Other types of systems may also lead to departure from both mirror symmetry and from axisymmetry, e.g., stochastic mass loss from the mass losing star together with an accreting companion which blows the jets, or an accreting close single star companion, together with a wider companion to the AGB star which sole role is to cause departure from axisymmetry (Soker \\& Hadar 2002). Therefore, I don't expect all of the PNs listed below to have triple star systems. Some, though, may have triple star systems. \\newline {\\bf He 3-1375.} This PN has a dense ring, and small lobes protrude from the nebula (Bobrowsky et al.\\ 1998). However, no single axis of symmetry, or point-symmetry, or a plane of symmetry, can be defined to this nebula. \\newline {\\bf IC 2149 (PN G 166.1+10.4).} This PN has extremely asymmetrical narrow structure (Balick 1987; Vazquez et al.\\ 2002). Vazquez et al.\\ (2002) suggest that this is the equatorial plane of the nebula. Another possibility is that the narrow structure is composed of unequal jets. \\newline {\\bf M 1-59 (PN G 023.9-02.3).} This bipolar PN has unequal-size lobes, both with departure from axisymmetry (Manchado et al.\\ 1996). A close accreting single companion which blew two jets, and a wider companion which sole role was to cause the departure from axisymmetry, may also account for the structure, although this model cannot by itself account for the unequal size of the two lobes. \\newline {\\bf NGC 6210 (PN G 043.1+37.7).} This is a `messy' PN, with a general elliptical structure with unequal sides, and blobs, filaments, and jet-like structures around it (Balick 1987; Terzian \\& Hajian 2000). \\newline {\\bf NGC 1514 (PN G 165.5-15.2)} This PN has a general axisymmetrical structure, but with large departure from exact axisymmetry revealed both in its image (Balick 1987) and in its kinematics (Muthu \\& Anandarao 2003). \\newline {\\bf NGC 6886(PN G 060.1-07.7).} Two cylindrical-type lobes protrude from a spherical structure. The two lobes are bent relative to the symmetry axis to the same side. Bobrowsky et al.\\ (2004) proposed that the lobes were formed by two jets (or a collimated fast wind [CFW]) blown by a companion at orbital separation of $\\sim 30 AU$. The jets were bent by the ram pressure of the wind from the mass losing AGB stellar progenitor of the PN. This model by itself, however, cannot account for the observation that the two lobes are unequal in size and intensity.   I thank an anonymous referee for helpful comments. This research was supported in part by grants from the Israel Science Foundation."
    },
    "0402/astro-ph0402478_arXiv.txt": {
        "abstract": "We have studied the dynamical evolution of rotating star clusters with mass spectrum using a Fokker-Planck code. As a simplest multi-mass model, we first investigated the two-component clusters. Rotation is found to accelerate the dynamical evolution through the transfer of angular momentum outward, as well as from the high masses to the low masses. However, the degree of acceleration depends sensitively on the assumed initial mass function since dynamical friction, which generates mass segregation, also tends to accelerate the evolution, and the combined effect of both is not linear or multiplicative. As long as dynamical friction dominates in the competition with angular momentum exchange the heavy masses lose random energy and angular momentum, sink towards the centre, but their remaining angular momentum is sufficient to speed them up rotationally. This is gravo-gyro instability. As a consequence, we find that the high mass stars in the central parts rotate faster than low mass stars. This leads to the suppression of mass segregation compared to the non-rotating clusters. From the study of multi-component models, we observe similar trends to the two-component models in almost all aspects. The mass function changes less drastically for clusters with rotation. Unlike non-rotating clusters, the mass function depends on $R$ and $z$. Our models are the only ones that can predict mass function and other quantities to be compared with new observations. ",
        "introduction": "\\medskip  Since the pioneering studies of Lupton \\& Gunn (1987) and its subsequent applications to fit actual globular cluster observations it is evident that galactic globular clusters exhibit some degree of rotation, and that there is also a consistent amount of observed flattening of their shapes (White \\& Shawl 1987). More recently, van Leeuwen et al. (2000) and Anderson \\& King (2003) measure globular cluster rotation in proper motion, in order to derive the true direction of its rotational axis. Though not dominant the amount of rotational energy could also not be neglected on the other hand side, as was seen by the models of Einsel \\& Spurzem, which showed that even a moderate fraction of rotational energy in a cluster leads to significantly faster core evolution.  This is the third in a series of studies on the dynamical evolution of rotating stellar system by using orbit-averaged 2D Fokker-Planck (FP) models that include the effects of initial rotation. In the previous two papers, pre-collapse evolution (Einsel \\& Spurzem 1999, Paper I) and the evolution after core-collapse (Kim et al. 2002, Paper II) of rotating clusters composed of equal mass stars were studied. In the present study we explore the dynamical evolution of the rotating stellar systems with mass spectrum as a natural extension to the previous works. We expect that the exchange of angular momentum between different mass species would significantly affect the course of dynamical evolution, as the energy exchange is known to play important role in the evolution of multi-mass clusters.   Direct integration of the Fokker-Planck equation is used as a statistical method. Comparisons between results obtained with FP method and results from $N$-body simulation show that the approximations and assumptions which were used in FP models are reasonable, but need to be checked carefully by comparison of different methods. The simplest implementation of an FP model is a one-dimensional (1D) FP model where the distribution function $f$ is assumed to depend only on the energy $E$ of stars. All physical properties depend on the distance from the cluster center only (Cohn 1980; Lee, Fahlman \\& Richer 1991). The use of 1D models is inspired by the fact that the shape of globular clusters is approximately spherical, but it ignores the anisotropy of the velocity dispersion. Observations of globular clusters and theoretical models suggested that there exists a difference in the velocity dispersion between radial and tangential direction, especially for stars in the outer regions of the clusters (Takahashi 1995, 1996, 1997). The two-dimensional FP model where distribution function depends on energy and angular momentum with spherical geometry has been pioneered by Cohn (1979) but studied extensively only recently because of difficulties in numerical integration of the equations (Takahashi, Lee \\& Inagaki 1997, Takahashi \\& Lee 2001). In our case, modelling axisymmetric systems with only two integrals requires a careful check in particular. Any third integral is neglected here completely, which means that the diffusion properties of orbits in the axisymmetric potential are treated as a function of two integrals only.   Although there are successes in previous FP models to explain the dynamical evolution of star clusters, only few of them considered the natural and important physical property of the existence of an initial angular momentum in the star cluster. Kim et al. (2002) have reviewed briefly the previous studies where the initial rotation is considered. They presented the first post-core collapse studies on the evolution of rotating stellar systems and found that the global shape of the rotational structure of the globular cluster changes little, though the strength of the rotation (measured using the magnitude of the rotational velocity or the $z$-component of the angular momentum) decreases continuously with time due to the outward transfer of the angular momentum. An incorporation of mass loss and enhanced two-body relaxation processes accelerated the evolution of the star cluster not only in core-collapse time, but also the dissolution time of the cluster. They found an approach to self-similar evolution in late core-collapse.  The previous studies where the initial rotation is considered, assumed a star cluster with equal mass stars. Inclusion of the initial mass spectrum known to change the evolution of star clusters significantly, shortening of core-collapse times for example (Lee, Fahlman \\& Richer 1991). We also expect that there may be important physical processes between different mass species concerning the exchanges of the angular momentum. Until now, no study has been done on  rotating star clusters with mass spectrum except for one preliminary work by Spurzem \\& Einsel (1998). In this paper, thus we present the study of dynamical evolution of the rotating star cluster with the initial mass function (IMF hereafter). We first investigate the evolution of two-component models as a simplest extension of the single mass models. Then we extend our study to the multi-mass models represented by ten different mass species. The difference of our models as compared to the preliminary work by Spurzem \\& Einsel (1998) is that first our improved code includes more accurate numerical integration and discretisation procedures as described in Paper II, and here we do an extensive parameter study, which did not exist before.  The outline of the present paper is as follows; In section 2, the FP equations for the multi-mass system and the initial models are presented. We concentrate on the evolution up to the core collapse of two component models in section 3 and multi-mass models in section 4. We further discuss the evolution beyond the core collapse for both two-component and multi-component models section 5. We summarize our main results in section 6. ",
        "conclusions": "We have studied the dynamical evolution of the rotating stellar systems with the mass spectrum by solving the orbit-averaged 2D FP equation in ($E,J_z$) space. Numerical simulations are performed both for simple two component clusters and for clusters with a power-law mass function represented by ten mass species. In order to explore the evolution after core collapse we add the heating by three-body processes. We have employed rotating King models as initial models, where the velocity dispersions for all mass components are equal, i.e., no mass segregation at the beginning. The rotating King models are characterized by two parameters: initial central potential ($W_0$) and degree of initial rotation ($\\omega_0$). Clusters with two different central potential $W_0 = 6$ and $3$ are studied extensively. For models only until core-collapse, we have studied the evolution of the clusters with six different mass functions for two-component models, while three models are studied for clusters with a continuous mass function (power-law). For models where the evolution beyond the core-collapse are explored, we consider clusters with a mass function M2C ($m_2/m_1 = 5, M_1/M_2 = 10$) for the two component model and clusters with $\\alpha_0 = -2.35$ for the model with continuous mass spectrum.  Our results show that, as in equal mass system, the presence of the initial rotation accelerates the dynamical evolution as manifested by rapid core collapse and dissolution. The degree of the acceleration depends on the amount of the initial rotation and the shape of mass function. As the ratio of the mass ($m_2/m_1$) increases the degree of the acceleration decreases for two component models since both mass segregation and the rotation compete in acceleration of core-collapse. If $m_2/m_1$ is very large, the mass segregation alone could significantly reduce the core-collapse time and there is not much room for rotation to accelerate further. For models with the power-law mass function, the acceleration rate of core-collapse due to the rotation is larger for the model with a steeper slope of the mass function for a given initial rotation parameter ($\\omega_0$). The shortening of the life-time (dissolution of cluster) due to the rotation is observed far beyond the core bounce. The increase of mass loss rate, resulting from the enhanced two-body relaxation process causes the faster dissolution of cluster.  The evolution of $\\sigma_c$ on $\\rho_c$ can be approximated by a power-law except for the early evolutionary stage regardless of the degree of rotation both for pre- and post-collapse. The measured power law index $\\gamma$ for pre-collapse is very close to the value obtained for single mass system, while we have obtained a slightly shallower slope for post-collapse phase. The development of mass segregation, a consequence of the evolution of the multi-mass system is visible clearly both for non-rotating and rotating models. The evolution of $\\Omega_c$ on $\\rho_c$ shows a power-law behaviour, too. The slope of the power-law is, however, larger than that obtained for the single mass cluster. While the angular momentum is transferred only outward for the equal mass system, the exchange of the angular momentum between different mass species occurs for the multi-mass system, resulting a faster increase of $\\Omega_c$ on $\\rho_c$. Due to a cooperation of the central concentration of the massive stars (mass segregation) and the transfer of angular momentum from high mass to low mass stars, the radii where $V_{rot}$ reaches the maximum value goes outward. The maximum rotation of the most massive stars even increases at the early times, while it shows a monotonic decrease for the single mass system.  \\begin{figure} \\epsfig{figure=fig3_30.eps, height=0.50\\textwidth, width=0.50\\textwidth} \\vspace{-3mm} \\caption{Comparison of observed radial profile of the rotational velocity over 1D velocity dispersion for the galactic globular cluster M15 (data provided by Gebhardt 2002) with the result of a selected two-component model with $(W_0,\\omega_0)=(6,0.6)$, $(m_2/ m_1, M_1/M_2)=(5,10)$.} \\label{fig3-30} \\end{figure}  There are a few observations regarding the direct measurement of the rotation parameter ($V_{rot}$) (Gebhardt et al. 1995). Recently Gerssen et al. (2002) reported the kinematical study of the central part of the globular cluster M15, including $V_{rot}$. M15 is known as the globular cluster which contains a collapsed core. The radial profile of $V_{rot}/\\sigma$ (rotational velocity over one dimensional velocity dispersion) shows a steep increase and a rapid decline followed by a slow rise. The steep rise of $V_{rot}/\\sigma$ near the cluster center can not be explained by a single mass model (Kim et al. 2002). We have shown the run of $V_{rot}/\\sigma$ for two component model in Fig. \\ref{fig3-30} at two selected epochs, one in pre-collapse ($t/\\tau_{rh,0} = 0.42$) and the other after core bounce ($t/\\tau_{rh,0} = 2.41$), together with the observed data by Gebhardt (2002). The radius is measured in current half-mass radius. It is evident that cluster loose the angular momentum through the tidal boundary. The behaviour of the mass-weighted $V_{rot}/\\sigma$ (Fig. \\ref{fig3-30}(a)) are roughly similar to that of observed profile of M15 except for the central region. Note that the known half-mass radius of M15 is $\\sim 3^{'}.09$ (Djorgovski 1993). For the high mass component, $V_{rot,i}/\\sigma_i$ where $i$ represents the individual mass component, at $t = 2.41$ is higher than that obtained at $t = 0.42$, especially beyond $r_h$. It is mainly due to the lower velocity dispersion after core bounce, not due to the higher value of the rotation velocities. The highly rotating central region of M15 is, as in single mass system, not explained with the current multi-mass models, although there is difference in rotational structure between the single mass and the multi-mass systems. From the kinematical study of the central region of M15, Gerssen et al. (2002) proposed the presence of intermediate mass black hole (IMBH) with $M = 3.9\\times10^3 M_{\\odot}$. On the contrary, Baumgardt et al. (2002) demonstrated that the rise of velocity dispersion into the center can be explained with the clustering of the remnant stars, which are expected to be overwhelmingly populated in the cluster core than the normal stars. However, they did not rule out the possibility of the presence of IMBH. It may be necessary to include the remnant stars and the IMBH in current 2DFP models to investigate the role of the initial rotation on the observed strong increasing of $V_{rot}/\\sigma$ near the center of M15.  The models presented here still neglect many important physical processes occurring in real star clusters. The stellar evolution and primordial binaries are known to affect the early evolution of globular clusters. The possible existence of intermediate mass black hole in the center could also affect the course of dynamical evolution. These will be the task of future works.  This work was supported by the Korea Research Foundation Grant No. D00268 in 2001 to HML and by SFB439 to RS."
    },
    "0402/astro-ph0402152_arXiv.txt": {
        "abstract": "Gamma rays with energy above 10~GeV interact with optical-UV photons resulting in pair production. Therefore, a large sample of high redshift sources of these gamma rays can be used to probe the extragalactic background starlight (EBL) by examining the redshift dependence of the attenuation of the flux above 10~GeV. GLAST, the next generation high-energy gamma-ray telescope, will have the unique capability to detect thousands of gamma-ray blazars to redshifts of at least $z=4$, with sufficient angular resolution to allow identification of a large fraction of their optical counterparts. By combining established models of the gamma-ray blazar luminosity function, two different calculations of the high energy gamma-ray opacity due to EBL absorption, and the expected GLAST instrument performance to produce simulated fluxes and redshifts for the blazars that GLAST would detect, we demonstrate that these gamma-ray blazars have the potential to be a highly effective probe of the optical-UV EBL. ",
        "introduction": "}  In the last few years the study of galaxy formation and evolution has seen tremendous progress. Instruments at many different wavelengths have begun to penetrate to the relevant redshifts. One important prediction of models of galaxy formation and evolution is the nature of the radiation field produced by star formation. One way to probe the resulting extragalactic background light (EBL) is to measure the attenuation through pair production of gamma rays from distant sources. However, without a large sample of sources distributed across a wide redshift range, it is difficult to distinguish between extragalactic absorption and characteristics peculiar to individual sources. The Large Area Telescope (LAT) instrument on The Gamma-Ray Large Area Space Telescope (GLAST) will observe gamma rays with energies from 20~MeV to $>300$~GeV. The GLAST LAT will be the first instrument able to probe the intergalactic radiation field by observing the absorption of gamma rays from a large number of extragalactic point sources as a function of redshift over a wide range. Ground-based telescopes can measure the attenuation of TeV emission by intergalactic IR radiation (Stecker, DeJager, \\& Salamon 1992; Macminn \\& Primack 1996; Madau \\& Phinney 1996). However, these telescopes will measure the spectra of relatively small number of sources, making it more difficult to resolve the question of whether differences between sources are due to intergalactic attenuation or intrinsic peculiarities. Furthermore, the high pair production opacity of the IR radiation limits TeV probes of the EBL to a narrow, low-redshift range. GLAST, on the other hand, will observe thousands of sources, and will measure less drastic attenuation of GeV photons by optical and UV radiation. The energy range and capabilities of GLAST are thus ideal for probing the EBL to cosmological distances.  This paper reports our first modeling of the ability of GLAST to measure the extragalactic background light absorption. In order to do this, we need 1) models of the intergalactic radiation field, 2) the luminosity function of extragalactic gamma-ray sources, and 3) parameters of the instrument. In Section 2.1 we briefly review models for the intergalactic radiation field and the resulting gamma-ray opacity as a function of redshift. In Section 2.2, we describe the two gamma-ray blazar luminosity functions used. In Section 2.3 we describe the parameters used to simulate GLAST. In Section 3 we discuss the simulation procedure, including the two different models of blazar input spectra and the two models for the intergalactic radiation field. In Section 4 we present our results and conclusions. ",
        "conclusions": "}  Extragalactic attenuation of gamma-rays by low-energy background photons produces a distortion in the spectra of gamma-ray blazars that increases with increasing redshift. Because we cannot distinguish the difference between extragalactic attenuation and intrinsic effects in individual blazar spectra, statistical analysis of a large sample of blazars such as those presented in this paper is a powerful tool to study EBL absorption. Although AGILE, the next GeV mission (Tavani et al. 2001), will produce a significant increase in the total number of blazars and therefore refine the blazar luminosity function and evolution, GLAST will be the first mission to observe a large sample of high redshift blazars with sufficient statistics to separate intrinsic differences between blazars from redshift dependence of EBL absorption.  Our results indicate that the redshift dependence of the attenuation should be easily detectable by GLAST even when the diffuse background is taken into account and possible high-energy intrinsic rolloffs are considered.  Selection effects, both from GLAST itself and from optical coverage of redshift determinations, will primarily affect sources with low flux. These sources will have poorly measured flux ratios, and will suffer from optical selection effects due to their more poorly determined positions. Other biases include the locations of optical telescopes, source clustering, and other effects. It will be important to catalog these effects explicitly; in particular, insuring adequate optical coverage may require active preparation and participation.  GLAST will be able to measure the differences in blazar attenuation in the cosmologically interesting range in redshift from $z=1$ up to $z=5$. This is in contrast to ground-based observations of TeV attenuation by IR radiation, which will only be able to measure differences well below $z=1$, where the IR becomes opaque. As the energy threshold of the ground based experiments drops over time, their redshift range will increase, but will remain limited to low redshifts except for exceptional, statistically insignificant special cases, especially given their generally small fields of view. More than establishing that EBL attenuation occurs, GLAST will be able to distinguish between different EBL models. This would validate EBL attenuation as a direct cosmological probe.  We emphasize that this analysis will require redshift determinations of a large fraction of GLAST blazars. This is another example of the importance of cross-wavelength studies: by using optical measurements of blazar redshifts, gamma-ray measurements can uniquely probe the optical-UV EBL. A redshift measurement for thousands of high-redshift sources is not a trivial undertaking, but the effort will be well rewarded.  Even after observation of a redshift-dependent effect, the possibility would remain that the spectral evolution of gamma-ray blazars might coincidentally mimic redshift-dependent EBL absorption. For example, if blazars that formed in the early universe suffered more internal attenuation than blazars that formed later, the same effect could be produced. Note that blazars are variable, and there are some indications that their spectra can become harder when they flare (Sreekumar et al. 1996). Evolution in flaring probability could produce the same effect as actual spectral evolution from a statistical standpoint (for example, a higher percentage of high-redshift blazars might be observed in a quiescent phase), although one would expect the GLAST flux limit to produce a selection effect in the opposite direction. In any case, observation of a redshift-dependent spectral softening will provide an important constraint. Theorists will have to decide the likelihood of an evolutionary conspiracy."
    },
    "0402/astro-ph0402014_arXiv.txt": {
        "abstract": " ",
        "introduction": "Kinetic temperature and density are fundamental parameters for our understanding of the interstellar medium (ISM). Usually, symmetric rotors such as NH$_3$ are used to probe a cloud's kinetic temperature, while linear molecules, e.g. CS, probe its density.  However, different spatial distributions of the tracers (``chemistry'') often complicate the picture (see, e.g., \\cite{tafalla}) as they often trace physically different and spatially non-coexisting gas components. It is thus desirable to trace all relevant physical parameters with a single molecule. Promising candidates exist among slightly asymmetric rotors, which have properties qualifying them as tracers for physical conditions. Methanol, CH$_3$OH, is a slightly asymmetric rotor. It is ubiquitous and associated with different regimes of star formation, from quiescent, cold ($\\mathrm T \\sim 10$ K), dark clouds, to ``hot core'' sources near high-mass (proto)stellar objects, where [CH$_3$OH/H$_2$] values $\\sim 10^{-7}-10^{-6}$ are observed \\cite{menten}. Up to now an extremely poor knowledge of the CH$_3$OH collisional rates and of their propensity rules has prevented realistic systematic studies exploiting methanol's full potential as an interstellar tracer. Recently, this situation has changed with the calculation of collisional rate coefficients by \\cite{pottage1,pottage2}, for collisions with helium, for both CH$_3$OH-$A$ and CH$_3$OH-$E$, for levels up to $(\\mathrm J,\\mathrm K)=9$.\\\\ Here we would like to focus on general aspects connected to the analysis of complex molecules' spectra, of which CH$_3$OH is one of the simplest examples, (for details on methanol excitation and on its probing properties see \\cite{leurini}). ",
        "conclusions": ""
    },
    "0402/hep-th0402086_arXiv.txt": {
        "abstract": "Quintessential inflation describes a scenario in which both inflation and dark energy (quintessence) are described by the same scalar field. In conventional braneworld models of quintessential inflation gravitational particle production is used to reheat the universe. This  reheating mechanism is very inefficient and results in an excessive production of gravity waves which violate nucleosynthesis constraints and invalidate the model. We describe a new method of realizing quintessential inflation on the brane in which inflation is followed by `instant preheating' (Felder, Kofman \\& Linde 1999). The larger reheating temperature in this model results in a smaller amplitude of relic gravity waves which is consistent with nucleosynthesis bounds. The relic gravity wave background has a `blue' spectrum at high frequencies and is a generic byproduct of successful quintessential inflation on the brane. ",
        "introduction": "One of the most remarkable discoveries of the past decade is that the universe is accelerating. An accelerating universe is supported by observations of high redshift type Ia supernovae treated as standardized candles \\cite{sn1,sn2} and, more indirectly, by observations of the cosmic microwave background and galaxy clustering \\citep{wmap,tegmark03}. Within the framework of general relativity, cosmic acceleration should be sourced by an energy-momentum tensor which has a large negative pressure (dark energy). The simplest form of dark energy is undoubtedly the cosmological constant for which $p = -\\rho$ = constant. However, due to its un-evolving nature, the cosmological constant must be set to an extremely small value in order to dominate the expansion dynamics of the universe at precisely the present epoch. This gives rise (according to one's viewpoint) either to a fine-tuning problem or to a `cosmic coincidence' problem. For this reason theorists have suggested making dark energy a dynamical quantity associated with a Lagrangian and having well defined equations of motion (see \\cite{ss00,carroll01,pr02,sahni02a,paddy03} for reviews of dark energy). Perhaps the simplest dynamically evolving dark energy models are quintessence fields -- scalar fields which couple minimally to gravity and which roll down a steep potential \\cite{peebles88,wetterich88}. Although quintessence models do not resolve the `cosmic coincidence' conundrum they do alleviate, to some extent, the fine-tuning problem faced by the cosmological constant since they approach a common evolutionary path from a wide range of initial conditions.  Braneworld models \\cite{randall,shiromizu} add an interesting new dimension to scalar field dynamics on the brane. The presence of a quadratic density term (due to high energy corrections) in the Friedman equation on the brane fundamentally alters the expansion dynamics at early epochs by greatly increasing the Hubble parameter and hence the damping experienced by the scalar field as it rolls down its potential \\cite{maartens}. Consequently, inflation on the brane can be realized by very steep potentials -- precisely those used to describe quintessence. The braneworld scenario therefore provides us with the opportunity to unify inflation  and dark energy through a mechanism called quintessential inflation \\cite{pv99,copeland,lidsey1,sss02,majumdar,nkd}. Models of quintessential inflation have a single major drawback: they are usually derived from {\\em non-oscillating} potentials for which the standard reheating mechanism does not work. Indeed, in order to ensure that the inflaton survives until today one usually invokes a method of reheating based on the quantum mechanical production of particles in the time-varying gravitational field after inflation \\cite{ford,spokoiny}. This method of reheating is very inefficient, and leads to a `kinetic regime' of prolonged duration when braneworld corrections are no longer important and the scalar field rapidly drops down a steep potential, resulting in $p_\\phi \\simeq \\rho_\\phi \\simeq {\\dot\\phi^2}/2$ and $a \\propto t^{1/3}$. Gravity waves, created quantum mechanically during the kinetic regime have a `blue tilt' and, for a prolonged kinetic regime, their energy density can dominate the energy density of the universe and violate nucleosynthesis constraints \\cite{sss02} (see also \\cite{star79,sahni90,ss92,giovan98,lmw00,kobayashi,hiramatsu,elmw03}). Thus conventional braneworld models of quintessential inflation run into serious problems associated with copious graviton production which renders them unviable for an extended region in parameter space. As we shall show in this paper, this problem is easily circumvented if, instead of gravitational particle production, one invokes an alternative method of reheating, namely `instant preheating' proposed by Felder, Kofman and Linde \\cite{FKL,FKL1,kls97} (see also \\cite{kls94,stb95,R,chiba}). This method results in a much higher reheat temperature and therefore in a much shorter duration kinetic regime. As a result the amplitude of relic gravity waves is greatly reduced and there is no longer any conflict with nucleosynthesis constraints. (For other approaches to reheating in quintessential inflation see \\cite{curvaton1,dimo,curvaton2,bi1,bi2}. An alternate approach to dark energy in braneworld models is provided in \\cite{DDG,ss03}.) Before we apply the instant preheating method to quintessential inflation on the brane let us briefly review scalar field dynamics in braneworld cosmology. ",
        "conclusions": "In this paper we have successfully applied the instant preheating mechanism discussed in \\cite{FKL,FKL1} to braneworld inflation. Braneworld cosmology has the attractive property of allowing inflation to occur even for steep potentials thus greatly expanding the class of potentials which give rise to inflation. We have demonstrated that instant preheating works remarkably well for non-oscillating potentials such as an exponential. In such models the instant preheating mechanism results in a much higher energy density for radiation than the mechanism based on gravitational particle production \\cite{ford,spokoiny}.  One of the main results of this paper is that a single scalar field on the brane can successfully describe inflation at early epochs and quintessence at late times. Braneworld models of quintessential inflation followed by instant preheating do not over-produce gravity waves and are therefore entirely consistent with the supernova data on the one hand and nucleosynthesis constraints on the other. Although the amplitude of the gravity wave background in quintessential inflation is below the projected sensitivity of LISA, the `blue spectrum' of gravity waves with wavelengths $\\lleq 10^{10}$ cm, makes it likely that that they could lie within the sensitivity range of future space-based gravity wave detectors (see for example \\cite{seto01}).  We would like to end by mentioning that the main features of instant preheating in braneworld models are sufficiently robust and are likely to carry over to other inflationary models in which enhanced damping is an important feature of early scalar field dynamics. In this paper we have laid stress on the `heavy damping' experienced by a scalar field during the RS phase in order to construct a successful scenario of quintessential-inflation on the brane. It should be emphasised however that the RS phase also occurs in braneworld models which have a Gauss-Bonnet (GB) term in the bulk; see equations (\\ref{action}) - (\\ref{gr}). Indeed it is to this class of models to which one must turn for a scenario, which not only gives rise to quintessential-inflation, but also satisfies all other observational constraints -- particularly those on the spectral index and the scalar to tensor ratio placed by WMAP+SDSS \\cite{spergel03,tegmark03}. It is well known that steep inflation in the RS scenario comes into conflict with observations \\cite{suji04} while GB brane world can rescue these models \\cite{ssr}. Indeed, it appears that steep inflation in GB models which commences at an intermediate energy scale between RS and GB regimes -- can be easily reconciled with observations \\cite{ssr}. Figure \\ref{exponential.ps} summarizes this result by showing the values of the scalar to tensor ratio $R$ and the scalar spectral index $n_S$ in the GB inflationary model ; for details the reader is referred to \\cite{ssr}. It is interesting that the $R(n_S)$ curve shows a minimum for steep inflation which begins at an intermediate energy scale in GB inflation and which agrees with observations. (The upper right branch of the curves corresponds to steep inflation which commenced deep in the GB regime; the left branch corresponds to the RS regime.) We conclude that a successful scenario of quintessential inflation on the Gauss-Bonnet braneworld has been constructed which agrees with CMB+LSS observations and also generates an interesting blue spectrum for gravity waves on small scales.  \\bigskip"
    },
    "0402/astro-ph0402129.txt": {
        "abstract": "This is a review on cosmological perturbation theory. After an introduction, it presents the problem of gauge transformation. Gauge invariant variables are introduced and the Einstein and conservation equations are written in terms of these variables. Some examples, especially perfect fluids and scalar fields are presented in detail. The generation of perturbations during inflation is studied. Lightlike geodesics and their relevance for CMB anisotropies are briefly discussed. Perturbation theory in braneworlds is also introduced. ",
        "introduction": "\\index{introduction} The idea that the large scale structure of our Universe might have grown our of small initial fluctuations via gravitational instability goes back to Newton (letter to Bentley, 1692\\cite{NewtonB}).  The first relativistic treatment of linear perturbations in a Friedmann-Lema\\^\\i tre universe was given by Lifshitz (1946)\\cite{Lif46}. There He found that  the gravitational potential cannot grow within linear perturbation theory and he concluded that galaxies have not formed by gravitational instability.  Today we know that it is sufficient that matter density fluctuations can grow. Nevertheless, considerable initial fluctuations with amplitudes of the order of $10^{-5}$ are needed in order to reproduce the cosmic structures observed today. These are much larger than typical statistical fluctuations on scales of galaxies and we have to propose a mechanism to generate them. Furthermore, the measurements of anisotropies in the cosmic microwave background show that the amplitude of fluctuations is constant over a wide range of scales, the spectrum is scale independent.  As we shall see, standard inflation generically produces such a spectrum of nearly scale invariant  fluctuations.  In this course I present gauge invariant cosmological perturbation theory. I shall start by defining gauge invariant perturbation variables. Then I present the basic perturbation equations. As examples for the matter equations we shall consider perfect fluids and scalar fields. The we briefly discuss lightlike geodesics and CMB anisotropies (this section will be very brief since it is complemented by the course on CMB anisotropies by A. Challinor). Finally, I shall make some brief comments on perturbation theory for braneworlds, a topic which is still wide open in my opinion. ",
        "conclusions": "In this course I have given an introduction to cosmological perturbation theory. Perturbation theory is an important tool especially to calculate CMB anisotropies and polarisation since these are very small and can be determined reliably within linear cosmological perturbation theory. To determine the evolution of the cosmic matter density, linear perturbation theory has to be complemented with the theory of weakly non-linear Newtonian gravity and with N-body simulations. To finally understand the formation of galaxies non-gravitational highly non-linear physics, like heating and cooling mechanisms, dissipation, nuclear reactions etc. have to be taken into account. This very difficult subject is still in its infancy.  To make progress in our understanding of braneworlds, linear perturbation theory can also be most helpful. We can use it to determine e.g. the propagating modes of the gravitational field on the brane, light deflection and redshift in weak gravitational fields and the Newtonian limit. The condition that linear perturbations on the brane at low energy and large distances reduce to those resulting from Einstein gravity is non--trivial and has, to my knowledge, not yet been fully explored to limit  braneworld models.  \\acknowledgement I thank the organizers for a well structured school in a most beautiful environment.  % ----------------------------------------------------------------"
    },
    "0402/astro-ph0402618_arXiv.txt": {
        "abstract": "We calculate the one-point probability distribution function (PDF) for cosmic density $\\delta$ in non-linear regime of the gravitational evolution. Under the local approximation that the evolution of cosmic fluid fields can be characterized by the Lagrangian local dynamics with finite degrees of freedom, the analytic expressions of PDF are derived taking account of the smoothing effect. The validity and the usefulness of the local approximation are then discussed comparing those results with N-body simulations in a Gaussian initial condition. Adopting the ellipsoidal collapse model (ECM) and the spherical collapse model (SCM) as Lagrangian local dynamics, we found that the PDFs from the local approximation excellently match the simulation results in the case of the cold dark matter initial spectrum. As for the scale-free initial spectra given by $P(k)\\propto k^n$, N-body result suffers from spurious numerical effects, which prevent us to give a detailed comparison. Nevertheless, at the quality of N-body data, the model predictions based on the ECM and the SCM quantitatively agree with N-body results in cases with spectral index $n<0$. For the index $n\\ge0$, choice of the Lagrangian local dynamics becomes crucial for an accurate prediction and a more delicate modeling is required, however, we find that the model prediction based on the ECM provides a better approximation to the N-body results of cumulants and PDFs. ",
        "introduction": "\\label{sec:intro} The probability distribution function (PDF) of the cosmological density fluctuation is a fundamental statistical quantity characterizing the large-scale structure of the universe. In a standard picture of cosmic structure formation based on the cold dark matter scenario, the gravitational evolution of the dark matter distribution plays an essential role for the hierarchical nature of observed luminous distributions. Usually, the evolution of dark matter distribution is believed to be developed from a small initial fluctuation with Gaussian random distribution. While the PDF of density fluctuation retains Gaussian shape in a linear regime, the deviation from Gaussian distribution becomes significant in the non-linear regime of gravitational evolution.  A number of studies in quantifying the non-Gaussian properties of density field have been developed theoretically and observationally. From the numerical and the observational study, a systematic analysis using the cosmological N-body simulation or the observed galaxy distribution yield various phenomenological prescription for the density PDF in the non-linear regime \\citep[e.g.,][]{SH1984,H1985,GY1993,UY1996}. Among them, the lognormal distribution has been long known to fit to the simulations quite accurately \\citep[e.g.,][]{CJ1991,CMS1993,BK1995,TW2000}. Recently, \\citet{KTS2001} critically examined this issue using the high resolution N-body simulation with Gaussian initial conditions, and found that the accuracy of the lognormal model remains valid, irrespective of the nature of initial spectra. The weak dependence of the initial spectra was later investigated using the phenomenological models with dark halo approach \\citep{THK2003}.   On the other hand, from the analytical study, a perturbative construction of the PDFs has been exploited by \\citet{B1992,B1994a} employing a field-theoretical approach and the predictions including the smoothing effect excellently match the N-body simulations in the weakly non-linear regime. Beyond the perturbative prediction, however, no exact treatment is present and the non-perturbative approximation or the phenomenological approach taking account of the empirical simulation results are necessary. \\citet{FG1998} and \\citet{SG2001} proposed to use a spherical collapse model as a non-perturbative approximation to predict the higher-order moments and PDFs. In their treatment, one assumes that the Lagrangian dynamics of the local density field is simply described by spherical collapse model. Although this approximation clearly misses the non-locality of the gravity in the sense that the evolution of the local density field can be determined by the one-to-one local mapping, the advantage of this treatment is that one can easily calculate the higher-order correction of the moments and PDFs. Further, it turns out that the spherical collapse approximation exactly recovers the leading order results of perturbation theory.  Recently, we generalize the idea of spherical collapse approximation to the local approximation in which the evolution of local density field is characterized by the Lagrangian local dynamics with finite degrees of freedom \\citep{OKT2003}. As a demonstration, the PDFs were computed using the ellipsoidal collapse model. In the ellipsoidal collapse model, the local density at a position is expressed as the multivariate function of initial parameters, i.e., the principal axes of the ellipsoid given at the same position. Thus, the relation between the initial and evolved density field cannot be described by the one-to-one local mapping. As a consequence, the local approximation with ellipsoidal collapse model successfully explains the stochastic nature seen in the simulation, i.e., the joint probability between the initial and the evolved density fields, as has been reported by \\citet{KTS2001}. In addition, the leading-order results from the ellipsoidal collapse model correctly reproduce those obtained from the exact perturbation theory.   In the present paper, we extend the previous study to the quantitative comparison between the local approximation and the N-body simulations. Evaluating the PDF of local density fields taking account of the smoothing effect, we consider the validity and the limitation of the local approximation with spherical and ellipsoidal collapse models. The PDFs from the spherical collapse model were previously compared with the N-body simulations in the case with cold dark matter (CDM) power spectrum \\citep{SG2001}. In this paper, taking account of the smoothing effect, the N-body results with the scale-free initial spectra as well as the CDM spectrum are compared.  This paper is organized as follows. In section \\ref{sec:LCM}, we start to review the local approximation of one-point statistics developed by \\citet{OKT2003} and briefly show how to compute the PDF and the moments from the Lagrangian local collapse model. As representative models of Lagrangian local dynamics, the spherical and the ellipsoidal collapse model are considered. Then, we consider the smoothing effect and discuss how to incorporate it into the model predictions. Based on this, the perturbative calculation of cumulants up to the two-loop order is presented and the qualitative behaviors of the model prediction are discussed in section \\ref{sec:perturbation_ECM}. In section \\ref{sec:results}, the validity and the usefulness of the local approximation for one-point statistics is investigated by comparing the PDFs and cumulants from the local collapse models with those obtained from the N-body simulations. Finally, section \\ref{sec:conclusion} is devoted to discussion and conclusions. ",
        "conclusions": "\\label{sec:conclusion} In the present paper, we critically examined the validity and the usefulness of the local approximation to the PDF and the cumulant predictions. Adopting the ellipsoidal and the spherical collapse model as representative model of the Lagrangian local dynamics, the PDFs and the cumulants are calculated taking account of the smoothing effect and the resultant predictions are compared with the N-body simulations with a Gaussian initial condition. Due to the cutoff of the density arising from the spurious numerical effects, the detailed comparison in cumulants becomes difficult, and the correction for the cutoff density should be self-consistently incorporated into the model prediction. At a level of the quality of the N-body data, however, the local approximation with both SCM and ECM successfully reproduces the N-body results for the PDFs and the cumulants, although a self-consistent calculation of local approximation presented in this paper (labeled by ``cutoff-2'') is still needed to be improved. This is indeed the case of the $\\Lambda$CDM model and the scale-free models with indices $n=-2,~-1$. For the scale-free model with $n=0$, while the discrepancy between the model prediction and the simulation result is manifest in the local approximation with SCM, the agreement with N-body results still remains good for the ECM prediction. The detailed discussion reveals that the prediction based on the local approximation sensitively depends on the slope of the initial spectrum and that the predictions for $n>0$ become more sensitive to the non-linear dynamics of the local collapse model. Thus, a more delicate modeling of the Lagrangian local dynamics is required for an accurate prediction. Taking this point carefully, we therefore conclude that the local approximation with SCM and ECM provides an excellent approximation to the N-body simulations for CDM and scale-free models with $n<0$ in both the linear and the non-linear regimes, $0\\simlt \\sigma_l\\simlt 5$, while the local approximation should be used with caution in the $n\\geq0$ cases.   In this paper, we found that the predictions based on the ellipsoidal collapse model somewhat improve the approximation, however, the degree of the improvement is not so large as long as a CDM-like initial spectrum (i.e., effective spectral index $n_{\\rm eff}=-3-d\\log\\sigma_l^2/d\\log R<0$) is concerned. Compared to the prediction from the spherical collapse model, the calculation of PDF from ellipsoidal collapse model is rather complicated and require a time-consuming numerical integration. It seems that the spherical collapse model provides a simpler prescription for the PDF in real space and is practically more useful than the ellipsoidal collapse model. However, if one considers the one-point statistics in redshift space, the situation might be changed drastically. As reported by \\citet{SG2001}, the local approximation with spherical collapse model only provides a good approximation to the redshift space PDFs when $\\sigma_l\\simlt0.4$. A part of this reason is ascribed to the fact that the model prediction cannot recover the linear perturbation result, referred to as the Kaiser effect \\citep{K1987}; the variance in the redshift space, $\\sigma_z^2$ is related to the one in the real space as: \\begin{equation} \\sigma_z^2=\\left(1+\\frac{2}{3}f_\\Omega+\\frac{1}{5}f_\\Omega^2\\right)\\sigma_l^2. \\label{eq:Kaiser_effect} \\end{equation} In contrast to the spherical collapse model, which leads to the incorrect prediction $\\sigma_z^2=(1+f_\\Omega/3)^2\\sigma_l^2$, the Kaiser effect (\\ref{eq:Kaiser_effect}) can be correctly recovered by means of the ellipsoidal collapse model. The derivation of equation (\\ref{eq:Kaiser_effect}) is presented in appendix \\ref{appen:Kaiser_ECM}. This fact is very interesting and also provides an important suggestion that the non-sphericity of the Lagrangian local dynamics play a crucial role in computing the one-point statistics in redshift space and is indeed essential for an accurate prediction. The detailed analysis of the model predictions in redshift space is now in progress and will be described elsewhere."
    },
    "0402/astro-ph0402332_arXiv.txt": {
        "abstract": "{ We studied the colliding galaxy NGC\\,7714 with two {\\it XMM-Newton} observations, six months apart. The galaxy contains two bright X-ray sources: we show that they have different physical nature. The off-nuclear source is an accreting compact object, one of the brightest ultraluminous X-ray sources (ULXs) found to date. It showed spectral and luminosity changes between the two observations, from a low/soft to a high/hard state; in the high state, it reached $L_{\\rm x} \\approx 6 \\times 10^{40}$ erg s$^{-1}$. Its lightcurve in the high state suggests variability on a $\\approx 2$ hr timescale. Its peculiar location, where the tidal bridge between NGC\\,7714 and NGC\\,7715 joins the outer stellar ring of NGC\\,7714, makes it an interesting example of the connection between gas flows in colliding galaxies and ULX formation. The nuclear source ($L_{\\rm x} \\approx 10^{41}$ erg s$^{-1}$) coincides with a starburst region, and is the combination of thin thermal plasma emission and a point-source contribution (with a power-law spectrum). Variability in the power-law component between the two observations hints at the presence of a single, bright point source ($L_{\\rm x} \\ga 3 \\times 10^{40}$ erg s$^{-1}$): either a hidden AGN or another ULX. ",
        "introduction": "{\\it XMM-Newton} and {\\it Chandra} studies of colliding or merging gas-rich galaxies at various stages of the evolutionary sequence (Toomre 1977) have revealed significant contribution to the X-ray emission both from diffuse hot gas, associated to starburst processes, and from accreting point sources, generally associated to a young stellar population. See, for example: the Mice (Read 2003); the Antennae (Zezas et al.~2002; Fabbiano et al.~2003a); M\\,82 (Griffiths et al.~2000); the Cartwheel (Gao et al.~2003). A peculiar feature in many of these systems is the presence of accreting X-ray sources brighter than the Eddington limit for a stellar-mass black hole (BH) ($L_{\\rm Edd} \\approx 10^{39}$ erg s$^{-1}$); they are commonly known as ultraluminous X-ray sources (ULXs). The ages and masses of the compact objects in ULXs, the nature of the donor stars, and the geometry of emission are still unclear, and hotly debated. It is also unclear precisely why ULXs are preferentially found in interacting galaxies (Swartz et al.~2003), and what this can reveal about their mechanism of formation.  The interacting system Arp 284 (Arp 1966) is an exceptional example of a recent ($\\sim 100$--$200$ Myr ago), direct impact (Struck \\& Smith 2003). It consists of the nuclear starburst galaxy NGC\\,7714 (classified as SB(s)b pec\\footnote{NED: NASA Extragalactic Database}) and its fainter, currently inactive companion NGC\\,7715 (Im pec). NGC\\,7714 is located at a redshift distance of $37.3$ Mpc (Huchra et al.~1999, for $H_0 = 75$ km s$^{-1}$ Mpc$^{-1}$). It has a prominent stellar ring, three tidal arms/tails, and is connected to NGC\\,7715 by a gas and stellar bridge (Arp 1966). Its low inclination (viewing angle of $45^{\\circ}$) and low foreground absorption ($n_{\\rm H} = 4.9 \\times 10^{20}$ cm$^{-2}$; Dickey \\& Lockman 1990) make it a good target for X-ray studies.  {\\it ROSAT}/HRI observations of NGC\\,7714 (Papaderos \\& Fricke 1998) have revealed two strong ($L_{\\rm x} > 10^{40}$ erg s$^{-1}$) X-ray sources, separated by $\\approx 22\\arcsec$. The brighter one coincides with the starburst nucleus; the fainter one is located approximately where the gas/stellar bridge joins a collisional stellar ring, but does not correlate with any bright counterpart at other wavelengths, nor is it located in a starburst region. It was suggested (Papaderos \\& Fricke 1998) that the off-nuclear source might be a compact region of hot shocked gas. This could be due either to the collision of a fast starburst-driven outflow with the colder, denser gas in the galactic bridge; or, to the infall of high-velocity H\\,{\\footnotesize{I}} clouds along the bridge onto the outer H\\,{\\footnotesize{I}} disk of NGC\\,7714 (a somewhat similar situation to accretion-disk hot spots in X-ray binaries). However, the limited wavelength coverage and resolution of {\\it ROSAT} did not allow detailed individual analyses of the two sources.  Here we present some preliminary results of our {\\it XMM-Newton} study of the system. We argue that the off-nuclear source is an accreting point source (in fact, one of the most luminous ULXs ever detected), and we discuss the spectral and timing properties of the two sources.      \\begin{figure} \\includegraphics[width=8.8cm]{epic_color.eps} \\includegraphics[width=8.8cm]{dssmap3.ps} \\caption{Top panel: true-color image of the NGC\\,7714 field obtained by coadding the MOS and pn observations of 2002 Jun 7 and Dec 8. The image was smoothed with a $3\\times3$ pixel boxcar. Red: $0.2$--$1$ keV; green: $1$--$2$ keV; blue: $2$--$12$ keV. Size of the image: $7\\farcm8 \\times 6\\farcm2$; North is up, East is to the left. Bottom panel: on the same scale, the positions of the three brightest X-ray sources in the field are overplotted (red circles) onto a Digitized Sky Survey $B$ image; the positions of a few other, fainter X-ray sources detected with {\\it XMM-Newton}/EPIC are also overplotted (green circles and ellipse). } \\label{FigVibStab} \\end{figure}    \\begin{figure} \\includegraphics[width=9cm]{hstmap3.ps} \\caption{Position of the ULX and of the X-ray nucleus (red circles with radius $= 3\\arcsec$) overplotted on an HTS/WFPC2 image in the $V$ band (f555w filter). The ULX is located where the tidal bridge meets the collisional stellar ring. North is up, East is to the left. } \\label{FigVibStab} \\end{figure} ",
        "conclusions": "We have used {\\it XMM-Newton} to study the interacting galaxy system NGC\\,7714/15. We have reported here the main properties of the two brightest sources: the starburst nucleus and an off-nuclear ULX. The X-ray spectrum of the off-nuclear source suggests that it is an accreting BH, and rules out the possibility that it is due to thermal-plasma emission from a hot spot, as previously speculated. Its X-ray flux varies by a factor of 2 over a six months' interval; the source appears softer in the low state, unlike the typical behaviour of Galactic BH candidates but in agreement with the behaviour of many ULXs. Its spectrum in the low/soft state can be fitted by a disk-blackbody model with $kT_{\\rm in} \\approx 1$ keV: this is inconsistent with a Shakura-Sunyaev disk, but can be explained with a slim-disk model.  In the high state, its emitted isotropic luminosity is $\\approx 6 \\times 10^{40}$ erg s$^{-1}$ in the $0.3$--$12$ keV band, implying a bolometric luminosity $\\approx 1.5 \\times 10^{41}$ erg s$^{-1}$ from a reasonable extrapolation of the power-law spectrum (photon index $\\Gamma \\approx 2$). BH masses of $\\sim$ a few $10^2$--$10^3 M_{\\odot}$ would be required to satisfy the Eddington limit. Furthermore, variability on timescales of $\\approx 2$ hr is detected in the high state.  The ULX is located at the junction of the tidal bridge (consisting of gas and young stars) with the collisional outer ring (consisting of an old stellar population, with no gas). We have pointed out that ULXs are often found in tidally interacting systems, associated with metal-poor molecular clouds, tidal dwarfs, or H{\\footnotesize I} structures formed in the galactic collision.  The nucleus has an X-ray luminosity $\\approx 10^{41}$ erg s$^{-1}$ in the $0.3$--$12$ keV band. Thermal plasma emission contributes for $\\approx 3 \\times 10^{40}$ erg s$^{-1}$, constant over the two observations, and is probably extended (marginally resolved in the EPIC/MOS images). A point-like power-law component contributes for $\\approx 5 \\times 10^{40}$ erg s$^{-1}$ and $\\approx 8 \\times 10^{40}$ erg s$^{-1}$ in the two observations. The power-law component in the X-ray spectra of starburst nuclei is generally due to unresolved high-mass X-ray binaries. The amount of variability in our case implies that one single source contributes for at least $3 \\times 10^{40}$ erg s$^{-1}$. This suggests that there is either a hidden AGN or another ULX in the nuclear region."
    },
    "0402/astro-ph0402104_arXiv.txt": {
        "abstract": "{We present the first results from our search for close stellar and sub-stellar companions to young nearby stars on the northern sky. Our infrared imaging observations are obtained with the 3.5\\,m Calar Alto telescope and the AO system ALFA. With two epoch observations which were separated by about one year, we found two co-moving companion candidates, one close to HD\\,77407 and one close to GJ\\,577. For the companion candidate near GJ\\,577, we obtained an optical spectrum showing spectral type M4.5; this candidate is a bound low-mass stellar companion confirmed by both proper motion and spectroscopy. We estimate the masses for HD\\,77407\\,B and GJ\\,577\\,B to be $\\sim 0.3$ to $0.5\\,M_{\\odot}$ and $\\sim 0.16$ to $0.2\\,M_{\\odot}$, respectively. Compared to Siess al. (2000) models, each of the two pairs appears co-eval with HD\\,77407\\,A,B being 10 to 40\\,Myrs and GJ\\,577\\,A,B being $\\ge$ 100\\,Myrs old. We also took multi-epoch high-resolution spectra of HD\\,77407 to search for sub-stellar companions, but did not find any with $3\\,M_{Jup}$ as upper mass ($m\\sin\\,i$) limit (for up to 4 year orbits); however, we detected a long-term radial velocity trend in HD\\,77407\\,A, consistent with a $\\sim 0.3$\\,M$_{\\odot}$ companion at $\\sim 50$\\,AU separation, i.e. the one detected by the imaging. Hence, HD\\,77407\\,B is confirmed to be a bound companion to HD\\,77407\\,A. We also present limits for undetected, but detectable companions using a deep image of HD\\,77407\\,A and B, also observed with the Keck NIRC2 AO system; any brown dwarfs were detectable outside of 0.5\\,arcsec (17\\,AU at HD\\,77407), giant planets with masses from $\\sim 6.5$ to 12\\,M$\\rm_{Jup}$ were detectable at $\\ge 1.5$\\,arcsec. ",
        "introduction": "Most nearby stars are quite old, so that close sub-stellar companions are too faint to be detected directly. If we consider young stars, their companions are also young and therefore self-luminous due to accretion and contraction, see e.g. Wuchterl \\& Tscharnuter (2003), who consider objects even younger than in our observations; other teams (e.g. Baraffe et al. 1998, Burrows et al. 1997) also show quantitatively how sub-stellar objects get fainter when they get older. However they use arbitrary initial conditions, so that their models should not be used for objects younger than $\\sim 10$\\,Myrs (Baraffe et al. 2002). If such young companions are also nearby, they should be well separated from their primaries. Hence, young nearby stars are most suitable for direct imaging of sub-stellar companions, see e.g. Jayawardhana \\& Greene (2001) for a comprehensive overview. There are $\\sim 200$ stars known with ages from $\\sim 1$ to 100\\,Myrs within $\\sim 100$\\,pc (e.g. Montes et al. 2001b, Wichmann et al. 2003), about one third of them being in the northern sky. Most of our targets still show lithium absorption and/or Ca II, H and K emission; some of them also show H$\\alpha$, and/or are classified as members of young associations by Montes et al. (2001b).  The direct detection of close sub-stellar companions is currently less difficult in the near infrared, because in these wavelengths the brightness difference between companion and primary star is low and detectors work well; in the thermal infrared, the brightness difference is even lower, but the detectors are not yet as sensitive as in JHK. Nevertheless these close companions are much fainter than their host star. The application of an adaptive optics system increases the resolution to the diffraction limit and advances the dynamic range so that the detection of these faint objects becomes feasible.  The determination of the spectral type and hence companionship of faint companion candidates using JHK colors is difficult, because such objects are too faint and located in the PSF wing of their host star, so that colors cannot be measured well. To solve simultaneously both extinction and spectral type, three-band imaging would be indispensable. Nevertheless one needs a further observation to confirm the proper motion of a possible companion. Finally four images are necessary which makes such a photometric search inefficient. On the other hand, an astrometric survey (in one band) only needs two observations (1st and 2nd epoch), i.e. allows one to study many more targets with a minimum of observation time.  Most of the nearby stars have high proper motion and therefore they are well suited for an astrometrical survey. In this technique each star with at least one faint object nearby is observed in two epochs. Companion and primary star show the same motion relative to non-moving background stars. The proper motion of our target stars is high enough so that an epoch difference of one year is sufficient in most cases to find co-moving companions.  For a co-moving companion (candidate), spectroscopic confirmation is always necessary, either by taking a spectrum of the companion (showing its late spectral type) and/or by taking spectra of the primary (showing its secular acceleration due to the companion). We are searching for sub-stellar companions also on the southern sky, with speckle and normal imaging at the ESO NTT, and now also with AO (NAOS-CONICA at the ESO VLT). Four sub-stellar companions to young (nearby) stars have been confirmed by both proper motion and spectroscopy: The $\\sim 12$\\,Myrs young TWA-5 B (Lowrance et al. 1999, Neuh\\\"auser et al. 2000), the $\\sim 300$ Myrs old Gl\\,569\\,B,C (Mart\\'in et al. 2000, Lane et al. 2001), the $\\sim 35$\\,Myrs young HR\\,7239\\,B (Lowrance et al. 2000, Guenther et al. 2001), and the $\\sim 300$\\,Myrs old HD\\,130948\\,B,C (Potter et al. 2002, Goto et al. 2002).  Here, we present some first results from our ongoing imaging program at the Calar Alto observatory. We present the instrument and data reduction in Sect.\\,2, the astrometric results in Sect.\\,3, photometry in Sect.\\,4, and spectroscopy in Sect.\\,5 \\& 6. Finally, in Sect.\\,7, we present the H-R diagram to determine masses and ages of the new companions and discuss our results. ",
        "conclusions": "HD\\,77407 is a slow rotating ($v \\cdot \\sin i \\simeq 7$ km/s) G0 star with a notable chromospheric excess emission in the H$\\alpha$, H$\\beta$ and Ca\\,II lines. It is detected as radio source with a flux of 1.67\\,mJy at 20\\,cm and also as EUV source (Wichmann et al. 2003). Montes et al. (2001a) and Wichmann et al. (2003) report EW(Li\\,I)=170...183\\,m\\AA, consistent with our own measurement EW(Li\\,I)=170\\,m\\AA, indicating that it is a very young star. It is identified as a member of the local association (20 to 150\\,Myrs) due to its galactic space motion (Montes et al. 2001a). Our RV monitoring revealed no sub-stellar spectroscopic companion to HD\\,77407, but some scatter probably due to stellar activity (typical for young stars), and also a long-term trend consistent with the mass estimate of the co-moving companion detected in the imaging.  \\begin{figure*} [htb] \\resizebox{\\hsize}{!}{\\includegraphics[angle=0]{siess.eps}} \\caption{H-R diagram with Siess et al. (2000) models with (z=0.02 no overshooting) for HD\\,77407 and GJ\\,577 primary and secondary, our observed $M_{H}$ value versus the temperature obtained from our observed (or derived) spectral type. The isochrones (dotted lines) are labeled by ages in Myrs. The object mass is given in units of solar mass. Furthermore we have plotted data for main-sequence stars (solid black line) from Schmidt-Kaler (M$_{V}$) and Kenyon-Hartmann (H-V and T$_{eff}$) for spectral types F8V to K0V. GJ\\,577\\,A lies on that main-sequence for spectral type G5V as observed, inconsistent with the ZAMS by Siess et al. by $\\sim$500\\,K.} \\end{figure*}  GJ\\,577 is a G5\\,V star with high levels of photospheric magnetic activity and a photometrically determined rotation period of around 4 days (Messina \\& Guinan 1998, Messina et al. 1999). Based on the kinematic criteria this star can be considered as a member of the Hyades supercluster with an age of 600\\,Myrs. A moderate Ca\\,II H and K emission is observed in the spectra as well as an EW(Li\\,I) of 145\\,m\\AA . This Lithium absorption is too strong for a member of the Hyades supercluster and is close to the weakest Li-lines of the local association with an age of 20 to 150\\,Myrs (Montes et al. 2001a, 2001b). Halbwachs et al. (2003) included GJ\\,577 in their spectroscopic observing program, but did not find any close low mass companion, nor any long-term trend, consistent with our data.  If we use the measured absolute H-band magnitudes of HD\\,77407\\,B and GJ\\,577\\,B as well as the stellar age given above, we can determine the companion mass and its effective temperature by using the evolutionary models for low-mass stars from Siess et al. (2000). The spectral type is converted to effective temperature using the scale for main-sequence dwarf stars by Kenyon \\& Hartmann (1995).  In Fig.\\,8, we show the H-R diagram for both new visual pairs and compare their locations with tracks and isochrones by Siess et al. (2000) with a metallicity z=0.02 and no overshooting.  GJ\\,577\\,B appears to be older than 100\\,Myrs or already on the zero-age MS (ZAMS), also consistent with the 20 to 150\\,Myrs given by Montes et al. (2001a, 2001b). Here, it is not possible to determine the age better from the companion alone, because the isochrones lie very close together at its location in the H-R diagram.  We used Siess et al. (2000) models as well as Baraffe et al. (1998) models. We took those sets of models which were available for low-mass objects down to 0.1\\,M$_{\\sun}$. For Baraffe et al. (1998) models (mixing length parameter $\\alpha=1.0$, He abundance $Y=0.275$ and solar metallicity [M/H]=0) and for the Siess et al. (2000) models (with metallicity z=0.02 no overshooting and with Kenyon-Hartmann conversion) the primaries lie below the main sequence, which is unphysical and may indicate problems with that particular set of models. We note that the primary stars are not saturated in the 2MASS images and that the Hipparcos parallaxes of the primaries should not be affected by the much fainter companions. Furthermore we used M$_{V}$ data from Schmidt-Kaler and converted them to M$_{H}$ using (V-H) from Kenyon-Hartmann (1995). Those data are consistent with the 2MASS photometry and lie also under the Siess et al. ZAMS (see Fig.8). For a metallicity [M/H]=-0.5, mixing length parameter $\\alpha$=1 and He abundance Y=0.25 for Baraffe models and metallicity z=0.01 for Siess models, each of the two pairs appears to be co-eval. However it is extremely unlikely that these two young stars are that metal-poor. Hence we do not show those models here.  GJ\\,577\\,B was also detected by McCarthy et al. (2001) and Lowrance et al. (2003). The former show their J-band image having resolved A and B. They give $J = 11.15\\pm0.15$\\,mag and $I \\simeq 13$\\,mag, consistent with spectral type mid-M. Furthermore, McCarthy et al. (2001) quote Lowrance (2001) as having measured the proper motion of GJ\\,577\\,B to be consistent with A. McCarthy et al. (2001) do not show nor mention a spectrum, and the dissertation of Lowrance (2001) is not available to us. Lowrance et al. (2003) show that GJ\\,577\\,B is actually a double star and is called GJ\\,577\\,B\\&C. The components B and C are both close to the sub-stellar limit or brown dwarfs. They have a combined spectral type of M5 to M6 which is marginally consistent with our results.  We can use the ALFA and Keck images to obtain limits for undetected, but detectable faint companion candidates, i.e. for determining the limiting dynamic range achieved in terms of magnitude difference versus separation. We measure the $3 \\sigma$ flux level (for HD\\,77407 for both Keck and ALFA, so that we can compare them), and ratio it to the primary to determine the curves shown in Fig.\\,9.  \\begin{figure*} [htb] \\resizebox{\\hsize}{!}{\\includegraphics[angle=0]{dynamic.eps}} \\caption{Dynamic ranges achieved with ALFA and Keck AO for HD\\,77407, magnitude difference in H between primary and companions versus separation in arcsec. The two companions detected are indicated by circles. The Keck system is obviously much more sensitive than ALFA. We also indicate magnitude differences expected for companions at the upper mass limits for planets and brown dwarfs - computed for an age of 50\\,Myrs using Baraffe et al. (2003). Any brown dwarf should have been detected outside of 0.5\\,arcsec (17\\,AU), and any giant planet above $\\sim 6.5$\\,M$\\rm_{Jup}$ would have been detected outside of 1.5\\,arcsec. The upper x-axis scale in AU is given for the distance of HD\\,77407.} \\end{figure*}  Any additional stellar companion (above $\\sim 75$\\,M$\\rm_{Jup}$ for 50\\,Myrs following Baraffe et al. 2003) should have been detected at a separation of $\\ge 0.4$\\,arcsec, the radius of the semi-transparent coronograph used. Also, any brown dwarf companion above $\\sim 40$\\,M$\\rm_{Jup}$ would have been detected outside of 0.4\\,arcsec. Brown dwarfs of any mass ranging from $\\sim 12$ to $\\sim 75$\\,M$\\rm_{Jup}$ were detectable outside of 0.5\\,arcsec, which is 17\\,AU at the distance of HD\\,77407. We reached a magnitude difference of $\\Delta H=9$\\,mag at 0.5\\,arcsec separation. Giant planets with masses from $\\sim 6.5$ to 12\\,M$\\rm_{Jup}$ were detectable outside of 1.5\\,arcsec (50\\,AU), where we reached $\\Delta H=12$\\,mag.  We will report other detected companion candidates (both background objects as well as possibly bound secondaries) as well as other possibly negative (null) results (i.e. primaries without any detected companion candidates) elsewhere.  \\begin{table} [htb] \\centering \\caption{Stellar properties of HD\\,77407\\,B and GJ\\,577\\,B} \\label{prop} \\begin{tabular}{l|l|l|l} & HD\\,77407\\,B & GJ\\,577\\,B\\\\ \\hline Age $[$Myrs$]$    & 10 ... 40       & $>$ 100   \\\\ Mass [M$_{\\sun}$] & 0.3 ... 0.5     & 0.16 ... 0.2  \\\\ Spec type           & M0 ... M3     & M4 ... M5 \\\\ T$_{\\rm eff}$~$[K]$ & 3850 ... 3470 & 3370 ... 3240 \\\\ \\end{tabular} \\end{table}"
    },
    "0402/astro-ph0402274_arXiv.txt": {
        "abstract": "We present high resolution interferometric observations of the cool atomic and cold molecular ISM of the TDG candidate Arp\\ 245N, an object resembling a dwarf galaxy in the northern tidal tail of the interacting system NGC\\,2992/3. We observed the HI line with the NRAO VLA and the CO(1$\\to$0) transition with the OVRO millimeter interferometer at $5''-6''$ angular resolution (750\\,pc linear resolution). These datacubes offer the required spatial and velocity resolution to determine whether the mass concentration near the tip of the  tail is a genuine feature, and hence a good TDG candidate, or an artefact caused by a fortuitous alignment of our line of sight with the direction of the  tail. A preliminary analysis seems to confirm that  Arp\\,245N is a self--gravitating entity. ",
        "introduction": "Tidal Dwarf Galaxies (TDGs) are objects resembling actively star forming dwarf galaxies and are assembled from the debris (tidal tails and bridges) launched into the IGM by violent galaxy interactions in which at least one member is a gas--rich galaxy. They are composed of stars and gas from the outskirts of one or both of the parent galaxies involved in the interaction. The recent surge of interest in TDGs started with papers by Mirabel, Lutz, \\& Maza (1991) on the Superantennae and Mirabel, Dottori, \\& Lutz (1992) on the Antennae (see also the review by Duc \\& Mirabel 1999). Several groups of authors have since embarked on the exciting topic of TDGs as witnessed by these proceedings (see e.g., the contributions by Duc et al.\\ and by Braine et al.\\, this volume).  Currently outstanding questions are: (i) Are TDGs really recycled objects made of collisional debris or pre-existing galaxies involved in a three--body interaction? (ii) Are  TDG  genuine density enhancements in the tidal tails or are they merely due to projection effects along the line of sight? (iii) Do TDGs form self gravitating entities or are they simple transient condensations? (iv) Are TDGs Dark Matter (DM) dominated, like galaxies in general, and dwarf galaxies in particular, or are they nearly devoid of DM, as theory predicts? (v) Finally do TDGs leave the potential well of their progenitors and hence constitute a sizeable fraction of the known dwarf galaxy population or do they eventually fall back and merge, leaving no trace?  All these questions have actually been raised for the particular TDG candidate identified in the northern tidal tail of Arp\\,245, an interacting  system composed of two spiral galaxies, NGC 2992 and NGC 2993. Although Arp\\,245N was observed at many wavelengths and is one of the best studied TDG candidates,  its  nature as a tidal object or as a real entity have  been challenged. Smith \\& Struck (2001) argued that TDG Arp\\,245N could actually be a preexisting edge-on disk galaxy that is interacting with the other two galaxies. Hibbard et al.\\ (this volume) point out that this system is viewed from an unfavorable perspective, making the projection effects particularly severe.  The combination of high--resolution  kinematical and morphological data  is critical to tackle all these issues. We have therefore carried out HI and CO interferometric observations of the system. ",
        "conclusions": "Arp\\,245N is a typical TDG candidate in the sense that it is a major HI concentration associated with recent star formation which resides near the  tip of the tidal tail. Because the latter is seen close to edge--on, the question thus arises whether the apparent concentration is a genuine feature (Hibbard \\& Mihos, this volume). Bournaud et al.\\ (2003, 2004; see also the contributions by Amram et al., this volume) have run extensive numerical models and shown the characteristic shape in position--velocity space of a {\\em bona fide} TDG and that of a spurious feature. They show that the kinematical signature of projection effects is a change in the sign of the velocity gradient along the tail before reaching its apparent tip. For curved tails that are extended enough in 3D space, a loop-like feature may even be seen in a position--velocity (pV) diagram along the tails.  Figure 2 shows such a pV diagram along the tidal tail connecting the TDG candidate and NGC\\,2992 using the intermediate and high resolution VLA HI data. The signal was actually integrated over a band with a width similar to that of the tail. The loop expected for projection effects is clearly seen on the figure. The part of the tail that is seen curving back towards NGC 2992 (as seen projected on the sky) is actually consistent with our earlier numerical simulations of the system (Duc et al., 2000; see the face--on view in their Fig.~10).  The HI morphology of the tidal tail, as seen at high resolution in Fig.\\ 1, may also give some clues as to the geometry of the system. Its U--like shape could be interpreted as being due to bending of the tidal tail near its apparent tip. The tail is actually not seen perfectly edge--on (as indicated, in the optical,  by the large width of the stellar tail). Thanks to the higher spatial resolution provided by the VLA in its B--configuration, we can hence 'resolve' the projection effects -- which was impossible with the early C--array data. On this map, the HI column density seems to peak in the part of the tail which points back to NGC 2992. This is where OVRO detected the bulk of the molecular gas and where the brightest HII regions in the tail are found. The velocities of all these phases match. The spatial and velocity coincidence between the CO, H$\\alpha$ and HI emission peaks at this location in  the tail is a strong indication  that a genuine condensation  is present there and that this is likely the progenitor of a Tidal Dwarf galaxy.  At the same position, a pV diagram perpendicular to the tidal tail shows  a small scale velocity gradient similar to that expected for a rotating body (see Fig.~3). The peak--to--peak velocity range is 100\\,km\\,s$^{-1}$. A word of caution is warranted here, though. As Duc et al.\\ (2000) mentioned, the  simulated pV diagram along the same direction  in the numerical model  shows a similar gradient. Further simulations are required to disentangle the embedded TDG from the rest of the tail. We should then be able to determine its dynamical mass and, comparing it with the luminous mass (corresponding to the HI condensation in the B--array), probe its dark matter content. Not taking into account the line of sight crowding, and considering all the matter present at the apparent tip of the tail, one derives a dynamical mass similar to the luminous one and equal to $\\sim 2 \\times 10^9$\\,M$_\\odot$. For the above-mentioned reasons, these are most likely overestimates.  In summary, a first analysis of new high resolution HI and CO datacubes tends to  support the existence of a bound entity within the northern tail of  Arp\\,245. However, they also show the kinematical signature expected when  part of the  tail is  bending away along the line of sight, and eventually back to NGC 2992. Because of these projection effects, the size and mass of the embedded TDG candidate derived from low resolution data are probably overestimates."
    },
    "0402/astro-ph0402568_arXiv.txt": {
        "abstract": "{ Based on recent work on spectral decomposition of the emission of star-forming galaxies, we assess whether the integrated 2-10 keV emission from high-mass X-ray binaries (HMXBs), $L_{2-10}^{\\rm HMXB}$, can be used as a reliable estimator of ongoing star formation rate (SFR). Using a sample of 46 local ($z \\mincir 0.1$) star-forming galaxies, and spectral modeling of {\\it ASCA}, {\\it BeppoSAX}, and {\\it XMM}-Newton data, we demonstrate the existence of a linear SFR--$L_{2-10}^ {\\rm HMXB}$ relation which holds over $\\sim$5 decades in X-ray luminosity and SFR. The total 2-10 keV luminosity is {\\it not} a precise SFR indicator because at low SFR (i.e., in normal and moderately-starbursting galaxies) it is substantially affected by the emission of low-mass X-ray binaries, which do not trace the current SFR due to their long evolution lifetimes, while at very high SFR (i.e., for very luminous FIR-selected galaxies) it is frequently affected by the presence of strongly obscured AGNs. The availability of purely SB-powered galaxies -- whose 2-10 keV emission is mainly due to HMXBs -- allows us to properly calibrate the SFR--$L_{2-10}^{\\rm HMXB}$ relation. The SFR--$L_{2-10}^{\\rm HMXB}$ relation holds also for distant ($z \\sim 1$) galaxies in the {\\it Hubble} Deep Field North sample, for which we lack spectral information, but whose SFR can be estimated from deep radio data. If confirmed by more detailed observations, it may be possible to use the deduced relation to identify distant galaxies that are X-ray overluminous for their (independently estimated) SFR, and are therefore likely to hide strongly absorbed AGNs. ",
        "introduction": "X-ray emission of star-forming galaxies (SFGs) consists of various components including discrete sources, such as X-ray binaries and supernova remnants (SNRs), diffuse hot gas, Compton scattering of ambient FIR photons, and possibly an active nucleus. The resulting integrated spectra harbor the signatures of these emission components.  \\begin{figure} \\vspace{3.4cm} \\special{psfile=0500fig1a.ps hoffset=0 voffset=107 hscale=17.0 vscale=19.2 angle=-90} \\hspace{.5cm} \\special{psfile=0500fig1b.ps hoffset=108 voffset=107 hscale=17.0 vscale=19.2 angle=-90}  \\caption{The integrated \"stellar\" emission in the template X-ray spectrum of a star-forming galaxy proposed by Persic \\& Rephaeli (2002). Two different cases are shown: \"standard\" ({\\it left}) and \"top-heavy\" ({\\it right}). Shown for the standard case in ascending order at 7 keV: SNRs, faint LMXBs, HMXBs, and bright LMXBs. On the right, in ascending order at 3 keV, only SNRs and HMXBs are shown. The assumed 0.5-50 keV luminosities and number abundances are: log$L_{\\rm x} = 37.7$ for HMXBs (50 objects) and high-luminosity LMXBs (70 objects), 37.0 for SNRs (20 objects), and 36.7 for low-luminosity LMXBs (130 objects). The spectral components are normalized in energy flux in the 0.5-50 keV band. The spectrum is absorbed through a HI column density $N_{\\rm H}= 10^{22}$ cm$^{-2}$.} \\end{figure}  Persic \\& Rephaeli (2002, hereafter PR02) have quantitatively assessed the roles of the various X-ray emission mechanisms in SFGs. They have used an equilibrium stellar-population synthesis model of the Galactic population of X-ray binaries (Iben et al. 1995a,b) to deduce birthrates for interacting binaries; these, combined with estimates of the duration of the X-ray bright phase, have allowed PR02 to make realistic estimates of the relative (Galactic) abundances of high-mass and low-mass X-ray binaries (HMXBs, LMXBs). The abundance of SNRs (both Type II and Ia explosions) was also consistently estimated. From the literature PR02 derived typical spectra for these classes of source. The spectral properties and relative abundances of the various classes of stellar sources determine the composite X-ray spectrum arising from a stellar population of Galactic composition. Therefore, the PR02 \"stellar\" contribution has then no essential degrees of freedom: fixed by the synthetic model of the Galactic population of binaries and by the observed X-ray spectra of the contributing components, it represents the X-ray spectrum emitted by a Galactic mix of HMXBs and LMXBs. As such, the stellar component of the PR02 template spectrum is likely to be appropriate for a quietly star forming galaxy like the Milky Way (see Fig.1-{\\it left}). \\footnote{The diffuse non-stellar part of the PR02 synthetic spectrum has both a thermal and a non-thermal component. The former is related to the SN-powered outgoing galactic wind and is mostly relevant at energies $\\mincir$1 keV. The latter is due to Compton scattering of the SN-accelerated, radio-emitting relativistic electrons off the FIR and CMB radiation fields, and it arguably dominates the spectrum of SFGs at energies of $\\magcir$30 keV (see Persic \\& Rephaeli 2002, 2003). }  The PR02 approach has the flexibility of handling also the extreme cases of 'no ongoing SF' and 'very high ongoing SF' by switching off the HMXB and SNR components and, respectively, the LMXB component (see Fig.1-{\\it right}). By letting the amplitudes of the various spectral components vary (while keeping their profiles fixed) in the spectral fit, one can use the PR02 procedure to determine the SF state of a given galaxy. Thus, this treatment is generally valid and not limited to any particular regime of SF activity.  Based on their survey of galactic X-ray emission mechanisms, PR02 concluded that in the 2-10 keV energy range X-ray binaries of both types, HMXBs and LMXBs, are the most prominent components with the required spectral shapes (see Fig.1-{\\it left}). These have either a power-law (PL) form for HMXBs, or cut-off PL for LMXBs, with observed ranges of spectral parameters (HMXBs: $\\G \\simeq 1.0-1.4$, see White et al. 1983; LMXBs: $\\G \\simeq 0-1.4$, $E_{\\rm c} \\sim 5-10$ keV, see Christian \\& Swank 1997) that provide good fits to the spatially unresolved {\\it BeppoSAX} spectra of the most extensively observed nearby starburst galaxies (SBGs), M~82 and NGC~253 (see PR02), and to the {\\it XMM-Newton} spectra of a number of distant Ultra-Luminous Infra-Red Galaxies (ULIRGs; Franceschini et al. 2003).  From a broader perspective, the template spectrum of PR02 provides a framework that may prove especially useful for interpreting low spatial resolution data on SFGs, either distant (from {\\it Chandra} and {\\it XMM}: e.g., Hornschemeier et al. 2001, 2003; Alexander et al. 2002; Bauer et al. 2002; Franceschini et al. 2003) or nearby (from {\\it ASCA} and {\\it BeppoSAX}: e.g., Cappi et al. 1999; Dahlem et al. 1998; Della Ceca 1996, 1997, 1999; Moran et al. 1999). In particular, this template spectrum would allow the general possibility of measuring the {\\it ongoing} star formation rate (SFR) in galaxies from their X-ray spectra, or perhaps -- for some galaxies --  directly from their X-ray luminosities. The basic notion is that ongoing SFR in a galaxy can be measured based on stellar X-ray sources which are both sufficiently bright for their collective emission to be unambiguously identified, and sufficiently short-lived so that they trace the 'instantaneous' SFR. Of the three main types of stellar galactic X-ray sources,  \\noindent {\\it (i)} SNRs (which are X-ray bright over timescales $t_{\\rm x} \\sim 10^3$ yr) are the evolutionary outcome of massive ($8 \\mincir M/M_\\odot \\mincir 40$) stars that explode on timescales ($5 \\mincir \\tau_{\\rm ev}/{\\rm Myr} \\mincir 50$, see Maeder \\& Meynet 1989) short compared with a typical SB duration ($\\tau_{\\rm SB} \\mincir 100$ Myr). Hence SNRs do trace the istantaneous SFR; however, their relative emission in integrated SFG spectra is quite modest and hard to identify;  \\noindent {\\it (ii)} LMXBs ($t_{\\rm x} \\sim 10^7$ yr) do contribute significantly to the X-ray emission but, due to the long delay between their formation and the onset of their X-ray emission (the donor star has $M \\mincir M_\\odot$), they do not trace the current SFR; finally,  \\noindent {\\it (iii)} HMXBs ($t_{\\rm x} \\sim 10^4$ yr) provide a suitable combination of short delay between binary formation and onset of X-ray emission (the donor star has $M \\magcir 8\\,M_\\odot$) and significant -- sometimes dominant -- relative X-ray emission.  \\noindent Consequently, the measurement of ongoing SFR in galaxies hinges on our ability to separate out the HMXB contribution to the 2-10 keV luminosity, $L_{2-10}$.  In this paper we extend the work of PR02 by examining ways in which $L_{2-10}$ can be used as an astrophysically motivated SFR estimator for SFGs. The possibility of using $L_{2-10}$ as an integral measure of the SFR has already been suggested (e.g.: Bauer et al. 2002; Grimm et al. 2003; Ranalli et al. 2003; Gilfanov et al. 2004). In this paper we propose that by using the HMXB portion of the hard X-ray luminosity, the SFR--$L_{\\rm 2-10}^{\\rm HMXB}$ relation is universal and extends from normal galaxies (low SFR) to very actively starbursting galaxies (very high SFR). Given the ways in which $L_{\\rm 2-10}$ can be contaminated, i.e. systematically from LMXBs at low-SFR regimes (PR02), and occasionally but quite frequently from AGNs at high-SFR regimes (Franceschini et al. 2003), our results will emphasize the potential and the limitations of using -- over $\\sim$5 decades in X-ray luminosity -- the 2-10 keV luminosity as an independent gauge of SFR in SFGs.  The paper is organized as follows. In section 2 we discuss galactic SFR indicators linked to the presence of short-lived massive ($\\magcir 8 M_\\odot$) main-sequence stars. Section 3 describes the sample of SFGs, which comprises objects with SBs of various strengths. In section 4 we discuss the effectiveness of the 2-10 keV luminosity of HMXBs as a SFR indicator. Section 5 summarizes our main results. The conclusions are in section 6.  \\begin{table*} \\caption[] {Data I: The sample of ultra-luminous IR galaxies (ULIRGs)$^{(a)}$.} \\begin{flushleft} \\begin{tabular}{ l l l l l l l l} \\noalign{\\smallskip} \\hline \\hline \\noalign{\\smallskip} Object         & D$^{(b)}$  & $f_{0.5-2}$               &  $f_{2-10}$              & Instr. & $f_{60}$ & $f_{100}$  & Notes$^{(c)}$     \\\\ & [Mpc]      &[$10^{-14}$ erg s$^{-1}$]  &[$10^{-14}$ erg s$^{-1}$] &        &   [Jy]   &  [Jy]      &                   \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} IRAS~05189-2524&  170$^{[7]}$  &     ....          & 360.0$^{[7]}$     & SAX   &  13.25$^{[8]}$ &    11.84$^{[8]}$ & AGN$^{[7]}$        \\\\ IRAS~12112+0305&  291$^{[3]}$  &   1.53$^{[2]}$    &   1.62$^{[2]}$    & XMM   &  8.18$^{[3]}$  &    9.46$^{[3]}$  & SB                 \\\\ Mkn~231        &  169$^{[3]}$  &  29.06$^{[2,11]}$ &  84.54$^{[2,11]}$ & XMM   &   32.0$^{[3]}$ &    30.3$^{[3]}$  & SB + AGN$^{[2]}$   \\\\ Mkn~273        &  151$^{[3]}$  &  17.00$^{[4]}$    &  70.0$^{[4]}$     & ASCA  &  21.7$^{[3]}$  &    21.4$^{[3]}$  & SB + AGN$^{[4]}$   \\\\ IRAS~14348-1447&  330$^{[3]}$  &   3.08$^{[2]}$    &   1.91$^{[2]}$    & XMM   &  6.82$^{[3]}$  &    7.31$^{[3]}$  & SB                 \\\\ IRAS~15250-3609&  213$^{[3]}$  &   2.31$^{[2]}$    &   2.31$^{[2]}$    & XMM   &  7.10$^{[3]}$  &    5.93$^{[3]}$  & SB                 \\\\ Arp~220        &   73$^{[3]}$  &   8.0$^{[6]}$     &  18.0$^{[6]}$     & SAX   & 104.0$^{[3]}$  &   112.0$^{[3]}$  & SB                 \\\\ NGC~6240       &   97$^{[3]}$  &  64.0$^{[5]}$     & 190.0$^{[5,9]}$   & ASCA, SAX & 22.94$^{[3]}$ &   26.49$^{[3]}$& AGN$^{[9]}$       \\\\ IRAS~17208-0014&  170$^{[3]}$  &   7.56$^{[2]}$    &   4.54$^{[2]}$    & XMM   &  9.53$^{[3]}$  &   11.05$^{[3]}$  & SB                 \\\\ IRAS~19254-7245&  246$^{[3]}$  &  10.91$^{[2,10]}$ &  17.58$^{[2,10]}$ & XMM   &  5.5$^{[3]}$   &     5.8$^{[3]}$  & SB + AGN$^{[2]}$   \\\\ &               &                   &                   &       &                &                  & (The Superantennae)\\\\ IRAS~20100-4156&  517$^{[3]}$  &   1.68$^{[2]}$    &   2.24$^{[2]}$    & XMM   &  5.2$^{[3]}$   &     5.2$^{[3]}$  & SB                 \\\\ IRAS~20551-4250&  171$^{[3]}$  &  52.93$^{[2]}$    &  99.56$^{[2]}$    & XMM   & 12.8$^{[3]}$   &    10.0$^{[3]}$  & SB + AGN$^{[2]}$   \\\\ IRAS~22491-1808&  309$^{[3]}$  &   0.61$^{[2]}$    &   0.65$^{[2]}$    & XMM   &  5.54$^{[3]}$  &    4.64$^{[3]}$  & SB                 \\\\ IRAS~23060+0505&  692$^{[3]}$  & 270.00$^{[1]}$    & 250.0$^{[1]}$     & ASCA  &  1.2$^{[3]}$   &     0.8$^{[3]}$  & AGN$^{[1]}$        \\\\ IRAS~23128-5919&  178$^{[3]}$  &  12.04$^{[2]}$    &  21.66$^{[2]}$    & XMM   &  10.8$^{[3]}$  &    11.0$^{[3]}$  & SB + AGN$^{[2]}$   \\\\  \\noalign{\\smallskip} \\hline \\hline \\end{tabular} \\end{flushleft} \\smallskip  $^{(a)}$ References: [1] Brandt et al. 1997; [2] Franceschini et al. 2003; [3] Genzel et al. 1998; [4] Iwasawa 1999; [5] Iwasawa \\& Comastri 1998; [6] Iwasawa et al. 2001; [7] Severgnini et al. 2001; [8] Sanders et al. 2003; [9] Vignati et al. 1999; [10] Braito et al. 2003; [11] Braito et al. 2004.  $^{(b)}$ Distances are taken from Genzel et al. (1998), or are computed consistently assuming H$_0=75$ km s$^{-1}$ Mpc$^{-1}$, $q_0=0.5$ otherwise.  $^{(c)}$ Component(s) dominating the 2-10 keV emission; other name(s).  \\end{table*}  \\begin{table*} \\caption[] { Data II: The sample of local normal and starburst galaxies (SBG sample)$^{(a)}$.} \\begin{flushleft} \\begin{tabular}{ l  l  l  l  l  l  l  l  l  } \\noalign{\\smallskip} \\hline \\hline \\noalign{\\smallskip} Object$^{(b)}$  &   D$^{(c)}$     & $f_{0.5-2}^{(d)}$ &  $f_{2-10}$       & Instr.  & $B_{\\rm T}^{0\\,(e)}$ & $f_{60}$ & $f_{100}$& Notes$^{(f)}$\\\\ & [Mpc]           &[$10^{-12}$ erg/s] &[$10^{-12}$ erg/s] &         &                      &  [Jy]    &  [Jy]    &           \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} NGC~0055$^{\\rm R}$    &  1.3$^{[26]}$ & 1.8$^{[2,20]}$&  0.68$^{[2,20]}$  &ROSAT, ASCA& 7.63 &  77.00$^{[8]}$ & 174.09$^{[8]}$ &             \\\\ NGC~0253$^{\\rm GR}$   &  3.0$^{[26]}$ & 2.5$^{[1,20]}$&  5.0$^{[1,20]}$   &  SAX    & 7.09   & 967.81$^{[8]}$  &1288.15$^{[8]}$ &            \\\\ NGC~0628              &  9.7$^{[26]}$ & ....          &  0.249$^{[15]}$   &  ASCA   & 9.76   &  21.54$^{[8]}$  &  54.45$^{[8]}$ &            \\\\ NGC~0891$^{\\rm R}$    &  9.6$^{[26]}$ & 0.83$^{[20]}$ &  1.9$^{[20]}$     &  ASCA   & 9.37   &  66.46$^{[8]}$  & 172.23$^{[8]}$ &            \\\\ NGC~1569$^{\\rm R}$    &  1.6$^{[26]}$ & 0.54$^{[3]}$  &  0.22$^{[3]}$     &  ASCA   & 9.42   &  54.36$^{[8]}$  &  55.29$^{[8]}$ &            \\\\ NGC~1808$^{\\rm R}$    & 10.8$^{[26]}$ & 0.65$^{[20]}$ &  0.76$^{[20]}$    &  ASCA   &10.43   & 105.55$^{[8]}$  & 141.76$^{[8]}$ &            \\\\ NGC~2146$^{\\rm R}$    & 17.2$^{[26]}$ & 0.82$^{[5]}$  &  1.11$^{[5]}$     &  ASCA   &10.58   & 146.69$^{[8]}$  & 194.05$^{[8]}$ &            \\\\ NGC~2276$^{\\rm R}$    & 36.8$^{[26]}$ & 0.21$^{[20]}$ &  0.44$^{[20]}$    &  ASCA   &11.75   &  14.29$^{[8]}$  &  28.97$^{[8]}$ &            \\\\ NGC~2403$^{\\rm R}$    &  4.2$^{[26]}$ & 1.6$^{[20]}$  &  0.93$^{[20]}$    &  ASCA   & 8.43   &  41.47$^{[8]}$  &  99.13$^{[8]}$ &            \\\\ NGC~2782              & 37.3$^{[26]}$ & 1.3$^{[25]}$  &  ....             &  ROSAT  &12.01   &   9.17$^{[8]}$  &  13.76$^{[8]}$ &            \\\\ NGC~2903$^{\\rm R}$    &  6.3$^{[26]}$ & 0.79$^{[20]}$ &  0.686$^{[15]}$   &  ASCA   & 9.11   &  60.54$^{[8]}$  & 130.43$^{[8]}$ &            \\\\ NGC~3034$^{\\rm GR}$   &  5.2$^{[26]}$ & 5.8$^{[21]}$  & 15.5$^{[21]}$     &  RXTE   & 5.58   &1480.42$^{[8]}$  &1373.69$^{[8]}$ &     = M~82 \\\\ NGC~3079              & 20.4$^{[26]}$ & 4.58$^{[2]+}$  &  0.78$^{[2]}$     &ROSAT, ASCA&10.41 &  50.67$^{[8]}$ & 104.69$^{[8]}$ &            \\\\ NGC~3256$^{\\rm GR}$   & 37.4$^{[26]}$ & 0.69$^{[16]}$ &  0.586$^{[16]}$   &  ASCA   & 8.34   & 102.63$^{[8]}$  & 114.31$^{[8]}$ &            \\\\ NGC~3310              & 18.7$^{[26]}$ & 0.74$^{[20]}$ &  0.21$^{[27]}$    &  ASCA   &10.95   &  34.56$^{[8]}$  &  48.19$^{[8]}$ &            \\\\ NGC~3367$^{\\rm R}$    & 43.6$^{[26]}$ & 0.18$^{[20]}$ &  0.16$^{[20]}$    &  ASCA   &11.92   &   6.44$^{[8]}$  &  13.48$^{[8]}$ &            \\\\ NGC~3556$^{\\rm R}$    & 14.1$^{[26]}$ & 0.44$^{[20]}$ &  0.60$^{[20]}$    &  ASCA   & 9.83   &  32.55$^{[8]}$  &  76.90$^{[8]}$ &    = M~108 \\\\ NGC~3628              &  7.7$^{[26]}$ & 4.16$^{[2]+}$  &  0.98$^{[2]}$     &ROSAT, ASCA & 9.31&  54.80$^{[8]}$& 105.76$^{[8]}$ &             \\\\ Arp~299$^{\\rm GR}$    & 41.6$^{[17]}$ & 0.57$^{[20]}$ &  1.08$^{[20,23]}$ &  ASCA   &11.85   & 113.05$^{[8]}$  & 111.42$^{[8]}$ &= NGC~3690 + IC~694\\\\ &               &               &                   &         &        &                 &                &{\\small AGN at $\\magcir$10 keV}$^{[32]}$\\\\ NGC~4038/39$^{\\rm GR}$& 25.4$^{[26]}$ &0.72$^{[20,23]}$& 0.53$^{[20,23]}$  &  ASCA  &10.62   &  45.16$^{[8]}$  &  87.09$^{[8]}$ & = The Antennae\\\\ NGC~4449$^{\\rm R}$    &  3.0$^{[26]}$ & 0.826$^{[4]}$ &  0.482$^{[4]}$    &  ASCA   & 9.94   &  36.0$^{[11]}$  &  73.0$^{[11]}$ &            \\\\ NGC~4631$^{\\rm R}$    &  6.9$^{[26]}$ & 40.0$^{[2]+}$  &  1.15$^{[2]}$     &ROSAT, ASCA& 8.61 &  85.40$^{[8]}$  & 160.08$^{[8]}$ &           \\\\ NGC~4654$^{\\rm R}$    & 16.8$^{[26]}$ & 0.06$^{[20]}$ &  0.09$^{[20]}$    &  ASCA   &10.75   &  13.39$^{[8]}$  &  37.77$^{[8]}$ &            \\\\ NGC~4666              & 14.1$^{[26]}$ & 0.16$^{[19]}$ &  0.29$^{[19]}$    &  SAX    &10.68   &  37.11$^{[8]}$  &  85.95$^{[8]}$ &            \\\\ NGC~4945$^{\\rm G}$    &  5.2$^{[26]}$ & 1.3$^{[10]}$  &  5.4$^{[10]}$     &  SAX    & 7.43   & 625.46$^{[8]}$  &1329.70$^{[8]}$ & {\\small AGN} \\\\ &               &               &                   &         &        &                 &                & {\\small at $\\magcir$10 keV}$^{[10,31,32,22,24]}$  \\\\ NGC~5236$^{\\rm G}$    &  4.7$^{[26]}$ & 3.5$^{[18]}$  &  4.7$^{[18]}$     & ASCA    & 7.98   & 265.84$^{[8]}$  & 524.09$^{[8]}$ &  = M~83    \\\\ NGC~5253$^{\\rm G}$    &  3.2$^{[26]}$ & 0.32$^{[14]}$ &  ....             &ROSAT    &10.47   &  29.84$^{[8]}$  &  30.08$^{[8]}$ &            \\\\ NGC~5457$^{\\rm R}$    &  5.4$^{[26]}$ & 0.54$^{[20]}$ &  0.68$^{[20]}$    & ASCA    & 8.21   &  88.04$^{[8]}$  & 252.84$^{[8]}$ &  = M~101   \\\\ NGC~6946$^{\\rm R}$    &  5.5$^{[26]}$ & 3.0$^{[20]}$  &  1.2$^{[20]}$     & ASCA    & 7.78   & 129.78$^{[8]}$  & 290.69$^{[8]}$ &            \\\\ NGC~7469$^{\\rm G }$   & 65.2$^{[17]}$ & ....          & 29.8$^{[9]}$      & ASCA    &12.64   &  27.33$^{[8]}$  &  35.16$^{[8]}$ & AGN$^{[28,29,30]}$   \\\\ NGC~7552$^{\\rm G }$   & 19.5$^{[26]}$ & 1.0$^{[13]}$  &  0.36$^{[13]}$    &Einstein &11.13   &  77.37$^{[8]}$  & 102.92$^{[8]}$ &            \\\\ NGC~7679              & 71.0$^{[6]}$ & 3.3$^{[6]}$    &  6.0$^{[6]}$      & SAX     &12.89   &   7.40$^{[8]}$  &  10.71$^{[8]}$ & AGN$^{[6]}$\\\\ IC~342$^{\\rm R}$      &  3.9$^{[26]}$ & 1.8$^{[20]}$  &  11.0$^{[20]}$    &ASCA     & 6.04   & 180.80$^{[8]}$  & 391.66$^{[8]}$ &            \\\\  \\noalign{\\smallskip} \\hline \\hline \\end{tabular} \\end{flushleft} \\smallskip  $^{(a)}$ References: [1] Cappi et al. 1999; [2] Dahlem et al. 1998; [3] Della Ceca et al. 1996; [4]  Della Ceca et al. 1997; [5] Della Ceca et al. 1999; [6] Della Ceca et al. 2001; [7] de Naray et al. 2000; [8] Sanders et al. 2003; [9] Guainazzi et al. 1994; [10] Guainazzi et al. 2000; [11] Hunter et al. 1986; [12] Iwasawa et al. 1993; [13] Maccacaro \\& Perola 1981; [14] Martin \\& Kennicutt 1995; [15] Mizuno et al. 1998; [16] Moran et al. 1999; [17] NED; [18] Okada et al. 1990; [19] Persic et al. 2003; [20] Ranalli et al. 2003; [21] Rephaeli \\& Gruber 2002; [22] Done et al. 2003; [23] Sansom et al. 1996; [24] Schurch et al. 2002; [25] Schulz et al. 1998; [26] Tully 1988; [27] Zezas et al. 1998; [28] Perez-Olea \\& Colina 1996; [29] Nandra et al. 2000; [30] Blustin et al. 2003; [31] Madejski et al. 2000; [32] Della Ceca et al. 2002.  $^{(b)}$ The superscripts G, R indicate whether an object is included in the Genzel et al. (1998) or Ranalli et al. (2003) samples, respectively.  $^{(c)}$ Distances are taken from Tully (1988) if $cz \\leq 3000$ km s$^{-1}$, or are computed consistently assuming H$_0=75$ km s$^{-1}$ Mpc$^{-1}$ otherwise.  $^{(d)}$ Fluxes marked with a cross refer to the 0.1-2.0 keV band.  $^{(e)}$ Blue apparent magnitudes, corrected to face-on and for Galactic absorption, from RC3 (de Vaucouleurs et al. 1991).  $^{(f)}$ Other name(s); spectrally dominating component(s) at energies $\\magcir$10 keV. \\end{table*}  \\begin{table*} \\caption[] { Data III: The sample of {\\it Hubble} Deep Field North galaxies (HDFN sample)$^{(a)}$.} \\begin{flushleft} \\begin{tabular}{ l  l  l  l  l  l  l  l } \\noalign{\\smallskip} \\hline \\hline \\noalign{\\smallskip} Source &  $z$  &$F_{1.4\\, {\\rm GHz}}$& $L_{1.4\\, {\\rm GHz}}$ &      SFR$^{(b)}$    & $f_{2-10}$             &  $L_{2-10}$            &  Instr. \\\\ &       &       [$\\mu$Jy]     &[erg s$^{-1}$Hz$^{-1}$]&[$M_\\odot$ yr$^{-1}$]&[erg s$^{-1}$cm$^{-2}$] &  [erg s$^{-1}$]        &         \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 134    & 0.456 &         210         & $1.03 \\times 10^{30}$ &      25.8           &$~~2.8 \\times 10^{-16}$ & $ 1.38 \\times 10^{41}$ & Chandra \\\\ 136    & 1.219 &         180         & $8.68 \\times 10^{31}$ &     217.1           &$~~1.9 \\times 10^{-16}$ & $ 9.17 \\times 10^{41}$ & Chandra \\\\ 148    & 0.078 &          96         & $1.16 \\times 10^{28}$ &       0.3           & $<4.1 \\times 10^{-17}$ & $<4.95 \\times 10^{38}$ & Chandra \\\\ 188    & 0.410 &          83         & $3.23 \\times 10^{29}$ &       8.1           &$~~5.8 \\times 10^{-17}$ & $ 2.26 \\times 10^{40}$ & Chandra \\\\ 194    & 1.275 &          60         & $3.24 \\times 10^{30}$ &      80.9           &$~~2.0 \\times 10^{-16}$ & $ 1.08 \\times 10^{42}$ & Chandra \\\\ 246    & 0.423 &          36         & $1.50 \\times 10^{29}$ &       3.8           &$~~7.5 \\times 10^{-17}$ & $ 3.13 \\times 10^{40}$ & Chandra \\\\ 278    & 0.232 &         160         & $1.84 \\times 10^{29}$ &       4.6           &$~~1.6 \\times 10^{-16}$ & $ 1.84 \\times 10^{40}$ & Chandra \\\\  \\noalign{\\smallskip} \\hline \\hline \\end{tabular} \\end{flushleft} \\smallskip  $^{(a)}$ All data are taken, or derived, from Ranalli et al. (2003). Distant dependent quantities have been computed assuming H$_0=75$ km s$^{-1}$ and $q_0=0.5$.  $^{(b)}$ The SFR has been computed using Condon's (1992) formula ${\\rm SFR}(>5\\, M_\\odot$) = $L_{1.4\\, {\\rm GHz}} / (4 \\times 10^{28} {\\rm erg~s^{-1} Hz^{-1}} ) \\, M_\\odot {\\rm yr}^{-1}$, which yields the SFR for stars in the mass range $5 \\leq M/M_\\odot \\leq 100$ (assuming a Salpeter-like stellar IMF with $dN/dM \\propto M^{-2.5}$, and a non-thermal radio spectral index $\\alpha=0.8$). \\end{table*} ",
        "conclusions": "Our main suggestion -- and the central theme of investigation -- has been the notion that the 2-10 keV collective emission of HMXBs, $L_{2-10}^{\\rm HMXB}$, is a meaningful gauge of ongoing galactic SFR. As spelled out in Section 4, given the universality of the accretion process onto NS/BH in binaries, the linear dependence of $L_{\\rm FIR}$ on the SFR, the direct link between the FIR and HMXB emissions, a tight relation is predicted between SFR and $L_{2-10}^{\\rm HMXB}$. Whether the latter relation is also linear depends largely on whether the stellar mass range and the shape of the IMF are the same in all galaxies, or vary systematically with the SFR. In the former case we expect the relation to be linear, whereas if the IMF becomes progressively more top-heavy (i.e., the lower mass cut-off becomes progressively higher) with increasing SFR, a non-linear relation is expected.  As we have mentioned, the use of HMXBs as a galactic SFR estimator was already proposed by Grimm et al. (2003) (see also David et al. 1992). Within the general agreement between their background picture and ours, the main result of Grimm et al. is, however, quite different from ours: their resulting relation is non-linear (SFR $\\propto L^{0.6}$) for small SFR and low values of $L_{2-10}$ (i.e., SFR $\\mincir 4.5 \\, M_\\odot {\\rm yr}^{-1}$ and $L_{2-10} \\mincir 2.6 \\times 10^{40} {\\rm erg ~ s}^{-1}$), and linear for higher values of the SFR and $L_{2-10}$. Gilfanov et al. (2004) argue that such non-linearity in the low-SFR limit may be caused by non-Gaussianity of the probability distribution of the integral distribution of discrete sources. We can offer no clear explanation for the discrepancy between our result and theirs. One important difference between our approach and that of Grimm et al. (2003) is their {\\it a priori} minimization of the LMXB contribution because -- as they put it -- `owing to the absence of optical identifications of a donor star in the X-ray binaries detected by {\\it Chandra} in other galaxies, ... there is no obvious way to discriminate the contribution of low-mass X-ray binaries'. Grimm et al. argue that, given the long evolution timescales of LMXBs, in a galaxy the LMXB population should be proportional to the galaxy stellar mass, whereas the HMXB population should be proportional to the SFR, so that the relative importance of LMXBs should be characterized by the ratio of the stellar mass to the SFR. Using dynamical estimates for the stellar masses of galaxies, and SFR derived from a variety of indicators (UV, H$\\alpha$, FIR, radio, that give a wide range of values of the SFR -- see their Table 3), Grimm et al. claimed that they selected only galaxies where HMXB emission is expected to exceed LMXB emission by a factor of $\\magcir$3. For their sample galaxies, consequently, Grimm et al. supposed that the plain 2-10 keV luminosity is a measure of the collective HMXB emission in that band. In contrast, our procedure involves the identification of the HMXB emission from a given SFG by means of a spectral decomposition of the observed X-ray spectrum.  \\begin{figure} \\vspace{4.3cm} \\special{psfile=0500fig7a.ps hoffset=-10 voffset=-53 hscale=30.0 vscale=30.0 angle=0} \\hspace{.5cm} \\special{psfile=0500fig7b.ps hoffset=85 voffset=-53 hscale=30.0 vscale=30.0 angle=0} \\caption{ The distribution of 60$\\mu$-to-100$\\mu$ flux density ratios as a function of the SFR ({\\it left}) and of the SB 2-10 keV to FIR luminosity ratio ({\\it right}). Because $f_{60}/f_{100}$ is an approximate measure of the temperature, these plots show that: {\\it (i)} in spite of a relatively large scatter, a SFR--temperature correlation is present in the data, whereby galaxies with higher (lower) SFR tend to have higher (lower) dust temperature, $T \\sim 50 \\, (25) \\gr$; and {\\it (ii)} the ratio of HMXB 2-10 keV luminosity to dust FIR emission remains roughly constant irrespective of the dust temperature. Symbols are as in Fig.4. } \\end{figure}  The need in our approach to single out $L_{2-10}^{\\rm HMXB}$ rather than $L_{2-10}$ is further spurred by the mix of starbursts of widely varying strengths, from low/moderate (local SBGs) to extreme (ULIRGs). While the FIR emission is mostly related to current SF in both cases (see section 2.1), the 2-10 emission is not in the former case but it is in the latter. Had we restricted to the local SBG sample, we would have found SFR $\\simeq L_{2-10} / (5 \\times 10^{39} {\\rm erg ~s}^{-1}) M_\\odot {\\rm yr}^{-1}$ (see Fig.5-{\\it left}), in agreement with Ranalli et al. (2003). In fact, the 2-10 keV emission from local SBGs has approximately a similar mix of SB-related (HMXB) emission and SB-unrelated (LMXB) emission, so that for this sample the primary correlation, SFR $\\propto L_{2-10}^{\\rm HMXB}$, propagates into SFR $\\propto L_{2-10}$. Some extra scatter plagues the latter relation because the value $f \\sim 0.2$ appearing in the equality $L_{2-10}^{\\rm HMXB} = f \\, L_{2-10}$ most likely has some appreciable scatter within the sample. Only when the SB-dominated ULIRGs and HDFN galaxies are introduced into the plot and are used as calibrators, does the linearity of the SBG SFR--$L_{2-10}$ relation break down (see Fig.5-{\\it left}). The linearity of the relation, extended to the whole luminosity range, is re-established by using the SB-related luminosity $L_{2-10}^{\\rm HMXB}$ (see Fig.5-{\\it right}). [Of course, the calibration of our relation is lower than that of Ranalli et al.'s relation by the factor $f = 0.2$.]  One key assumption underlying our SFR--$L_{2-10}$ relation is that the FIR luminosities -- from which the SFR values are computed -- are really due to the SB and are not contaminated by other contributions in any important way. In order to check this assumption, we convert the observed $f_{60}/f_{100}$ ratios to dust temperatures using the table in Helou et al. (1988). The distribution of 60$\\mu$-to-100$\\mu$ flux density ratios (see Fig.7-{\\it left}) implies a range of estimated dust temperatures \\footnote{Because galaxies have multiple emission components with widely varying parameters, the ratio $f_{60}/f_{100}$ cannot be expected to yield a precise temperature. The uncertainty affecting dust temperature estimates is largely intrinsic, and can be traced back to substantial uncertainties that are essentially of astrophysical origin (e.g., the frequency dependence of the dust emissivity, e.g., Helou et al. 1988).  } $25 \\gr \\mincir T \\mincir 50 \\gr$ (if the dust emissivity is proportional to $\\nu^{-1}$), the lower bound typically referring to more quiescent (i.e., 'normal') galaxies (e.g., NGC~628) and the upper bound to more actively star forming objects (e.g., the well-known SB-dominated galaxy Arp~220; Silva et al. 1998).  This range of values seems realistic for the samples of SFGs considered in this paper, thus confirming the soundness of the assumption that our FIR luminosities are not significantly contaminated by emission unrelated to ongoing SF activity, and verifying earlier suggestions that galaxies with higher SFRs have higher dust temperatures (e.g., David et al. 1992; but there may be exceptions to this trend, see Fig.7-{\\it left}). Notice also that the ratio of the two SB-related luminosities, $L_{2-10}^{\\rm HMXB}/L_{\\rm FIR}$, remains approximately constant when moving from lower temperatures (normal and starburst galaxies) to higher temperatures (ULIRGs), log$\\, (L_{2-10}/L_{\\rm FIR}) \\simeq -4.3 \\pm 0.3$ (see Fig.7-{\\it right}). In particular, note that, when properly accounting for LMXB and AGN emission, SBGs and ULIRGs imply an X-to-FIR ratio consistent with that for the pure SBs arguably represented by the AGN-free ULIRGs.  In the simple scenario, in which the formation of OB stars and of accreting NSs/BHs during episodes of active SF activity only depends on local conditions, the same FIR--X-ray emission relation is expected to hold for local ($z \\mincir 0.1$) and distant ($z \\sim 1$) star-forming galaxies. Indeed, the sample of $z \\sim 1$ HDFN galaxies we have used does follow the SFR-$L_{\\rm x}$ relation established in local galaxies. It should be emphasized that distant galaxies are in a substantially active SF phase: their 2-10 keV emission must then be HMXB-dominated, if no central AGN is contaminating the emission. If an independent estimate of the SFR is also available (e.g., from deep radio observations), the location of the galaxy in the SFR-$L_{2-10}$ plane can be used to evaluate any excess X-ray luminosity which, if detected, would presumably imply the presence of an AGN. For example, if according to eq.(2) a distant galaxy were very overluminous for its SFR (e.g., if $L_{2- 10} \\sim 10^{43}$ erg s$^{-1}$ for SFR $\\sim 100$ $M_\\odot$ yr$^{-1}$), we might suspect a significant AGN contribution to the observed 2-10 keV emission (evolutionary considerations rule out significant LMXB emission). This method, based on joint X-ray and (say) radio observations, could be a very efficient tool to detect AGNs that are buried deeply (i.e., absorbed through $N_{\\rm H} > 10^{23}$ cm$^{-2}$) in the nuclei of distant SFGs and are not detectable in other bands [e.g., the exemplary cases of NGC~6240 (Vignati et al. 1999), NGC~4945 (Guainazzi et al. 2000), Arp~299 (Della Ceca et al. 2000)]. We therefore caution against using $L_{2-10}$ as a gauge of the SFR.  Finally, from Fig.6-({\\it left}) we see that the emission of SFGs at $\\sim$1 keV is also an indicator of the SFR, although the scatter of the SFR$-L_{\\rm x}$ relation is higher in the soft band than in the hard band (see Fig.6, where the equivalent $L_{\\rm FIR}$--$L_{\\rm x}$ relation is shown). Part of the scatter seen in Fig.6-({\\it left}) is due to observational errors -- such as uncertainties in determining the absorption and metallicity of the soft thermal plasma, or occasional differences in the definitions of the soft spectral bands -- but most of it is probably intrinsic. In fact, the soft component which is systematically observed in the spectra of SFGs is interpreted as sub-keV thermal emission originating from SN-powered outgoing galactic winds (e.g., Dahlem et al. 1998; Franceschini et al. 2003). Spatially resolved data (Strickland et al. 2000) and simulations (Strickland \\& Stevens 2000) have shown that this emission occurs at the boundary between the hot, tenuous wind fluid proper and the cool, denser ISM where conditions are optimal for thermal emission of X-rays. This implies that only a small fraction of the wind mass is involved in X-ray emission, and that local conditions (i.e., density, clumpiness, and chemical composition of the ISM) crucially determine the level of this emission. The small mass fraction of thermally X-ray emitting galactic gas is subject to a large scatter among galaxies. These considerations suggest that most of the scatter observed in Fig.6-({\\it left}) is probably intrinsic. Indeed, the pure {\\it thermal} emission does not generally correlate with $L_{\\rm FIR}$ (see Fig.6c of Franceschini et al. 2003), implying that galactic-wind emission is {\\it not} a SFR indicator (in spite of winds being an immediate SB outcome). Then, we suggest that the observed $L_{0.5- 2} - L_{\\rm FIR}$ correlation is not primary and mainly due to SB-powered galactic winds in the SFGs, but that it is probably induced by the portion of the hard spectral component of SFGs showing up in the 0.5-2 keV band. Notice how the $L_{\\rm x} - L_{\\rm FIR}$ relation improves when going from left to right in Fig.6, i.e. when the reference X-ray luminosity changes from 'total soft' through 'total hard' to 'HMXB hard'."
    },
    "0402/astro-ph0402465.txt": {
        "abstract": "We present electro-photometric UBV and high-speed U-band flickering observations of  the recurrent nova T~CrB during a period when its U brightness varies by more than 2 mag. The V band is dominated by the ellipsoidal variability of the red giant, however, the variability of the hot component also causes $\\sim 0.15$ mag variations in V. We define a set of parameters which characterise the flickering. The Fourier spectra of all 27 nights are similar to each other. The power spectral density of the variations has a power law component ($\\propto$f$^{-1.46}$ on average). We do not detect a dependence of the Fourier spectra and autocorrelation function on the brightness of the object. Having subtracted the contribution of the red giant, we show that the flickering amplitude correlates with the average flux of the accreting component. A comparison with CH~Cyg and MWC~560 indicates that the flickering of T~CrB is more stable (at least during the time of our observations), than that in the jet-ejecting symbiotic stars. The data are available in electronic form from the authors. ",
        "introduction": "T~CrB (HD 143454) is an interacting binary star which consists of a red giant and a white dwarf (Selvelli et al. 1992, % Belczy{\\'n}ski \\& Miko{\\l}ajewska 1998, and references therein). The star has undergone two nova eruptions  (Nova CrB 1866, 1946) and is thus classified as a recurrent nova (and, due to the presence of the cool giant, plus emission lines seen at outburst, also as a symbiotic star). The red giant fills Roche lobe, and thus the accretion flow onto the white dwarf (WD) is via L$_{1}$, which is typical for cataclysmic variables. Sharing characteristics of three (partly overlapping) types of interacting binaries, T~CrB is therefore an important object for our understanding of the different processes in interacting binaries.  Stochastic brightness variations (flickering), occurring on time scales of seconds and minutes with amplitudes ranging from a few millimagnitudes up to more than an entire magnitude are a phenomenon typical for cataclysmic variables, and it is rarely observed in symbiotic stars. For example, to-date it is detected in only 8 of the  220 known symbiotics (Dobrzycka et al. 1996; Belczy{\\'n}ski et al 2000, Sokoloski et al. 2001). In T~CrB flickering with amplitude of $\\Delta$U$\\sim$0.1-0.5 mag has been observed on a time scale of minutes (Ianna 1964, Lawrence et al. 1967 Bianchini \\& Middleditch 1976, Walker 1977, Bruch 1980). The flickering amplitude is somewhat smaller in B and V bands (Raikova \\& Antov 1986, Hric et al. 1998). In addition, on some occasions such flickering disappears (Bianchini \\& Middleditch 1976, Oscanian 1983, Miko{\\l}ajewski et al. 1997). In our previous investigation (Zamanov \\& Bruch 1998) we showed that the flickering of T~CrB is indistinguishable from the flickering observed in dwarf novae, in spite of the vast difference in the  geometrical size of the systems.  The exact origin of the stochastic variations is not clear, but they are considered to be a result of accretion onto the WD through a disk. The possible mechanisms include unstable mass transfer, magnetic discharges, turbulence and instability in the boundary layer (e.g. Warner 1995, Bruch 1992).  Here we present new UBV and high-speed flickering observations of T~CrB, estimate the contribution of the red giant, analyse the U band variability, search for relations between the flickering quantities and the brightness of the object, and compare its behaviour with two other symbiotic stars (the \"nanoquasars\" CH~Cyg and MWC~560). ",
        "conclusions": "We have analysed the U band variability of the recurrent nova and symbiotic star T~CrB, and compared its behaviour with two other symbiotic stars  CH~Cyg and MWC~560. During the period of our observations T~CrB brightness varies in between U$=$13 and U$=$10 mag. The analyses we performed show that: (1) the V brightness during the last 22 years is dominated by the ellipsoidal variability of the red giant, however the hot component variability with $\\Delta U \\simgt $2 mag, introduces a shift in V with about 0.15 mag. No signs of variability of the red giant has been detected. (2) the power spectrum of the flickering does not change during our observations, remaining always with slope $\\gamma \\approx -1.5$  in spite of the changes in U. We do not detect changes in the power spectrum like those observed in CH~Cyg; (3) The calculated e-folding time of the ACF also does not show dependence on the changes in U; (4) The flickering amplitude is strongly correlated with the average flux of the hot component. (5) The differences in the dependence of the flickering amplitude between T~CrB, CH~Cyg and MWC~560 could be connected with jet ejection and the possible presence of a magnetic white dwarf in the last two. In general, in T~CrB, we have observed  flickering, which does not change considerably its characteristics (at least during the time of our observations). %. and appears % considerably less unstable than the those % in the jet ejectors CH~Cyg and MWC~560 (at least during the time of our observations).  In the future it would be very interesting to determine the behaviour of the flickering amplitude, ACF, power spectra, etc. of other symbiotic stars with flickering (in particular RS Oph, RT Cru, $o$ Ceti) as well as the flickering  of MWC~560 during phases without outflow, as well the connection of flickering with jet precession. Simultaneous spectral and photometric observations over a wide spectral range from UV to IR  could be very useful  to investigate in detail the flickering behaviour and its connection with accretion disk instabilities and jet ejections."
    },
    "0402/astro-ph0402512_arXiv.txt": {
        "abstract": "We have discovered 16 Type Ia supernovae (SNe~Ia) with the {\\it Hubble Space Telescope (HST)} and have used them to provide the first conclusive evidence for cosmic deceleration that preceded the current epoch of cosmic acceleration.  These objects, discovered during the course of the GOODS ACS Treasury program, include 6 of the 7 highest-redshift SNe~Ia known, all at $z>1.25$, and populate the Hubble diagram in unexplored territory. The luminosity distances to these objects, and to 170 previously reported SNe~Ia, have been determined using empirical relations between light-curve shape and luminosity.  A purely kinematic interpretation of the SN~Ia sample provides evidence at the $>$  99\\% confidence level for a transition from deceleration to acceleration or similarly, strong evidence for a cosmic jerk. Using a simple model of the expansion history, the transition between the two epochs is constrained to be at $z=0.46 \\pm 0.13$.  The data are consistent with the cosmic concordance model of $\\Omega_M \\approx 0.3, \\Omega_\\Lambda \\approx 0.7$ ($\\chi^2_{dof}=1.06$), and are inconsistent with a simple model of evolution or dust as an alternative to dark energy.  For a flat Universe with a cosmological constant, we measure $\\Omega_M=0.29 \\pm ^{0.05}_{0.03}$ (equivalently, $\\Omega_\\Lambda=0.71$).  When combined with external flat-Universe constraints including the cosmic microwave background and large-scale structure, we find $w=-1.02 \\pm ^{0.13}_{0.19}$ (and $w<-0.76$ at the 95\\% confidence level) for an assumed static equation of state of dark energy, $P = w\\rho c^2$.  Joint constraints on both the recent equation of state of dark energy, $w_0$, and its time evolution, $dw/dz$, are a factor of $\\sim 8$ more precise than its first estimate and twice as precise as those without the SNe~Ia discovered with {\\it HST}.  Our constraints are consistent with the static nature of and value of $w$ expected for a cosmological constant (i.e., $w_0 = -1.0$, $dw/dz = 0$), and are inconsistent with very rapid evolution of dark energy.  We address consequences of evolving dark energy for the fate of the Universe. ",
        "introduction": "Observations of type Ia supernovae (SNe~Ia) at redshift $z < 1$ provide startling and puzzling evidence that the expansion of the Universe at the present time appears to be {\\it accelerating}, behavior attributed to ``dark energy'' with negative pressure (Riess et al. 1998; Perlmutter et al. 1999; for reviews, see Riess 2000; Filippenko 2001, 2004; Leibundgut 2001).  Direct evidence comes from the apparent faintness of SNe~Ia at $z \\approx 0.5.$.   Recently expanded samples of SNe~Ia have reinforced the statistical significance of this result (Knop et al. 2003) while others have also extended the SN~Ia sample to $z \\approx 1$ (Tonry et al. 2003; Barris et al. 2004).  Observations of large-scale structure (LSS), when combined with measurements of the characteristic angular size of fluctuations in the cosmic microwave background (CMB), provide independent (though indirect) evidence for a dark-energy component (e.g., Spergel et al. 2003).  An independent, albeit more tentative investigation via the integrated Sachs-Wolfe (ISW) effect also provides evidence for dark energy (Scranton et al. 2003). The magnitude of the observed acceleration was not anticipated by theory and continues to defy a {\\it post facto} explanation.  Candidates for the dark energy include Einstein's cosmological constant $\\Lambda$ (with a phenomenally small value), evolving scalar fields (modern cousins of the inflation field; Caldwell, Dav\\'e, \\& Steinhardt 1998; Peebles \\& Ratra 2002), and a weakening of gravity in our 3 + 1 dimensions by leaking into the higher dimensions required in string theories (Deffayet, Dvali, \\& Gabadadze 2002).  These explanations bear so greatly on fundamental physics that observers have been stimulated to make extraordinary efforts to confirm the initial results on dark energy, test possible sources of error, and extend our empirical knowledge of this newly discovered component of the Universe.  Astrophysical effects could imitate the direct evidence from SNe Ia for an accelerating Universe. A pervasive screen of grey dust could dim SNe~Ia with little telltale reddening (Aguirre 1999a,b).  Luminosity evolution could corrupt the measurements if SNe~Ia at $z \\approx 0.5$ are intrinsically fainter than their low-redshift counterparts.  To date, no evidence for an astrophysical origin of the apparent faintness of SNe Ia has been found (Riess 2000; Coil et al. 2001; Leibundgut 2001; Sullivan et al. 2003). However, given the significance of the putative dark energy and the unique ability of SNe~Ia to illuminate it, we need a more definitive test of the hypothesis that supernovae at $z\\sim 0.5$ are intrinsically dimmer, or dimmed by absorption.  If cosmic acceleration is the reason why SNe Ia are dimmer at $z \\sim 0.5$, then we expect cosmic deceleration at $z>1$ to reverse the sign of the observed effect.  The combination of recent acceleration and past deceleration is a clear signature of a mixed dark-matter and dark-energy Universe and one which is readily distinguishable from simple astrophysical dimming (Filippenko \\& Riess 2001).  Furthermore, assuming SNe~Ia at $z > 1$ continue to trace the cosmological world model, measurements of SNe~Ia in the next redshift octave provide the unique ability to discriminate between a static and evolving dark-energy equation of state.  This would provide a vital clue to distinguish a cosmological constant from other forms of dark energy that change with time.  Ground-based efforts to look for past deceleration with SNe~Ia have offered hints of the effect, but ultimately they have suffered from insufficient signal-to-noise ratios (Tonry et al. 2003; Barris et al. 2004).  Discovering, confirming, and then monitoring transients at $I \\approx 25$ mag on the bright sky is challenging even with the largest telescopes and the best conditions.  A single SN~Ia at $z \\approx 1.7$, SN 1997ff, discovered with WFPC2 on the {\\it Hubble Space Telescope (HST)} (Gilliland, Nugent, \\& Phillips 1999), provided a hint of past deceleration; however, inferences drawn from a single SN~Ia, while plausible, are not robust (Riess et al. 2001; Ben\\'\\i tez et al. 2002; Mortsell, Gunnarsson, \\& Goobar 2001).  To study the early expansion history of the Universe, we initiated the first systematic, space-based search and follow-up effort to collect SNe~Ia at $z > 1$, carried out in conjunction with the Great Observatories Origins Deep Survey (GOODS) Treasury program (Giavalisco et al. 2003) conducted with the Advanced Camera for Surveys (ACS) aboard {\\it HST}. (The ability to detect SNe at $z>1$ with the Space Telescope was an application first envisioned during its planning; Tammann 1977, Colgate 1979).   A separate ``piggyback'' program was utilized to obtain target of opportunity (ToO) follow-up {\\it HST} observations of the SNe~Ia with ACS and NICMOS (the Near-Infrared Camera and Multi-Object Spectrograph).  Elsewhere we present a color-based method for discrimination of SNe~Ia at $z > 1$ from other transients (Riess et al. 2003) and the full harvest of the SN survey (Strolger et al. 2004).  We present the follow-up spectroscopy and photometry of 16 SNe~Ia in \\S 2, light-curve analysis in \\S 3, cosmological tests and constraints in \\S 4, and a discussion and summary in \\S 5 and \\S 6, respectively. ",
        "conclusions": "We have conducted the first space-based SN search and follow-up campaign using the ACS on board {\\it HST}.  The search parameters and the full list of 42 new SNe are provided elsewhere (Strolger et al. 2004).  We reviewed the sample of SNe~Ia harvested from the survey and examined its cosmological significance.  The key results can be summarized as follows.  (1) We obtained multi-color light curves and spectroscopic redshifts for 16 new SNe~Ia which uniformly sample the redshift range $0.2 < z < 1.6$.  Twelve of these are classified by their spectra, 2 from their red, early-type host galaxies, and 2 by photometric diagnostics.  Three of the SN spectra are at the highest redshifts yet observed for SNe.  Six of the SNe~Ia are among the seven highest-redshift known; all are at $z>1.25$. These data provide a robust extension of the Hubble diagram to $1 < z < 1.6$.  (2) Utilizing a simple kinematic description of the magnitude-redshift data, we find that the SNe~Ia favor recent acceleration and past deceleration at the 99.2\\% confidence level.  An alternate kinematic parameterization requires a positive jerk (third derivative of the scale factor).  The best-fit redshift of the transition between these kinematic phases is $z=0.46 \\pm 0.13$, although the precise value depends on the kinematic model employed.  (3) We have compared the goodness-of-fit of cosmological models and simple models of astrophysical dimming. The ``gold'' sample of 157 SNe~Ia is consistent with the ``cosmic concordance'' model ($\\Omega_M=0.3, \\Omega_\\Lambda=0.7$) with $\\chi^2_{dof}=1.06$.  The data reject at high-confidence simple, monotonic models of astrophysical dimming which are tuned to mimic the evidence for acceleration at $z \\approx 0.5$.  These models include either a universe filled with gray dust at high redshift, or luminosity evolution $\\propto z$.  More complex parameterizations of astrophysical dimming which peak at $z \\approx 0.5$ and dissipate at $z>1$ remain consistent with the SN data (but appear unattractive on other grounds).  (4) For a flat Universe with a cosmological constant, we measure $\\Omega_M=0.29 \\pm ^{0.05}_{0.03}$ (equivalently, $\\Omega_\\Lambda=0.71$).  When combined with external flat-Universe constraints including the CMB and LSS, we find for the dark-energy equation of state $w=-1.02 \\pm ^{0.13}_{0.19}$ (and $w<-0.76$ at the 95\\% confidence level) for an assumed static equation of state of dark energy, $P = w\\rho c^2$.  (5) Joint constraints on both the recent equation of state of dark energy and its time evolution are a factor of $\\sim 8$ more precise than its first estimate and twice more precise than those derived without the SNe~Ia discovered by {\\it HST}.  Both of these dark energy properties are consistent with a cosmological constant (i.e., with $w_0=-1.0$, $w'=0$) and are inconsistent with very rapid evolution of dark energy (i.e., $\\vert w' \\vert > $ a few).  The absence of rapid evolution places constraints on the time in which a simple scalar field could evolve to recollapse the Universe.  Specifically, the timescale to a potential recollapse is larger than $\\sim$30 Gyr. If dark energy is evolving towards more negative $w$, we cannot place any meaningful limit on the minimum time to a (speculative) Big Rip.  \\bigskip \\medskip  We wish to thank Richard Hook, Swara Ravindranath, Tomas Dahlen, Peter Garnavich, Duilia de Mello, Ed Taylor, Soo Kim, Rafal Idzi, Carl Biagetti, Lexi Moustakas, Marin Richardson, Vicki Laidler, Ann Hornschmeier, Ray Lucas, Norman Grogin, Claudia Kretchmer, Brian Schmidt, Stephane Blondin, and Anton Koekemoer for their help in the supernova search.  We are grateful to Dorothy Fraquelli, Sid Parsons, Al Holm, Tracy Ellis, Richard Arquilla, and Mark Kochte for their help in assuring rapid delivery of the data.  We thank Tom Matheson, Dan Stern, Hy Spinrad, Piero Rosati, Mario Nonino, Alice Shapley, Max Pettini and Dawn Erb for their efforts to obtain redshifts of some SN host galaxies.  We appreciate the guidance of Anton Koekemoer and Eddie Bergeron. Partly based on observations collected at the European Southern Observatory, Chile (Prog. Nr. 70.A-0497). Financial support for this work was provided by NASA through programs GO-9352 and GO-9583 from the Space Telescope Science Institute, which is operated by AURA, Inc., under NASA contract NAS 5-26555. Some of 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 NASA; the Observatory was made possible by the generous financial support of the W. M. Keck Foundation.  \\appendix"
    },
    "0402/astro-ph0402038_arXiv.txt": {
        "abstract": "We discuss star formation in the turbulent interstellar medium. We argue that morphological appearance and dynamical evolution of the gas is primarily determined by supersonic turbulence, and that stars form via a process we call gravoturbulent fragmentation. Turbulence that is dominated by large-scale shocks or is free to decay leads to an efficient, clustered, and synchronized mode of star formation. On the other hand, when turbulence carries most of its energy on very small scales star formation is inefficient and biased towards single objects.  The fact that Galactic molecular clouds are highly filamentary can be explained by a combination of compressional flows and shear. Some filaments may accumulate sufficient mass and density to become gravitationally unstable and form stars. This is observed in the Taurus molecular cloud. Timescales and spatial distribution of protostars are well explained by the linear theory of gravitational fragmentation of filaments.  The dynamical evolution, especially at late times, and the final mass distribution strongly depend on the global properties of the turbulence. In dense embedded clusters mutual protostellar interactions and competition for the available mass reservoir lead to considerable stochastic variations between the mass growth histories of individual stars. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0402/astro-ph0402662_arXiv.txt": {
        "abstract": "We present the first numerical model of the magnetohydrodynamical cosmic-ray (CR) driven dynamo of the type proposed by Parker (1992). The driving force of the amplification process comes from CRs injected into  the  galactic disk in randomly distributed spherical regions representing  supernova  remnants. The underlying disk is differentially rotating. An explicit resistivity is responsible for the dissipation   of the small-scale  magnetic field component. We obtain   amplification of  the large-scale magnetic  on a timescale 250 Myr. ",
        "introduction": "In 1992 Parker discussed the possibility of a new kind of galactic dynamo driven  by galactic  CRs accelerated in supernova remnants.  This dynamo contains a network of interacting forces: the buoyancy force of CRs,  the Coriolis force, the differential rotation  and magnetic reconnection. Parker estimated that such a dynamo  is able to amplify the large scale magnetic field on timescales of the order of $10^8 {\\rm yr}$.   Over the last decade we have investigated the different physical properties and consequences  of Parker's  idea and scenario by means of analytical calculations and numerical simulations (see eg. Hanasz \\&  Lesch 1998, Lesch \\& Hanasz  2003 and references therein). Here we present the first complete magnetohydrodynamical three-dimensional simulation  including the full network of  relevant interacting mechanisms.  It is the aim of our contribution to show that Parker's CR driven dynamo indeed acts efficiently  on  timescales comparable with the disk rotation time. In the next two Sections we  describe the physical elements of the model and the system of equations used in numerical simulations. Section 4 presents the numerical setup, Sections 5 and 6  inform the reader about the results on the  structure of the  interstellar medium including CRs and magnetic fields,  the strength of the amplified magnetic field and the  spatial structure of the  mean magnetic field. We  summarize our results very briefly in Section 7. ",
        "conclusions": "We have described the first numerical experiment in which the effect of amplification of the large scale galactic magnetic field was achieved by the (1) continuous (although intermittent in space and time) supply of CRs into the interstellar medium,  (2) shearing motions due to differential rotation and (3) the presence of an  explicit resistivity of the medium.   We observed in our experiment the growth of magnetic energy by seven orders of magnitude and the growth of magnetic flux by a factor of 1300 in 2150 Myr of the system evolution. We found that the large scale magnetic field grows on a timescale ~ 250 Myr, which is close to the period of galactic rotation.   Therefore the galactic dynamo driven by CRs  appears to work very efficiently, as it was suggested by Parker (1992). It is a matter of future work to verify whether the presented model is a fast dynamo, i.e. whether it works with a similar efficiency in the limit of vanishing resistivity."
    },
    "0402/astro-ph0402381_arXiv.txt": {
        "abstract": "Ever since the pioneering study of Spitzer, it has been widely recognized that grains play an important role in the heating and cooling of photo-ionized environments. This includes the diffuse ISM, as well as H\\,{\\sc ii} regions, planetary nebulae, and photo-dissociation regions. A detailed code is necessary to model grains in a photo-ionized medium since the interactions of grains with their environment include a host of microphysical processes. In this paper we will use the spectral synthesis code Cloudy for this purpose. A comprehensive upgrade of the grain model has been recently incorporated into Cloudy. One of these upgrades is the newly developed hybrid grain charge model. This model allows discrete charge states of very small grains to be modelled accurately while simultaneously avoiding the overhead of fully resolving the charge distribution of large grains, thus making the model both accurate and computationally efficient. A comprehensive comparison with the fully resolved charge state models of Weingartner \\& Draine (2001a) shows that the agreement is very satisfactory for realistic size distributions. The effect of the grain size distribution on the line emission from photo-ionized regions is studied by taking standard models for an H\\,{\\sc ii} region and a planetary nebula and adding a dust component to the models with varying grain size distributions. A comparison of the models shows that varying the size distribution has a dramatic effect on the emitted spectrum. The strongest enhancement is always found in optical/UV lines of the highest ionization stages present in the spectrum (with factors up to 2.5 -- 4), while the strongest decrease is typically found in optical/UV lines of low ionization lines or infrared fine-structure lines of low/intermediate ionization stages (with reductions up to 10 -- 25\\%). Changing the grain size distribution also affects the ionization balance, and can affect resonance lines which are very sensitive to changes in the background opacity. All these results clearly demonstrate that the grain size distribution is an important parameter in photo-ionization models. ",
        "introduction": "Grains are ubiquitous in the interstellar medium (ISM), and they can be detected either directly through their far-infrared emission or indirectly through extinction or polarisation studies. Despite the vast number of observations, many questions regarding grain composition and grain physics remain unanswered. Further study is therefore required, and detailed models are needed to interpret the results. Ever since the pioneering study of Spitzer (1948), it has been widely recognized that grains play an important role in the heating and cooling of the diffuse ISM (see also the more recent studies by Bakes \\& Tielens 1994, and Weingartner \\& Draine 2001a, hereafter WD). Grains also play an important role in the physics of H\\,{\\sc ii} regions and planetary nebulae (PNe; e.g., Maciel \\& Pottasch 1982, Baldwin et al.\\ 1991, hereafter BFM, Borkowski \\& Harrington 1991, Ercolano et al. 2003) and photo-dissociation regions (PDR's; e.g., Tielens \\& Hollenbach 1985).  The interactions of grains with their environment include a host of microphysical processes, and their importance and effects can only be judged by including all of these processes in a self-consistent manner. This can, in turn, only be done with a complete simulation of the environment. In this paper we use the spectral synthesis code Cloudy for this purpose. Cloudy is a well known and widely used photo-ionization code. This code is not only useful for modelling fully ionized regions, but calculations can also be continued into the PDR. In order to make the models realistic, the presence of a detailed grain model is usually required. The first grain model was introduced into Cloudy in 1990 to facilitate more accurate modelling of the Orion nebula (for a detailed description see BFM). In subsequent years, this model has undergone some revisions and extensions, but remained largely the same.  In the last couple of years, Cloudy has undergone several major upgrades, described in Ferland (2000a), Ferland (2000b), and van Hoof et al.\\ (2000b). This includes a comprehensive upgrade of the grain model. The latter was necessary for two reasons. First, the discovery of crystalline silicates in stellar outflows (e.g., Waters et al.\\ 1996), and other detailed observations of grain emission features by the Infrared Space Observatory (ISO), meant that the code had to become much more flexible to allow such materials to be included in the modelling. Second, even before the ISO mission it had already become clear that the photo-electric heating and collisional cooling of the gas surrounding the grains is dominated by very small grains (possibly consisting of polycyclic aromatic hydrocarbons or PAH's). The physics of very small grains could not be modelled very accurately with the original grain model. In view of these facts we have undertaken a comprehensive upgrade of the grain model in Cloudy. The two main aims were to make the code more flexible and versatile, and to make the modelling results more realistic (e.g., by improving the treatment of grain charging, the photo-electric effect, and stochastic heating). The new grain model has been introduced in version 96 of Cloudy.  We have used a three-pronged approach to improve the grain model in Cloudy. First we introduced a Mie code for spherical particles, which allows the user to use arbitrary grain materials and resolve any grain size distribution of their choosing to arbitrary precision. The latter is very important since most grain properties depend strongly (and more importantly non-linearly) on size. This upgrade, briefly described in Section~\\ref{resolv}, also opened the way for two other major improvements. First, it enabled the accurate modelling of stochastic heating effects for arbitrary grain materials and size distributions. Second, it enabled a much more realistic modelling of grain charging, photo-electric heating, and collisional cooling by the grains, as described in Section~\\ref{physics}. For this purpose we have developed a completely new grain charge model, which we call the hybrid model. This is described in detail in Section~\\ref{hybrid}.  In this paper we will study photo-electric heating by grains in photo-ionized environments in detail. In particular, we will study the effect that the distribution of grain sizes has on the relative intensities of emission lines. We will show that this effect is nothing short of dramatic, making the grain size distribution an important parameter in the modelling of photo-ionized regions such as H\\,{\\sc ii} regions and planetary nebulae. This will be described in Section~\\ref{heating}. Our conclusions will be summarized in Section~\\ref{conclusions}. ",
        "conclusions": ""
    },
    "0402/astro-ph0402348_arXiv.txt": {
        "abstract": "We present an analysis of sixteen galaxy clusters, one group and one galaxy drawn from the \\textit{Chandra} X-ray Data Archive. These systems possess prominent X-ray surface brightness depressions associated with cavities or bubbles that were created by interactions between powerful radio sources and the surrounding hot gas.  The central galaxies in these systems harbor radio sources with luminosities ranging between $\\sim 2\\times 10^{38}-7\\times 10^{44}~\\mbox{ergs s$^{-1}$}$. The cavities have an average radius of $\\sim 10~\\mbox{kpc}$, and they lie at an average projected distance of $\\sim 20~\\mbox{kpc}$ from the central galaxy.  The minimum energy associated with the cavities ranges from $pV \\sim 10^{55}~\\mbox{ergs}$ in galaxies, groups, and poor clusters to $pV \\sim 10^{60}~\\mbox{ergs}$ in rich clusters.  We evaluate the hypothesis that cooling in the hot gas can be quenched by energy injected into the surrounding gas by the rising bubbles. We find that the instantaneous mechanical luminosities required to offset cooling range between $1 pV$ and $20 pV$ per cavity. Nearly half of the systems in this study may have instantaneous mechanical luminosities large enough to balance cooling, at least for a short period of time, if the cavities are filled with a relativistic gas. We find a trend or upper envelope in the distribution of central X-ray luminosity versus instantaneous mechanical luminosity with the sense that the most powerful cavities are found in the most X-ray--luminous systems. Such a trend would be expected if many of these systems produce bubbles at a rate that scales in proportion to the cooling rate of the surrounding gas.  Finally, we use the X-ray cavities to measure the mechanical power of radio sources over six decades of radio luminosity, independently of the radio properties themselves.  We find that the ratio of the instantaneous mechanical (kinetic) luminosity to the 1.4~GHz synchrotron luminosity ranges from a few to roughly a thousand. This wide range implies that the 1.4 GHz synchrotron luminosity is an unreliable gauge of the mechanical power of radio sources. ",
        "introduction": "The cooling time of the intracluster gas in the cores of many galaxy clusters is shorter than 1~Gyr.  In the absence of heating, a ``cooling flow'' \\citep{fabi94} is established, in which the gas cools below X-ray temperatures and accretes onto the central cluster galaxy, where it accumulates in molecular clouds and forms stars. {\\it Chandra} images of cooling flow clusters have confirmed the existence of inwardly decreasing temperature gradients and short central cooling times, which are the distinguishing characteristics of a cooling flow. However, moderate-resolution \\textit{Chandra} and \\textit{ASCA} spectra and high resolution \\textit{XMM-Newton} spectra \\citep[e.g.\\ ][]{maki01,pete01,tamu01, kaast04} do not show the expected signatures of cooling below 2~keV,  reported to exist in lower resolution data from the \\textit{Einstein} and \\textit{ROSAT} observatories. This discrepancy would be difficult to understand unless the normal signatures of cooling below 2 keV are somehow suppressed, or if cooling is indeed occurring but at rates that are generally factors of $5-10$ lower than expected \\citep[e.g.\\ ][]{mole01,boeh02,pete03}. Several scenarios have been suggested that may suppress the cooling flux at low energies, and yet maintain the large cooling rates. The list includes differential absorption, efficient mixing, or inhomogeneous metallicity distributions \\citep[e.g.\\ ][]{fabi01}. However, these scenarios, as clever as they are, may require fine-tuning and are otherwise difficult to prove observationally. In any case, they leave the question of the repository for most of the cooling gas unanswered.  \\begin{deluxetable*}{clccccc} \\tabletypesize{\\scriptsize} \\tablecaption{Radio properties. \\label{tbl-cluster-radio}} \\tablewidth{400pt} \\tablehead{ \\colhead{} & \\colhead{} & \\colhead{$\\sigma$ (ref)\\tablenotemark{a}} & \\colhead{$S_{1400}$}   & \\colhead{} & \\colhead{$P_{1400}$} & \\colhead{$L_{\\rm{radio}}$} \\\\ \\colhead{System\\phn} & \\colhead{z} & \\colhead{(km s$^{-1}$)} & \\colhead{(Jy)} & \\colhead{$\\alpha$ (ref)\\tablenotemark{b}} &\t\\colhead{($10^{24}$ W Hz$^{-1}$)} & \\colhead{($10^{42}$ ergs s$^{-1}$)} } \\startdata Cygnus A    & 0.056  & \\nodata         & 1598$\\pm$41\\phn\\phn       & 0.7\\phm{2} (4)\\phm{22222,,,}         & 11800$\\pm$300\\phn\\phn     & 700$\\pm$20\\phn       \\\\ Hydra A     & 0.052  & \\phn \\phn \\phn322$\\pm$20 (8)\\phn  & 40.8$\\pm$1.3\\phn      & 0.92 (4,13)\\phm{222,,}     & 261$\\pm$8\\phn\\phn         & 19.9$\\pm$0.6\\phn     \\\\ A2597       & 0.085  & \\phn \\phn \\phn224$\\pm$19 (16) & 1.875$\\pm$0.056   & 1.35 (11)\\phm{2222,,,}       & 35$\\pm$1\\phn          & 6.8$\\pm$0.2         \\\\ MKW 3S      & 0.045  & \\nodata         & 115.0$\\pm$3.9\\phn\\phn     & 2.3\\phm{2} (4,13,17)\\phm{2,}   & 0.58$\\pm$0.02     & 3.6$\\pm$0.1        \\\\ A2052       & 0.035  & \\phn \\phn \\phn253$\\pm$12 (19) & 5.50$\\pm$0.21     & 1.2\\phn (1,4,13)\\phm{22,}    & 15.7$\\pm$0.6\\phn      & 2.08$\\pm$0.08       \\\\ A133        & 0.060  & \\nodata         & 0.167$\\pm$0.006   & 1.9\\phn (14,15)\\phm{22,,}     & 1.54$\\pm$0.06     & 1.94$\\pm$0.07       \\\\ A4059       & 0.048  & \\phn \\phn \\phn296$\\pm$49 (6)\\phn  & 1.284$\\pm$0.043   & 1.43 (13,18,21)\\phm{,} & 6.96$\\pm$0.2\\phn      & 1.73$\\pm$0.06       \\\\ A2199       & 0.030  & \\phn \\phn \\phn295$\\pm$6\\phm{9} (7)\\phn   & 3.58$\\pm$0.12     & 1.37 (1,4,17)\\phm{22,}   & 7.43$\\pm$0.2\\phn      & 1.56$\\pm$0.05       \\\\ Perseus     & 0.018  & \\phn \\phn \\phn246$\\pm$10 (8)\\phn  & 22.83$\\pm$0.68\\phn    & 1.0\\phn (10)\\phm{2222,,,}        & 16.5$\\pm$0.5\\phn      & 1.42$\\pm$0.04      \\\\ RBS 797     & 0.350  & \\nodata         & 0.0217$\\pm$0.0008 & \\nodata         & 9.0$\\pm$0.3       & 0.78$\\pm$0.03     \\\\ A1795       & 0.063  & \\phn \\phn \\phn297$\\pm$12 (3)\\phn  & 0.925$\\pm$0.028   & 0.98 (1,4,13)\\phm{22,}   & 9.0$\\pm$0.3       & 0.75$\\pm$0.02      \\\\ M87         & 0.0042 & \\phn \\phn \\phn330$\\pm$5\\phm{2} (2)\\phn   & 138.5$\\pm$4.9\\phn\\phn     & 0.81 (4,9,13)\\phm{22,}   & 5.48$\\pm$0.2\\phn      & 0.36$\\pm$0.01       \\\\ Centaurus   & 0.011  & \\phn \\phn \\phn256$\\pm$11 (5)\\phn  & 3.8            & 0.7\\phn (4,9,14,20) & 1.1            & 0.060            \\\\ A478        & 0.081  & \\nodata         & 0.0369$\\pm$0.0015 & \\nodata         & 0.60$\\pm$0.02     & 0.052$\\pm$0.002    \\\\ M84         & 0.0035 & \\phn \\phn \\phn278$\\pm$4\\phm{2} (8)\\phn  & 6.00$\\pm$0.15     & 0.63  (4,9,17)\\phm{22,}  & 0.162$\\pm$0.004   & 0.0089$\\pm$0.0002   \\\\ 2A 0335+096 & 0.035  & \\nodata         & 0.0367$\\pm$0.0018 & 0.9\\phn (12)\\phm{2222,,,}        & 0.104$\\pm$0.005   & 0.0077$\\pm$0.0004 \\\\ A262        & 0.016  & \\nodata         & 0.0657$\\pm$0.0023 & 0.6\\phn (4)\\phm{22222,,,}         & 0.039$\\pm$0.001   & 0.00210$\\pm$0.00007 \\\\ HCG 62      & 0.014  & \\nodata         & 0.0049$\\pm$0.0005 & \\nodata         & 0.0021$\\pm$0.0002 & 0.00018$\\pm$0.00002 \\\\ \\enddata  \\tablenotetext{a}{When no velocity dispersion was available, $\\left< \\sigma \\right> = 280$ km s$^{-1}$ was adopted. References are in parentheses.} \\tablenotetext{b}{The spectral index is defined so that $S \\sim \\nu^{-\\alpha}$. When no spectral index was available, $\\alpha = 1$ was adopted. References are in parentheses.} \\tablerefs{ (1) Becker, White, \\& Edwards 1991; (2) Bender, Saglia, \\& Gerhard 1994; (3) Blakeslee \\& Tonry 1992; (4) Burbidge \\& Crowne 1979; (5) Carollo, Danzinger, \\& Buson 1993; (6) Carter et al. 1985; (7) Fisher, Illingworth, \\& Franx 1995; (8) Heckman et al. 1985; (9) Kuehr et al. 1981; (10) Pedlar et al. 1990; (11) Sarazin et al. 1995; (12) Sarazin, Baum, \\& O'Dea 1995; (13) Slee 1995; (14) Slee \\& Siegman 1988; (15) Slee et al. 2001; (16) Smith, Heckman, \\& Illingworth 1990; (17) Spindrad et al. 1985; (18) Taylor, Barton, \\& Ge 1994; (19) Tonry 1985; (20) Wright et al. 1994; (21) Wright et al. 1996} \\end{deluxetable*}  The more appealing interpretation, which we address in this paper, posits that radiation losses are being balanced, or nearly so, by heating, implying that the large cooling rates of the last decade were indeed overestimated.   This suggestion has its own difficulties. Maintaining gas with a cooling time approaching 100 Myr at keV temperatures almost certainly requires the existence of one or more heating mechanisms operating in a self-regulating feedback loop. One such mechanism, thermal conduction from the hot outer layers of clusters, may be energetically feasible, in some instances \\citep[e.g.\\ ][]{tuck83,bert86,zaka03,voig02,voig04}. However, it generally requires fine tuning and can be unstable \\citep{bregdav,soke03}. Moreover, conduction operating alone at even the Spitzer rate cannot offset radiation losses \\citep{voig02,wise03} in all clusters, and is therefore unlikely to provide a general solution to the heating problem (recent simulations by Dolag et al.\\ 2004 support this conclusion). Additional heat sources, such as cosmic rays \\citep{boeh88,loew91}, and supernova explosions \\citep{mcna04} may contribute to heating the gas.  Nevertheless, these mechanisms are generally incapable of balancing radiation losses.  In this paper,  we evaluate whether the mechanical energy generated by active galactic nuclei (AGN) can balance radiation losses in cluster cores.  This possibility, which has a substantial legacy in the literature \\citep[e.g.\\ ][]{tabo93,binn95,tuck97,ciot01,sok01}, has been rejuvenated by the crisp, new {\\it Chandra} images of clusters showing the keV gas being displaced by radio sources harbored by central cluster galaxies.  The now ubiquitous signature of these interactions are X-ray surface brightness depressions projected on the radio lobe emission at 1.4 GHz, as is seen in Perseus \\citep{boeh93,fabi00,schm02,fabi02b,fabi03a,fabi03b}, Cygnus A \\citep{carilli}, and Hydra~A \\citep{mcna00,davi01,nuls02}. The displacement of the gas creates a low-density, rising bubble in pressure balance with the surrounding medium. X-ray surface brightness depressions that have no obvious association with the bright radio emission at 1.4 GHz, the so-called ghost cavities, have also been found, such as those in Abell~2597 \\citep{mcna01} and the outer depressions in Perseus \\citep{fabi00}.  These depressions are thought to have been created by interactions that occurred in the more distant past, but whose radio emission has faded over time.  This general scenario has been modeled theoretically using a variety of hydrodynamical, magnetohydrodynamical, and analytical techniques, which have successfully reproduced the gross characteristics of the cavities \\citep[e.g.\\ ][]{gull73,chur01,brug01,quil01,brig02,brug02b,reyn02,brug03,ddy03,bass03,binn03,kais03,math03,omma04,robi04}. The models do not, however, predict with any certainty the amount of mechanical energy provided by radio sources of a given luminosity, nor their frequency of recurrence. Whereas the cavities in several individual objects, such as Perseus \\citep{fabi03a} and Hydra A \\citep{mcna00}, contain enough enthalpy to balance cooling, at least for a short period of time, a systematic survey of cavities in systems with a broad range of properties is required to  determine whether this is generally true. In this paper we address this question by setting observational limits on the energetics and ages of cavities in 18 systems taken from the \\textit{Chandra} Data Archive.  We adopt $H_{0} = 70~\\mbox{km s$^{-1}$ Mpc$^{-1}$}$, $\\Omega_{M} = 0.3$, and $\\Omega_{\\Lambda} = 0.7$ in all calculations throughout this paper. ",
        "conclusions": "\\subsection{Trends with Radio Luminosity}  In Figure~\\ref{fig-Lr_Lm}, we present two plots showing the mechanical luminosity versus the total radio luminosity (\\emph{left}) and the total radio luminosity versus the ratio of the mechanical luminosity to the monochromatic, 1.4~GHz radio luminosity (\\emph{right}). In each plot we distinguish between radio-filled and ghost cavities, shown with filled symbols and empty symbols respectively.  The ``error bars'' for each point reflect the range of instantaneous mechanical luminosity implied by the range in possible ages. The data are taken or are derived from Tables~\\ref{tbl-cluster-radio} and \\ref{tbl-bubbles}.  The left-hand panel of Figure~\\ref{fig-Lr_Lm} shows a trend between the radio luminosity  and mechanical luminosity, with the sense that more luminous radio sources tend toward larger mechanical luminosities.   This trend seems to be shared by both the radio-filled cavities and the ghost cavities, in spite of the use of the current central radio power for both the filled cavities and the ghosts, to which the current central source may be unrelated. No segregation by FOM is seen. The relation between the two luminosities appears to be roughly a power law. To quantify this relation, we used a linear least-squares fit to the logarithms of the data, with errors in mechanical luminosity given by the extreme values for each system. We show in Figure~\\ref{fig-Lr_Lm} the best-fit line for the entire sample (\\emph{dashed line}), given by \\begin{equation} L_{\\rm{mech}} = 10^{25\\pm 3} \\left( L_{\\rm{radio}} \\right)^{0.44\\pm 0.06}, \\end{equation} and for the radio-filled cavities only (\\emph{dotted line}), given by \\begin{equation} L_{\\rm{mech}} = 10^{18\\pm 4} \\left( L_{\\rm{radio}} \\right)^{0.6\\pm 0.1}. \\end{equation} In both cases, the mechanical luminosity scales as the radio luminosity to approximately the one-half power over six decades of radio power, albeit with large scatter.  The relative contribution of cosmic scatter and observational uncertainty is hard to judge without precision radio data at a variety of wavelengths and without a better understanding of the bubble production timescale. Nevertheless, the existence of this trend demonstrates quantitatively that the radio sources are indeed creating the cavities.  The radio sources are not simply filling preexisting voids in the intracluster medium (ICM) created by other processes.  Furthermore, the synchrotron luminosity and mechanical luminosity do not scale in direct proportion to each other.  This relationship implies that the synchrotron luminosity cannot be used to infer the mechanical power of a radio jet in a simple fashion.  An important and poorly understood aspect of radio source physics is the degree of coupling between the mechanical (kinetic) luminosity of radio sources and their synchrotron luminosity. This coupling is theoretically tied to the magnetic field strength and age of the source \\citep[see ][]{ddy93,bick97}, neither of which can be measured reliably from radio data alone.  Radio sources are inefficient radiators.  The ratio of mechanical power to radio power is typically assumed to range between 10 and 100, almost entirely on theoretical considerations \\citep[]{ddy93,bick97}.  On the other hand, measurements of the the X-ray cavity sizes and surrounding gas pressures provide unique estimates of their ages and mechanical luminosities, independently of the radio properties themselves. We evaluate the ratio of mechanical energy to radio power by plotting the ratio of mechanical power in the bubbles to monochromatic, 1.4~GHz synchrotron luminosity, assuming $1pV$ of energy per radio lobe, against the radio luminosity in the right-hand panel of Figure~\\ref{fig-Lr_Lm}. This ratio ranges from a few to a few hundred for the powerful sources, which is broadly consistent with theoretical estimates \\citep[see ][]{ddy93, bick97}.  On the other hand, Abell~478 has a ratio exceeding a few thousand. To the extent that X-ray cavities provide a good measure of the mechanical energy of radio sources, the large variation in this ratio indicates that radio luminosity is not necessarily a reliable probe of the available mechanical energy.  There are several factors that can introduce scatter into our estimate of the ratio of radio to kinetic power.  The most important is probably intrinsic differences between the radio sources themselves, a consequence of dramatic changes in radio luminosity with time. Certainly, if radio outbursts are to compensate for radiative losses in cooling flows, then the absence of radio emission from some systems requires large variations of radio luminosity with time.  On the other hand,  the $pV$ energy of the bubbles alone would tend to underestimate the the mechanical luminosity of radio sources by factors of several if energy dissipating shocks are generated, or if the bubbles expand non-adiabiatically (they leak),  or if the internal energy of the bubbles is boosted with a relativistic plasma.   \\subsection{Heating by Radio-Induced Cavities}\\label{sec-heating}  \\citet{chur02} noted the conversion of enthalpy of the rising bubble into other forms of energy in the cluster atmosphere.  Here it is shown that, for an adiabatic bubble, this energy is dissipated in its wake.  If the mass in the bubble is negligible compared to the mass of the gas it displaces, then a bubble rises, because the gas falls in around it to fill the space it occupied.  This process is driven by the potential energy released as the surrounding gas moves inward. The energy is first converted to gas kinetic energy, then dissipated in the wake of the rising bubble.  In the notation of \\S~\\ref{sec-cavage}, the potential energy released when the bubble rises a small distance, $\\delta R$, is \\begin{equation} \\delta W = \\rho V g \\, \\delta R = - V {dp\\over dR} \\, \\delta R, \\end{equation} where $\\rho$ is the gas density, and we have used the equation of hydrostatic equilibrium to replace $\\rho g = -dp/dR$, where $p$ is the gas pressure.  This gives a differential equation for the energy dissipated in the bubble wake \\begin{equation} {dW\\over dR} = - V {dp\\over dR}. \\end{equation} If the bubble is adiabatic, with ratio of specific heats $\\gamma$, then $p V^\\gamma =$ constant and this equation can be integrated to give the energy dissipated as the bubble rises over a large distance, from $R_0$ to $R_1$, \\begin{equation} \\Delta W ={\\gamma\\over \\gamma - 1} ( p_0 V_0 - p_1 V_1) = H_0 - H_1. \\end{equation} Here, subscripts 0 and 1 label quantities at the corresponding radii, and the enthalpy of the bubble is $H = \\gamma p V /(\\gamma - 1)$. Note that the bubble is assumed to be small compared to $R$ (otherwise there can be a significant change in the density of the gas as it falls in around the bubble).  This would rarely be significant in a cluster, but when it is, then some of the potential energy goes into readjustment of the atmosphere as the bubble moves.  For a relativistic gas, $\\gamma = 4/3$, so that the enthalpy is $4pV$. The region where this is dissipated by an adiabatic bubble is determined by the pressure distribution of the atmosphere.  For clusters such as Hydra~A and Perseus, roughly half of this energy would be dissipated inside the cooling radius.  It is likely that the bubbles are not entirely adiabatic.  On the basis of our numbers, radio losses are generally negligible, but pieces may be broken away from bubbles, and the relativistic particles may leak.  Such effects will generally lead to a greater proportion of the bubble energy being deposited within the cooling radius.  It is important to note that our estimate of the mechanical luminosity relies critically on the assumption that the bubbles are close to local pressure equilibrium.  This is at least approximately true for the Hydra~A cluster \\citep{nuls02}.  However, according to the standard view of radio sources, bubbles may have been significantly overpressured while being formed \\citep[e.g.\\ ][]{hein98}.  In that case, the expanding bubble drives a shock, and the energy deposited by the expansion can be substantially larger than $pV$.  There may be additional heat input from the AGNs associated with radio outbursts.  This could take the form of spherical shocks (driven by poorly collimated outflows), direct injection of relativistic particles, inverse Compton heating \\citep{ciot01}, or other processes. Very substantial additional heat inputs would drive convection, leading to an isentropic core and mixing out of abundance gradients \\citep{brug02a}, but this is not a very strong constraint.  If such energy injection is significantly more than the bubble energy input, then it is inappropriate to associate it directly with the bubbles, but the mean heating power may be correlated with bubble mechanical power.  Finally, it should be noted that even for adiabatic bubbles, the free energy of a bubble decreases with time, and bubbles may even break up quickly, so that they disappear as X-ray cavities.  This means that the instantaneous estimate of bubble mechanical power that we have used varies with time, and may vary dramatically.  A much better controlled sample is needed to investigate such issues.  \\phn  \\subsection{Can Cavity Production Quench Cooling Flows?}\\label{sec-quenching}  We now turn to the question of whether radio sources deposit enough energy into the ICM to quench cooling. We use $L_{\\rm{X}}$, the total luminosity of the X-ray--emitting gas from within the cooling radius, as an estimate of the classical, or morphological, cooling luminosity in the absence of heating and $L_{\\rm{spec}}$, the spectral estimate of the cooling luminosity within the cooling radius, as the luminosity of the gas cooling to low temperatures.  The cooling luminosity, $L_{\\rm{X}} - L_{\\rm{spec}}$, must be offset by heating in order to prevent the gas from cooling to low temperatures. We note that this quantity ignores non--X-ray cooling, such as ultraviolet and optical emission, predicted to result from cooling by thermal conduction inside magnetic flux loops \\citep{sok04}, or along reconnected magnetic field lines between cold clouds and the ICM \\citep{sok03}. Any such emission would lower the cooling luminosity which must be balanced by heating. Figure~\\ref{fig-Lx_Lm} shows the mechanical luminosity plotted against $L_{\\rm{X}} - L_{\\rm{spec}}$ for our sample. The diagonal lines represent equality between cooling and heating, assuming energy inputs of $pV,$ $4 pV,$ and $16 pV$ per cavity. The data are derived from Tables~\\ref{tbl-cluster-xray} and \\ref{tbl-bubbles}. For RBS~797, an upper limit is shown.  The cooling luminosity for RBS~797 is poorly constrained by the spectrum, which consists of only $\\sim 9000$ counts after cleanin.  RBS~797, while very luminous, is the most distant cluster in our sample and has the shortest exposure time (see Section~\\ref{sec-sample}).  \\begin{figure} \\plotone{plot3.eps} \\caption{Mechanical luminosity vs.\\ total luminosity minus the spectroscopic estimate of the cooling luminosity.  Lines denoting $L_{\\rm{X}} - L_{\\rm{spec}} = L_{\\rm{mech}}$ are shown for the assumption of $pV,$ $4pV,$ and $16pV$ energy in the bubbles.  Symbols and error bars as in Figure~\\ref{fig-Lr_Lm}; the arrow denotes an upper limit.   \\label{fig-Lx_Lm}} \\end{figure}   Figure~\\ref{fig-Lx_Lm} shows several objects, such as Hydra~A, Cygnus~A, and M84, whose cavities can contain enough energy to balance radiative losses, at least temporarily, with nearly $1pV$ of heat input per cavity. The remaining objects, which require between a few and $\\sim 20 pV$ per cavity to balance cooling, would do so with varying degrees of difficulty.  As discussed above, $\\sim 2 pV$ would be deposited within the cooling radius by an adiabatic bubble containing relativistic plasma.  Up to $4 pV$ is available if the cavities are relativistic and nonadiabatic, and there may be further energy input if they are overpressured or produce a shock when they are formed.  Therefore, the objects that require $\\sim 4 pV$ or less may reasonably be supplied with enough energy in the cavities to balance cooling, depending on the detailed dynamics (the heat also needs to be distributed inside the cooling radius to match the distribution of radiative losses). Provided that the true radio cycling timescale ranges between $t_{\\rm{c_{s}}}$ and $t_{\\rm r}$, the cavities in one quarter to one-half of the objects in our sample contain enough energy to offset radiation losses. This would be true of the cavities in the remaining objects only if they are significantly nonadiabatic, as outlined in \\S~\\ref{sec-quenching}.  Bear in mind that our conclusions depend on the adopted cooling radius (see \\S 3.2), measurement uncertainties in the cavity sizes, and the cavity production timescale. Nevertheless, we can safely conclude that cooling can plausibly be balanced by bubble heating in some, but not all, systems.  It is unnecessary to balance the entire luminosity, $L_{\\rm{X}} - L_{\\rm{spec}},$ by bubble heating alone if there are other forms of heating present. A possible source of heating is thermal conduction, which, as demonstrated by \\citet{voig04}, could supply a significant amount of heat.  Using a sample similar to our own, Voigt \\& Fabian found that thermal conduction can reduce the cooling luminosity by factors of $\\sim 2-3$ in some objects. Although they are difficult to find in X-ray images, shocks associated with the expanding cavities can deposit additional energy into the ICM.  Deep {\\it Chandra} images of a growing number of objects, including Cygnus~A \\citep{wils03}, NGC 4636 \\citep{jones02}, M87 \\citep{form04}, and Perseus \\citep{fabi03a}, show surface brightness discontinuities that may be associated with weak shocks. In Cygnus~A and M87, the shocks imply that the radio source may provide several times the upper limit of the luminosity seen in the bubbles, under the assumption of $4pV$ of energy per bubble \\citep{wils03,form04}. It may therefore be a combination of heating mechanisms that leads to quenched cooling, as suggested by several authors \\citep{brig02,brig03,kim03,rusz02}.  It is important to note that that our sample is biased toward systems with visible evidence for X-ray cavities, and does not represent clusters as a whole.  Many clusters, including some with large cooling flows, do not contain cavities \\citep[e.g.\\ Abell 1068,][]{wise03,mcna04}. These objects may have very different reheating histories than the objects discussed here.  In this sense, the objects presented here represent the best-case examples for reheating the ICM by energetic bubbles.  Our analysis does not imply that all cooling flows can be quenched in this fashion.  \\subsection{Trends between X-ray and Mechanical Luminosities}\\label{sec-Lx-trends}  Figure~\\ref{fig-Lx_Lm} shows a trend between the X-ray luminosity and bubble mechanical luminosity, with the sense that systems with larger X-ray luminosities also have larger mechanical luminosities.  This trend extends over a dynamic range of $\\sim 1000$ in both X-ray and mechanical luminosity.  Just such a trend would be expected were the cooling and heating of the ICM coupled in some fashion. Several studies \\citep[e.g.][]{rosn89,binn95,davi01,quil01,chur02} have proposed that cooling is balanced by heating in a self-regulated feedback loop.  The feedback loop is driven by episodic radio activity fueled by cooling and accretion onto a central black hole.  The accretion energy is then returned to the ICM through an AGN outburst, including the action of the radio cavities, which temporarily arrests cooling.  At later times, the center of the system settles down and the cooling flow is reestablished.  During the cooling cycle,  molecular gas \\citep{edge01} accumulates and star formation ensues \\citep{jfn,mo89}, albeit at substantially lower levels than expected in steady-cooling models \\citep{fabi94}. Even if the radio bubbles are not the main source of heat, Figure~\\ref{fig-Lx_Lm} suggests that AGN feedback is intimately involved in the process that prevents a large cooling flow from forming.  The apparent correlation in Figure~\\ref{fig-Lx_Lm} should be treated with caution. As noted earlier, our sample was selected from clusters in the \\textit{Chandra} archive with fairly obvious cavities in their cores, and neglects those without obvious cavities. Other clusters are known to have substantial cooling luminosities commensurate with the observed levels of cold gas and star formation, yet contain no cavities and have low radio power.  A prime example is A1068 \\citep{wise03,mcna04}. Similar objects would appear in the lower right of this diagram, tending to weaken the correlation.  On the other hand, it is unlikely that we would have missed objects with powerful cavities, which would lie in the upper part of this diagram.  Therefore, the distribution of points may represent an upper envelope in mechanical luminosity as a function of  X-ray luminosity.  Such a distribution would be consistent with the feedback hypothesis, if objects like A1068 are in an extended cooling phase in which the central galaxies have experienced substantial levels of accretion in the past 100~Myr or so, when the radio source has not had a chance to create cavities capable of reducing or quenching cooling.  We have investigated the degree to which other systematic effects may lead to an unphysical luminosity-luminosity correlation. For example, \\citet{elvi78} pointed out that a sample of objects with a small range of fluxes and a large range of distances will show a correlation in a luminosity-versus-luminosity plot, even if there is no intrinsic correlation in the sample.  Our sample, however, has a large range of radio fluxes (from $\\sim 10$~mJy to $\\sim 10^{6}$~mJy). The cavities in our sample also cover a large range of projected angular sizes (from $\\sim 3$\" to $\\sim 35$\").  We believe, then, that these potential effects are unlikely to account entirely for the trends seen in Figures~\\ref{fig-Lr_Lm} and \\ref{fig-Lx_Lm}.  Selection bias may also contribute to the correlations. Small cavities are easily overlooked in distant objects, since cavities of a given linear size become more difficult to detect as their angular sizes decrease with increasing distance. Conversely, in very nearby objects, such as M87, we may miss larger bubbles that lie outside the detector.  Furthermore, other considerations, such as the bubble position, affect the detectability \\citep{enss02}.  The consequences of these and other effects on our selection function will be addressed in the future using a larger and better-defined sample of clusters, including a more sophisticated approach to placing limits on cavities that may exist in clusters but were missed by the observations.   \\subsection{Summary}  We have presented an analysis of 18 systems taken from the \\textit{Chandra} archive having clear evidence for cavities in their X-ray emission.  We find that the energy associated with the cavities is sufficient to substantially reduce or quench cooling in nearly half of the objects in our sample.  However, this mechanism alone probably does not provide a general solution to the cooling problem, unless X-ray cavities probe only a small fraction of the total kinetic luminosity of radio sources. In addition, we have discovered a trend between the cooling X-ray luminosity and the mechanical energy of the cavities, with the sense that more luminous systems produce larger and more energetic cavities. The  trend, or envelope, may have been established by a self-regulated cooling and feedback mechanism acting in many systems. The existence of such a mechanism in relatively nearby clusters, where the detailed physics can be examined, may provide significant insight on the process of galaxy formation that prevails at large redshifts \\citep[e.g.\\ ][]{voit03}. A similar mechanism may regulate the growth of galaxy halos during the dissipative stages of their development \\citep{dubi94} and may be an agent responsible for the detailed correlation between black hole mass and velocity dispersion of spheroids \\citep{fabi02b}. We have measured for the first time the distribution of the ratio of kinetic luminosity to monochromatic radio luminosity for a sample of radio sources. The ratio varies widely, with most objects ranging between few and a few hundred, assuming $1pV$ of energy per cavity.  X-ray cavities provide a unique probe of the mechanical power of radio jets, independently of the radio properties themselves.  Our future plans include expanding the sample size and acquiring better and more uniform radio data. In addition, we plan to extend our understanding of the detectability function of bubbles, using simulations of images with a wide range of exposure and signal-to-noise ratios.  \\phn  We acknowledge helpful discussions with Mangala Sharma, Liz Blanton, Frazer Owen, Hui Li, Phil Kronberg, and David De Young, and we thank Carlo Nipoti for pointing out an error in the Centaurus radio flux in an earlier version. This work was supported by Long Term Space Astrophysics Grant NAG5-11025, \\textit{Chandra} Archive Grant AR2-3007X, and a grant from the Department of Energy through the Los Alamos National Laboratory."
    },
    "0402/astro-ph0402454_arXiv.txt": {
        "abstract": "Using the SCUBA instrument on the JCMT, we have made a submillimeter mosaic at 850\\um\\ of a subarea of the {\\it Spitzer} First Look Survey (FLS).  Our image covers the central 151 \\sqam\\ of the northern extragalactic Continuous Viewing Zone (CVZ) field of the FLS to a median 3\\s\\ depth of 9.7 mJy.  The image contains ten 850\\um\\ sources detected at 3.5\\s\\ or higher significance, of which five are detected at $\\geq$4\\s.  We make the catalog of these SCUBA-selected FLS sources available to the community.  After correcting for incompleteness and flux bias, we find that the density of sources brighter that 10mJy in our field is (1.3$^{+1.1}_{-0.7}$)$\\times$$10^2$ \\persqdeg\\ (95\\% Poisson confidence limits), which is consistent with other surveys that probe the bright end of the submillimeter population. ",
        "introduction": "Launched in the second half of 2003, the Spitzer Space Telescope (formerly SIRTF), the fourth and final of NASA's Great Observatories, holds the promise of addressing many outstanding questions related to dust-enshrouded galaxy formation at high redshift.  One of the first science observations to be undertaken by {\\it Spitzer} is the {\\it Spitzer} First Look Survey (FLS), a first look at the mid-IR sky at sensitivities that are two orders of magnitude deeper than previous large-area surveys.  In addition to {\\it Spitzer} data at 3.6, 4.5, 5.8, 8.0, 24, 70, and 160\\um, the FLS has also been observed in deep ground-based campaigns at optical (KPNO 4m to $R$=25.5, 5\\s\\ in a 2\\arcsec\\ aperture) and radio (VLA to 115$\\mu$Jy, 5\\s, per 5\\arcsec\\ beam at 1.4GHz; Condon et al.\\ 2003) wavelengths\\footnote{See http://ssc.spitzer.caltech.edu/fls/}.  The {\\it Spitzer} FLS data will be released to the public in early 2004 and, together with the deep ancillary ground-based data, will provide the community with the first systematic look at the properties of faint {\\it Spitzer}-selected extragalactic sources.  While {\\it Spitzer} will discover many dust-enshrouded high-$z$ objects and will greatly help us understand their nature, the fact that it does not image at wavelengths longward of 160\\um\\ presents some important limitations.  For example, for typical dust temperatures (20--40K; e.g., Dunne et al.\\ 2000), the longest {\\it Spitzer} passband barely probes longward of the peak of the thermal dust emission for galaxies at even moderate redshifts, making it difficult to estimate their bolometric luminosities and hence infer quantities such as dust masses and star formation rates.  Moreover, with increasing redshift (or decreasing dust temperature), an object's dust emission peak shifts redward of the longest {\\it Spitzer} wavelength, causing strong negative k-corrections and making a galaxy at high redshift (or one with low dust temperatures) fade rapidly out of a Spitzer-selected sample.  In this Letter we present complementary long-wavelength imaging observations of a section of the {\\it Spitzer} FLS, obtained at 850\\um\\ with the Submillimetre \\notetoeditor{Editor: this is an instrument whose follows the British spelling convention of submillimetre/submillimeter} Common User Bolometer Array (SCUBA) on the James Clerk Maxwell Telescope (JCMT).  In a future paper we will discuss in detail the multiwavelength properties of sources in the area of our SCUBA map; the purpose of the present Letter is to quickly make available to the community the source catalog of objects from our SCUBA observations within the public-release {\\it Spitzer} FLS. ",
        "conclusions": "\\begin{figure} \\plotone{f3.lowres.ps} \\caption{\\label{radio-posn.fig} Positions of VLA radio detections compared to our 850\\um\\ SCUBA sources. Radio detections are shown as crosses, while SCUBA sources are represented with white circled points, with the size of the {\\it inner point} corresponding roughly to the 15\\arcsec\\ diameter of the SCUBA beam.  None of our SCUBA sources have radio detections. A HIGH QUALITY VERISION OF THIS FIGURE CAN BE OBTAINED FROM THE AUTHORS} \\end{figure}  In this Letter we presented our SCUBA observations of a 151 \\sqam\\ subarea in the northern CVZ field of the {\\it Spitzer} First Look Survey.  We found a total of 10 sources at S/N$>$3.5 and make their particulars available to the community. Our integrated source counts are consistent with those of the other two surveys of the bright end of the sub-mm population, namely the 8mJy Survey (Scott et al.\\ 2002) and the Hubble Deep Field supermap (Borys et al.\\ 2002, 2003).  Given that extragalactic sub-mm sources cluster strongly on the scales of current surveys (see, e.g., Figure~\\ref{maps.fig}), the fact that our number counts agree with those of the other two surveys gives an important confirmation of the numbers of SCUBA sources at the bright end of the population. Equally significantly, the agreement between our number counts and those of other surveys suggests that the subfield of the FLS that we imaged is not unrepresentative of the extragalactic sky.  We will study the multi-wavelength properties of the sub-mm sources in our map once the {\\it Spitzer} FLS data are released, In the meantime, we have compared the positions of our SCUBA detections with those of objects in the 1.4GHz VLA map of the FLS (Condon et al.\\ 2003).  There is no correspondence between the radio-selected and the sub-mm selected populations (see Figure 3) down to the limit of the VLA catalog (115 mJy, 5$\\sigma$).  This lack of radio detection of any of our sub-mm sources can be used to constrain their redshifts (Dunne, Clements \\& Eales 2000; see also Yun \\& Carilli, 2002): given the VLA flux limit and the rather narrow dynamic range of these data, most of our sub-mm sources appear to be at $z$$\\gtrsim$ 1.6--1.7 and the brightest two at $z$ $\\gtrsim$ 2.0, (however, these results are likely to be affected by flux boosting).  These redshift constrains are in line with our current knowledge of the redshift distribution of sub-mm selected sources: the median redshift of the population is believed to lie at $z \\sim$ 2--3, and evidence exists of a flux-redshift relation, such that more luminous sub-mm selected systems (such as those in our survey) reside at higher redshifts than the less luminous objects (Ivison et al.\\ 2002, Smail et al.\\ 2002, Webb et al.\\ 2003, Chapman et al.\\ 2003, Clements et al.\\ 2004). The lack of 1.4 GHz detection of our sub-mm sources is thus not unexpected: deeper radio data will be needed to detect and identify these systems.  We will explore these issues further once the {\\it Spitzer} FLS data are released.      We thank the Joint Astronomy Centre\\notetoeditor{Editor: note that this is an institution whose name that follows the British spelling convention of centre/center} staff who helped us obtain these data, and our colleagues who carried out some of these observations through the Canadian flexible scheduling scheme.  We also thank Colin Borys for providing in digital format the number counts shown in Fig.~\\ref{counts.fig}.  The JCMT is operated by the Joint Astronomy Centre\\notetoeditor{Editor: British spelling here, as above} on behalf of the Particle Physics and Astronoy Reseach Council of the United Kingdom, the Netherlands Organisation for Scientific Research, and the National Research Council of Canada."
    },
    "0402/astro-ph0402440_arXiv.txt": {
        "abstract": "In a parameter study extending to jet densities of $10^{-5}$ times the ambient one, I have recently shown that light large scale jets start their lives in a spherical bow shock phase. This allows an easy description of the sideways bow shock propagation in that phase. Here, I present new, bipolar, simulations of very light jets in 2.5D and 3D, reaching the observationally relevant scale of $>200$ jet radii. Deviations from the early bow shock propagation law are expected because of various effects.  The net effect is, however, shown to remain small. I calculate the X-ray appearance of the shocked cluster gas and compare it to Cygnus~A and 3C~317. Rings, bright spots and enhancements inside the radio cocoon may be explained. ",
        "introduction": "X-ray studies of galaxy cluster centers containing a radio jet have shown that the jets have a considerable impact on the cluster gas \\cite{CPH94,Cea02,Blanea01,Sea01}, forming rings, apparent spirals, and aligned features. Such systems have been claimed to be associated with very light jets \\cite{CHC97,Rosea99,mypap02d}. Parameter studies of very light jets have been carried out recently \\cite{CO02a,CO02b,Saxea02,mypap03a,Zanea03}. There, it has been shown that very light jets first form spherically symmetric bow shocks \\cite{mypap03a}. In that phase, they follow the expansion law (derived for a strong bow shock, which is applicable here): \\begin{equation}\\label{globeqmot} \\int_0^r{\\cal M}(r^\\prime)r^\\prime\\; \\mathrm{d}r^\\prime= 2\\int_0^t \\mathrm{d}t^\\prime \\int_0^{t^\\prime} E(t^{\\prime\\prime}) \\mathrm{d}t^{\\prime\\prime}\\enspace, \\end{equation} where ${\\cal M}(r)$ is an arbitrary spherically symmetric ambient gas mass distribution and $E(t)$ is the energy injection law. Here, I show 3D and 2.5D large scale simulations, updating the X-ray appearance of the shocked IGM, and exploring the accuracy of the spherical expansion law for the bow shock at late times. \\begin{figure*}[t] \\begin{center} \\rotatebox{0}{\\includegraphics[width=\\textwidth]{krause_fig1.ps}} \\caption{\\small Logarithm of number density (slice) for the 3D run at $2.04$~Myr. \\label{ml11}} \\end{center} \\end{figure*} ",
        "conclusions": "The bipolar simulations in connection with detailed studies of the early lives of very light jets reveal two parts of the bow shock: an inner elliptical part, and an outer cigar shaped one. These parts also appear in the X-ray emission. When viewed from an appropriate angle, they can appear as partial rings and bright spots. I suggest that the two X-ray rings in 3C~317 \\cite{Blanea01} are caused by this effect.  The X-ray appearance of the 2.5D simulation reproduces some important details of Cygnus~A's X-ray emission \\cite{Sea01}: the elliptical deformation in the shocked ambient gas region, the bright X-ray filaments inside the cocoon, and the fork like structure around the cocoon. The bow shock can be located at the interface, where the elliptical X-ray contours meet the spherical ones from the unaffected cluster gas. This good agreement also supports the idea of a very light, relativistic, and magnetised jet in Cygnus~A."
    },
    "0402/astro-ph0402395_arXiv.txt": {
        "abstract": "{ The complete near infrared (0.9-2.5 $\\mu$m) spectra in three different star forming regions (HH24-26, HH72 and BHR71) are presented and analyzed in the framework of shock excitation models. The spectra are dominated by H$_2$ rovibrational emission (vibrational state $\\nu$ $\\le$ 5, excitation energy $T_{ex}$ $\\le$ 35\\,000 K), while emission from ionized material, recognizable from [Fe II] and [S II] lines, is significantly fainter. The analysis of the H$_2$ excitation diagrams points to the existence of two different excitation regimes: whilst condensations observed only in the infrared appear to have temperatures rarely exceeding 3000 K and can be modelled in the framework of steady-state C-shock models, the infrared counterparts of Herbig Haro (HH) objects exhibit a temperature stratification with components up to more than 5000 K. The H$_2$ emission from representative HH objects (HH26A, HH72A and HH320A) has been successfully modelled by planar J-shocks with magnetic precursors, for which the main parameters (pre-shock density, speed) are derived. However, these same models are unable to reproduce the observed atomic and ionic emission, which probably arises from a distinct and perhaps more embedded region with respect to that traced by the H$_2$. Some of the physical parameters of such regions (fractional ionization, density) have been estimated in HH72, on the basis of the observed ionic lines. ",
        "introduction": "The physical effect exerted on the ambient medium by the impact of jets from accreting protostars results in the formation of shock waves. The compressed and heated gas radiates away the accumulated thermal energy through the emission of ionic, atomic and molecular lines, in relative proportions which depend on the structure of the shock wave. During the gas cooling, a fundamental role is played by molecular hydrogen, whose high abundance compensates the low transition rates due to its homonuclear nature. In the near infrared, the association of H$_2$ emission with knots of shocked gas has been demonstrated by mapping at sub-arcsecond scale in the $\\nu$=1-0 S(1) line at 2.122 $\\mu$m (e.g. Eisl\\\"{o}ffel 2000), and the rovibrational lines are effective probes of the molecular gas at thousands of kelvin, which is missed by optical observations. H$_2$ line intensities and intensity ratios are also extensively used to distinguish between fluorescence and shock excitation as possible mechanisms at the origin of the emission, since strong H$_2$ lines with $\\nu$$\\ge$6 in the 1.0-1.4 $\\mu$m range are expected if fluorescence is responsible for the emission (Black $\\&$ van Dishoeck 1987). It has proved more difficult to derive the nature (C-ontinuous or J-ump type, e.g. Kaufman $\\&$ Neufeld 1996, Hollenbach $\\&$ McKee 1989, Draine 1980) and the physical parameters of the shock waves, such as the velocity, the preshock density, the strength of the transverse magnetic field and the temperature of the neutral gas. In the conventional scenario, C-type shocks can exist up to velocities typically $\\approx$~50 km s$^{-1}$, beyond which the H$_2$ molecule is collisionally dissociated and the kinetic temperature of the gas rapidly increases, giving rise to a discontinuity in the shock parameters. However, more recent models (e.g. Smith 1995, Le Bourlot et al. 2002, Flower et al. 2003) have shown that H$_2$ dissociation can be inhibited over a wider range of shock parameters: from slow ($v_{s}$ $\\la$ 25 km s$^{-1}$), partially dissociative J-type shocks, up to fast C-type shocks travelling at 70-80 km s$^{-1}$.  From an observational point of view, the presence of different shock components (J- or C-type) can be demonstrated by the modifications induced in the H$_2$ excitation diagram: the lower post-shock densities attained in C-shocks result in a larger departure from LTE conditions than in J-shocks of the same speed (Flower et al. 2003). In order to probe such deviations, it is essential to investigate spectroscopically the 1.0-1.4 $\\mu$m range, since several H$_2$ lines with different vibrational quantum number and of relatively high excitation energy ($>$ 15\\,000 K) lie at these wavelengths.  With the aim of observing such lines, we have undertaken a spectroscopic survey (from 0.9 to 2.5 $\\mu$m) of a sample of Herbig-Haro (HH) objects and H$_2$ jets, based on the observations gathered with the SOFI spectrometer at the ESO-NTT. We have already reported the results of this survey regarding HH43, HH111, HH240/241 and HH120 (Giannini et al. 2002, Nisini et al. 2002, hereafter Paper~I and II, respectively), showing that quite different excitation conditions can occur in HH objects: while in HH43 the bulk of the cooling occurs in H$_2$ lines, the spectra of the other three objects are dominated by iron lines. In HH43, the H$_2$ emission has been successfully fitted by means of high velocity, C-type shock models (Flower et al. 2003), whereas, in the other three cases, the strong ionic emission testifies in favour of a dissociative component in the shock structure. Here we present the observations obtained in three other regions, namely HH24-26, HH72 and HH320/321 (BHR71).  All the spectra exhibit copious H$_2$ rovibrational emission, together, in some cases, with a fainter atomic and ionic component. Therefore, they represent valid tests of the ability of current shock models to predict simultaneously molecular and atomic/ionic emission and an opportunity to check the validity of the underlying assumptions of the models.  The structure of the paper is the following: we present the targetted regions in Sect. 2 and then describe the observations and the results obtained (Sect. 3). In Sect. 4, we derive the physical parameters of the emitting gas and model the observed emission, assuming shock excitation. Sect. 5 summarizes our conclusions. ",
        "conclusions": " \\begin{itemize} \\item[-] Strong H$_2$ emission lines with excitation energies up to 35\\,000 K are detected throughout all the Herbig-Haro objects present in the three regions, while infrared knots show lines with excitation energies rarely exceeding 15\\,000 K. This difference in the line emission reflects different excitation regimes: while the condensations observed only in the infrared are excited at a single temperature of $\\approx$\\,2000-3000 K, Herbig-Haro objects have a temperature stratification, with components up to more than 5000 K, which tend towards the values (of the order of 10$^4$ K) determined from optical transitions. We note that, in order to trace effectively the highest temperature components, it is essential to obtain observations in the 1.0-1.4 $\\mu$m range, where the lines with high vibrational quantum numbers are located.  \\item[-] Atomic and ionic emission in form of [{\\fe}], [{\\s}] and {[{\\n}]} transitions is detected in only a few HH objects. The ratios of lines of different species have been used to derive some of the physical parameters of the ionized gas (electron density, iron abundance and ionization fraction), which all indicate that the shocks present in these objects have lower excitation/ionization conditions than in cases where stronger [{\\fe}] emission occurs.  \\item[-] The observed H$_2$ emission in the HH objects can be reproduced by models of J-type shocks with magnetic precursors and ages of typically a few hundred years. The shock speeds are in the range from about 30 to about 50 km s$^{-1}$, and the preshock gas density is of the order of 10$^4$ cm$^{-3}$. In the case of the pure H$_2$ knot HH25C, it is possible to fit marginally the observational data by means of a steady state (C-type) shock model with a higher preshock density (10$^5$ cm$^{-3}$).  \\item[-] The same planar model that fits the H$_{2}$ emission systematically underestimates both the [{\\fe}] and [{\\ci}] lines fluxes observed in HH72A and HH26A. We believe that either a bow shock (with atomic and ionic emission originating at the apex of the shock and with H$_{2}$ arising in the wings) or a reverse shock in denser gas at the centre of the outflow is responsible for the ionic and atomic emission. Supporting evidence comes from the high electron density ($\\approx$ 5 $\\times$ 10$^4$ cm$^{-3}$) and visual extinction derived from the [{\\fe}] lines.  \\end{itemize}  \\emph{Acknowlegements}: We thank an anonymous referee for a detailed and constructive report."
    },
    "0402/astro-ph0402676_arXiv.txt": {
        "abstract": " ",
        "introduction": "Narrow-line Seyfert 1 galaxies are identified by their optical emission line properties.  They have narrow permitted optical lines (FWHM of H$\\beta<2000\\rm\\, km\\,s^{-1}$), weak forbidden lines ([\\ion{O}{III}]/H$\\beta<3$; this distinguishes them from Seyfert 2 galaxies), and frequently they show strong \\ion{Fe}{II} emission.\\cite{rf:1} In 1992, it was demonstrated by Boroson \\& Green\\cite{rf:2} that the optical emission-line properties around H$\\beta$ are strongly correlated with one another. A principal components analysis allowed the largest differences among optical emission line properties to be gathered together in a construct commonly known as ``Eigenvector 1''. The strongest differences hinge on the strength of the \\ion{Fe}{II} and [\\ion{O}{III}] emission, and the width and asymmetry of H$\\beta$. These are just the properties that define NLS1s.  During the 1990's, it was found that the X-ray properties are also manifested in these correlations: NLS1s are observed to have steeper soft X-ray spectra,\\cite{rf:3} steeper hard X-ray spectra,\\cite{rf:4,rf:5} and higher amplitude X-ray variability.\\cite{rf:6} Properties of UV spectra also appear in Eigenvector 1: NLS1s tend to have higher \\ion{Si}{III}]/\\ion{C}{III}] ratios, stronger low-ionization lines, weaker \\ion{C}{IV}, and stronger \\ion{N}{V}.\\cite{rf:7}  These sets of strong correlations are remarkable, because they involve dynamics and gas properties in emission regions separated by vast distances. This pervasiveness leads us to believe that we are observing the manifestation of a primary physical parameter. A favored explanation is that it is the accretion rate relative to the black hole mass onto the active nucleus.  This is easily understood from the X-ray variability properties as follows. NLS1s have systematically higher fractional amplitude of variability at a particular X-ray luminosity than do Seyfert 1 galaxies with broad optical lines in {\\it ASCA} observations (Fig.\\ 1).\\cite{rf:6} Since the {\\it ASCA} observations are all nearly the same length, a higher fractional amplitude of variability implies a shorter variability time scale.  A shorter time scale implies a smaller emission region, which corresponds to a smaller black hole mass, assuming that the geometry, etc., are uniform among Seyfert 1s.  Assuming a constant efficiency of conversion of gravitational potential energy to radiation, the X-ray luminosity corresponds to the absolute accretion rate.  Thus, for a given luminosity or accretion rate, NLS1s have a smaller black hole mass, and therefore have a higher accretion rate relative to the Eddington value.  The high-accretion-rate scenario is attractive and simple.  However, it appears to be incomplete.  In the process of analyzing the {\\it ASCA} X-ray spectra from NLS1s, I found that their properties did not appear to be uniform.  Specifically, there is a correlation between their fractional amplitude of variability and a parameter, $\\alpha_{xx}$, that measures the strength of the X-ray soft excess (Fig.\\ 1).\\cite{rf:5} Objects with very prominent soft excesses show high-amplitude variability, and objects with moderate soft excesses show lower-amplitude variability.  Thus, a connection between the shape of the X-ray spectrum and the variability among NLS1s is observed.  \\begin{figure} \\centerline{\\includegraphics[width=12.2 cm]{figure1_proc.eps}} \\caption{{\\it Left:} Excess variance, defined as the measured variance corrected for the measurement error and normalized by the mean, as a function of the 2--10 keV luminosity of a sample of Seyfert 1 galaxies observed with {\\it ASCA}.\\cite{rf:6}  {\\it Right:} The excess variance as a function of $\\alpha_{xx}$, a parameter defined as the point-to-point energy index between 0.7 and 4 keV of the best-fitting continuum spectral model, for a sample of NLS1s.\\cite{rf:5}} \\label{fig:1} \\end{figure}  The {\\it ASCA} data were invaluable for understanding the properties of NLS1s; however, because of the limited signal-to-noise ratio, we could not investigate the nature of the soft excess or the variability in detail. The large effective area of {\\it  XMM-Newton} vastly improves the situation.  We can gain some insight into the correlation shown in Fig.\\ 1 by looking at archival {\\it XMM-Newton} data from two NLS1s, 1H~0707$-$495 and Ton~S180, with different locations in the correlation plot. ",
        "conclusions": "In this talk, I presented a discussion of the X-ray spectral and variability properties exhibited by NLS1s, drawing from studies of {\\it ASCA} spectra,\\cite{rf:5,rf:6} and UV properties, drawing from studies of {\\it HST} spectra.\\cite{rf:11,rf:14}  I showed that although NLS1s are identified by fairly uniform optical spectral properties, they exhibit a {\\it range} of X-ray and UV spectral behaviors.  I also presented an explicit comparison of X-ray and UV properties of two NLS1s, 1H~0707$-$495 and Ton~S180.  These two objects show distinctly different X-ray spectral and variability behaviors.  This suggests that something about their central engines is different; perhaps there are differences in geometry resulting from different accretion rate.  There are also differences in their UV properties; here I discussed the differences in the \\ion{C}{IV} profile, but note that there are other patterns in the range of behaviors exhibited by NLS1s that are discussed in detail in Ref.~\\citen{rf:11}.  The spectral analysis and photoionization modeling presented in Ref.~\\citen{rf:14} indicates that the shape of the spectral energy distribution may be responsible for the range of UV spectral behaviors.  Specifically, when the UV continuum is blue and strong, and the X-ray continuum is weak, a resonance-line driven wind is produced that contributes blueshifted high-ionization emission lines.  This wind filters the continuum before it illuminates the intermediate- and low-ionization line-emitting region, causing it to emit strong low-ionization lines, and lower I.P.\\ intermediate-ionization lines. In contrast, when the X-rays are strong relative to the UV, the gas that would form the wind is overionized, and the resonance-line driving fails. In that case, the high-ionization lines are narrow and centered at the rest wavelength.  The spectral energy distribution is produced in the central engine. X-ray properties are key for understanding the central engine, because the X-rays are emitted very close to the black hole.  I propose that the differences seen in the X-ray properties of 1H~0707$-$495 and Ton~S180 are manifested in their spectral energy distributions, and perhaps also the geometry of the central engine, which then influences their UV emission line properties. This suggests a chain of causality linking the X-ray and UV behaviors perhaps arising from a range in some intrinsic parameter among members of the class of NLS1s."
    },
    "0402/astro-ph0402506_arXiv.txt": {
        "abstract": "We investigate the effect of a global change in the ionizing continuum level on the behavior of the strong optical broad emission lines seen in spectra of the nuclear emission-line regions of active galactic nuclei (AGN), including the Balmer lines, \\ion{He}{1} $\\lambda$5876, and \\ion{He}{2} $\\lambda$4686. Unlike most of the prominent heavy element lines found in the UV, the optical hydrogen and helium recombination lines' emissivities are strongly dependent on the incident continuum flux, since these lines arise out of excited states whose {\\em optical depths depend on the incident flux of photons}. Using photoionization calculations we determine the luminosity-dependent responsivities, $\\eta(r,L(t)) = \\Delta \\log L_{line}/\\Delta \\log L_{cont}$, of these lines for a general model of the broad emission line region (BLR), with the purpose of establishing them as important probes of the physical conditions within the BLR of AGNs. The dependence of these lines' emissivities on the incident photon flux invokes a dependence in their responsivities on distance from the central continuum source. In particular, the responsivities of these lines are generally anticorrelated with the incident photon flux. Thus, their responsivities vary with distance within the BLR for a fixed continuum luminosity and change with time as the continuum source varies. Consequently, after correcting for light-travel-time effects the response of the Balmer and optical helium lines should generally be strongest during low continuum luminosity states. Responsivity that depends on photon flux and continuum state may explain a number of outstanding problems currently under investigation in broad-line variability studies of these and other emission lines. These include the origin of the intrinsic Baldwin effect, measurements of luminosity-dependent lags (a ``breathing'' BLR) and luminosity-dependent variations in: integrated broad emission-line flux ratios (including \\ion{He}{2} $\\lambda$4686/H$\\beta$), broad line profile shapes, and radial velocity-dependent intensity ratios. The broad H$\\alpha$/H$\\beta$ and H$\\beta$/\\ion{He}{1} flux ratios and the Balmer emission-line responsivity are observed to decrease from the line center to the line wings. These, along with our findings, lead to the conclusion that the BLR velocity field diminishes with increasing distance from the central continuum source. This is consistent with recent reverberation studies that find a relationship between the emission-line lag and rms profile width for multiple lines in individual AGN, which implies that the velocity field is dominated by a central massive object. Finally, the responsivity of ionization-bounded clouds can account for much of the observed behavior of the optical recombination lines (e.g., the weak response of the Balmer line wings) previously attributed to a substantial contribution from matter-bounded clouds at small BLR radii. ",
        "introduction": "Historically, attempts at interpreting the observed variability behavior of the broad emission lines in active galactic nuclei (AGN) have tended to regard the broad emission line region (BLR) as a stationary entity (non-evolving). By this we mean that any variability in the observed emission-line intensities has been attributed solely to reverberation (light-travel-time) effects within a spatially extended BLR. Thus, variations in the measured BLR sizes and the amplitude of the line response have been assumed to arise primarily through differences in the temporal variability of the ionizing continuum, which manifests as a difference in shape and FWHM of the continuum autocorrelation function (ACF) from one event to the next (e.g., P\\'{e}rez, Robinson, \\& de~la~Fuenta 1992). In essence, higher frequency continuum variations are better able to probe smaller BLR sizes, while larger BLR sizes are better probed by lower frequency continuum variations.  While a difference in the continuum variability time scale (e.g., the width of the continuum ACF) explains in part some of the observed variability seen in the broad emission lines in AGNs, it is by no means the complete picture. Indeed, more sophisticated photoionization calculations (e.g., see Krolik et~al.\\ 1991; O'Brien et~al.\\ 1995; Bottorff et~al.\\ 1997; Kaspi \\& Netzer 1999; Horne, Korista, \\& Goad 2003) have modeled the line variability by including the local emission-line response to changes in the overall luminosity and/or shape of the ionizing continuum.  The models presented by Carroll (1985) and Netzer (1991) anticipated the non-linear response of some of the emission lines.  In this paper we illustrate, using a general description of the BLR, how large changes in the mean continuum level may explain much of the observed variability behavior of the optical recombination lines. In $\\S$~2 we describe the grid of photoionization models used in this study. In $\\S$~3 we describe the physical origin of line responsivity and its impact on the following issues related to BLR variability: the origin and slope of the intrinsic Baldwin effect ($\\S$~3.2; e.g., Kinney, Rivolo, \\& Koratkar 1990); variations in the measured continuum--emission-line time-delays ($\\S$~3.3; e.g., Peterson et~al.\\ 2002), variations in the emission line flux ratios ($\\S$~3.4; e.g., Balmer decrement: Tran, Osterbrock, \\& Martel 1992; H$\\beta$/\\ion{He}{2}~4686: Peterson \\& Ferland 1986); changes in profile shape ($\\S$~3.5; e.g., Wanders \\& Peterson 1996); and velocity-dependent line intensity ratios ($\\S$~3.6; e.g., Stirpe, de~Bruyn, \\& van~Groningen 1988). We emphasize that the variations in these quantities discussed here are {\\em not} a consequence of reverberation effects within a finite-sized BLR, although they are intimately related. Rather, they are due to global changes in the mean ionization state of the BLR gas. In $\\S$~4 we discuss the impact of these findings on the purported contribution of a very broad optically thin component to the emission-line response, and other related issues. It is our hope that many of these observed emission-line variations and trends reported in the literature may be examined in a new light. We summarize our main findings in $\\S$~5. ",
        "conclusions": "\\subsection{Optically Thin Gas in the So-Called Very Broad Line Region}  That the BLR may contain a significant component of ionized gas that is optically thin to hydrogen-ionizing photons has been suggested by investigators over many years. In nearly every case there is an association of this gas with that emitting the very broad wings of the lines. Most recently, Sulentic et~al.\\ (2000) reported that in the quasar PG~1416$-$129 the FWHM of the broad emission-line profile of H$\\beta$ increased by 50\\% from a high to a low continuum state separated in time by 10 yr. Their interpretation of this is that the wings of the Balmer lines are dominated by emission from an optically thin ``very broad line region'' (VBLR), while the core of the line is dominated by more highly variable optically thick gas. These are usually thought of as two distinct line-emitting regions. Ferland, Korista, \\& Peterson (1990), Peterson et~al.\\ (1993), and Corbin \\& Smith (2000) made similar arguments based on Balmer line -- continuum variations, while Gondhalekar (1987, 1990), O'Brien, Zheng, \\& Wilson (1989), \\& P\\'{e}rez, Penston, \\& Moles (1989) found less variable wings in Ly$\\alpha$ and \\ion{C}{4} $\\lambda$1549 of higher redshift quasars. Kassebaum et~al.\\ (1997) found similarly stronger variations in the line core than wings of H$\\beta$ in Mrk~335. It should be noted that while the conclusions regarding the core versus wing variations from monitoring campaign studies such as Peterson et~al.\\ (1993) and Kassebaum et~al.\\ (1997) were based on analyses of {\\em time-resolved} spectral observations of the continuum and H$\\beta$, most of the above studies relied on very few, often just two, observation epochs. Without sufficient time resolution of the variability, it is difficult to assess the actual line response. Additional arguments for an optically thin line-emitting region have been put forth through the analyses of line flux ratios in the profile wings: optical Balmer and \\ion{He}{2} lines (Marziani \\& Sulentic 1993; Corbin 1997), Balmer and \\ion{C}{4} $\\lambda$1549 (Corbin 1995), Balmer and Ly$\\alpha$ line profiles (Zheng 1992), and Balmer and \\ion{O}{1} $\\lambda$8446 profiles (Morris \\& Ward 1989).  While at any given continuum state the broad emission-line region may well contain clouds fully ionized in hydrogen, as the present model does, the results presented here suggest that there is no great need for a major line-emitting component consisting of vast amounts of optically thin gas. Most of the effects mentioned just above regarding the Balmer and helium recombination lines can be explained by the effects we described in $\\S$~3: weak Balmer line emission and response need not originate in clouds optically thin to the Lyman continuum. While an emission line such as \\ion{Ne}{8} $\\lambda$774 is likely to be emitted mainly within clouds that lack a hydrogen ionization front (e.g., see Hamann et~al.\\ 1998; Korista et~al.\\ 1997), we have also pointed out that by their very nature fully-ionized clouds are inefficient emitters of hydrogen recombination lines (refer again to Figure~1), and so are unlikely to be important energetically even if their covering fraction is large. This is especially true in the face of a variable ionizing continuum source; clouds will more likely be either ionization-bounded or too overionized to emit Balmer lines of any significance. It is also significant that the matter-bounded clouds in Figure~4 either have small positive ({\\em lightest gray})\\footnote{The {\\em only} clouds in Figure~4 whose hydrogen line responsivities are $\\approx$~0 lie above the diagonal line $\\log (U_Hc) \\approx 10.25$ and also within the lightest gray shaded region.} or negative (white) responsivities. Even if present, the latter type clouds never make their presence known in these emission lines, which is not surprising given their extremely low emitting efficiencies.  In regards to the comparisons of Balmer and Ly$\\alpha$ line profiles, as mentioned in $\\S$~3.4, the EW contours of Ly$\\alpha$ in the density-flux plane (see Figure~1, {\\em bottom right}) do not show the strong continuum flux dependencies of the Balmer lines. We suspect that this, along with the presence of a radially decreasing velocity field, accounts for the very large Ly$\\alpha$/H$\\beta$ ratios found in the emission-line wings by Zheng (1992), rather than these ratios being due to a major constituent of optically thin emitting gas.  Finally, we also point out that many of the above studies assumed that case~B emissivities apply to the hydrogen lines emitted by optically thick clouds in the BLR and attributed significant deviations from expected ratios to an optically thin emitting region. We conclude that a separate optically thin VBLR may be unnecessary, and in any case for the Balmer lines is probably energetically unimportant. The wings of the hydrogen recombination lines likely represent nothing more than the inner BLR where the Balmer line emission is inefficient and the responsivity low.   \\subsection{Effects of an SED that Varies with Continuum Level}  Paltani \\& Courvoisier (1994) and Paltani \\& Walter (1996) showed unambiguously that the UV continuum becomes bluer when the flux is higher in variable AGNs. This effect has also been found in the Sloan Digital Sky Survey quasar sample (Vanden Berk et~al.\\ 2004). In the particular case of NGC~5548, Maoz et~al.\\ (1993) reported that, after accounting for the Balmer continuum and UV \\ion{Fe}{2} emission, the 1800--2400~\\AA\\/ continuum becomes bluer with increased continuum flux. Additionally, Peterson et~al.\\ (2002) and Gilbert \\& Peterson (2003) found that after careful removal of the contaminating non-variable stellar optical continuum flux, the UV-optical continuum in NGC~5548 becomes bluer as it becomes brighter: $F_{5100} \\propto F_{1350}^{0.67}$. Extrapolation of this relation to higher energies might indicate a hardening of the incident ionizing continuum during brighter continuum states. Marshall et~al.\\ (1997) found larger amplitude variations in the extreme ultraviolet ($\\sim$~100~\\AA\\/) than at 1350~\\AA\\/ during a week-long period of the 1993 NGC~5548 monitoring campaign. Alternatively, as suggested by Paltani \\& Walter (1996), it may indicate the presence of a constant (or more slowly varying) continuum or pseudocontinuum component whose contribution is greater at longer wavelengths. Korista \\& Goad (2001) found that 20\\%--30\\% of the observed effect may be due to reverberation of the diffuse continuum emission of the broad line clouds. In the simulations presented here, the SED of the continuum incident on the broad-line clouds is the same for both high and low continuum states. How an emission line's responsivity would change with an {\\em incident} SED that changed shape as the UV-optical continuum level varied will depend on the details of how the flux of photons important to the creation and destruction of the line varies. We briefly discuss these effects here.  The shape of the ionizing continuum determines the heating for given values of gas density, ionizing photon flux, and chemical abundances. All else being equal, harder incident continua result in higher electron temperatures, and so major cooling lines (e.g., \\ion{C}{4} $\\lambda$1549) might be expected to increase their energy output in response. One should also keep in mind that the incident flux at energies corresponding to the {\\em Balmer} continuum is important to the destruction of hydrogen emission lines for the high gas densities encountered in the BLR (Shields \\& Ferland 1993). Also, harder continua during brighter UV-opt continuum states might also mean comparatively more photons above 1~ryd and/or 4~ryd, favoring the production of hydrogen and helium recombination lines. Kaspi \\& Netzer (1999) considered an SED whose break near 4~ryd moves to higher energies for brighter UV continuum states (their Figure~8d), with the effect of increasing the 4~ryd/1~ryd photon flux ratio with increasing $\\lambda$1350 continuum luminosity. However, in this variable SED scheme, the flux above 20~ryd (270~eV) remained constant. A comparison of their Figs.~7 (fixed SED) and 9 (variable SED) that show their model predictions of five UV emission lines' light curves versus observations illustrates some of the aforementioned effects of a variable continuum SED on line responsivity. While they included \\ion{He}{2}~$\\lambda$1640 in their simulations, such was not the case for the Balmer lines. Their prescription for a variable continuum SED increased the responsivities of the \\ion{He}{2} and Ly$\\alpha$ lines, whereas those of \\ion{C}{4}~$\\lambda$1549, \\ion{C}{3}] $\\lambda$1909, and \\ion{Mg}{2} $\\lambda$2800 decreased, presumably because the variations in the continuum beyond $\\sim$~150~eV were smaller in their variable SED scheme than for the constant SED. It is also apparent that the lags of \\ion{He}{2} and \\ion{C}{4} increased in the higher continuum states, because of a shift to larger responsivity-weighted radii.  \\subsection{The Importance of the Recombination Lines to the Method of Quasar Tomography}  The method of ``quasar tomography,'' proposed by Horne et~al.\\ (2003), is a unification of reverberation mapping and photoionization physics. Its goal is to place physical constraints on the recovered two-dimensional transfer function map $\\Psi(\\tau, v)$ of the broad emission lines, as well as to recover the gas distribution described by a five-dimensional map $f(r, \\theta, n_H, N_H, v)$, where $\\tau$ is the time delay, $v$ is the observed radial velocity, $r$ is the distance from the continuum source, and $\\theta$ is the angle from the line of sight. The sensitivity of the emissivities (EW contours in Figure~1) of the optical recombination lines to the continuum flux, and the resulting consequences to their responsivities, make these lines important candidates for inclusion in the quasar tomography. Horne et~al.\\ found that when applied to the mostly collisionally excited UV emission lines, a certain amount of ambiguity arises in the recovery of the maps due to the weak dependencies in their EWs (and so responsivities) along lines of constant ionization parameter ($\\propto \\Phi_H/n_H$) corresponding to their maximum emissivities. Their emissivities and responsivities are largely dependent on $U_H$, with smaller dependencies in gas density due to de-excitation or thermalization. The optical recombination lines, along with \\ion{Mg}{2} $\\lambda$2800, offer additional constraints in determining the physical conditions within the BLR, in that their emissivities and responsivities are dependent on the local continuum flux and so radius for a fixed central continuum luminosity, as well as changes in this flux as the central continuum source varies."
    },
    "0402/hep-ph0402033_arXiv.txt": {
        "abstract": "{\\\\ $~~~$ The decoupling of a cold relic, during a decaying-particle-dominated cosmological evolution is analyzed, the relic density is calculated both numerically and semi-analytically and the results are compared with each other. Using plausible values (from the viewpoint of supersymmetric models) for the mass and the thermal averaged cross section times the velocity of the cold relic, we investigate scenaria of equilibrium or non-equilibrium production. In both cases, acceptable results for the dark matter abundance can be obtained, by constraining the reheat temperature of the decaying particle, its mass and the averaged number of the produced cold relics. The required reheat temperature is, in any case, lower than about $20~{\\rm GeV}$.   \\\\ \\\\{\\sc Keywords}: Cosmology, Dark Matter \\\\ {\\sc PACS codes}: 98.80.Cq, 95.35.+d \\\\ \\\\ {\\sl\\bfseries Published in} {\\sl Astropart. Phys.} {\\bf 21}, 689 (2004)}    \\addtolength{\\textheight}{.5cm}   \\begin{document}  \\setcounter{page}{1} \\pagestyle{fancyplain}  \\addtolength{\\headheight}{.5cm}  \\rhead[\\fancyplain{}{ \\bf \\thepage}]{\\fancyplain{}{M{\\ftn ASSIVE} P{\\ftn ARTICLE} D{\\ftn ECAY AND} C{\\ftn OLD} D{\\ftn ARK} M{\\ftn ATTER} A{\\ftn BUNDANCE}}} \\lhead[\\fancyplain{}{ \\leftmark}]{\\fancyplain{}{\\bf \\thepage}} \\cfoot{} ",
        "introduction": "}\\label{sec:intro}  \\hspace{.562cm} The enigma of the Cold Dark Matter (CDM) constitution of the universe becomes more and more precisely defined, after the recently announced WMAP results \\cite{wmap, wmapl}, which determine the CDM abundance, $\\Omega_{\\rm CDM}h^2$, with an unprecedented accuracy: \\beq \\Omega_{\\rm CDM}h^2=0.1126_{-0.0181}^{+0.0161} \\label{cdmba}\\eeq at $95\\%$ confidence level. In light of this, the relic density of any CDM candidate $\\tilde\\chi$ (i.e, which decouples being non relativistic), $\\Omega_{\\tilde\\chi}h^2$, is to satisfy a very narrow range of values: \\beq {\\sf (a)}~~0.09\\lesssim \\Omega_{\\tilde\\chi}h^2~~\\quad\\mbox{and}\\quad~~ {\\sf (b)}~~\\Omega_{\\tilde\\chi}h^2\\lesssim0.13 \\label{cdmb}\\eeq which tightly restricts the parameter space of the theories which support the existence of $\\tilde\\chi$. The most popular of these are the $R$-parity conserving supersymmetric ({\\sc SUSY}) theories which identify $\\tilde\\chi$ to the stable lightest {\\sc SUSY} particle ({\\sc LSP}) \\cite{goldberg}. According to the standard scenario, (i) $\\tilde\\chi$ decouples from the cosmic fluid during the radiation-dominated (RD) era, (ii) being in chemical equilibrium with plasma and (iii) produced through thermal scatterings in the plasma. The condition (ii) is satisfied naturally by the lightest neutralino of the minimal {\\sc SUSY} standard model ({\\sc MSSM}), which turns out to have the required strength of interactions, being weakly interacting.  Although quite compelling, this scenario comes across with difficulties, especially when it is applied in the context of economical and predictive versions of {\\sc MSSM}. E.g., in most of the parameter space of the Constrained {\\sc MSSM} ({\\sc CMSSM}) \\cite{Cmssm}, $\\tilde\\chi$ turns out to be bino and its $\\Omega_{\\tilde\\chi}h^2$ exceeds the bound of Eq. (\\ref{cdmb}{\\sf b}). Several suppression mechanisms of $\\Omega_{\\tilde\\chi}h^2$ have been proposed so far: Bino-sleptons \\cite{ellis2} and particularly, for large $\\tan\\beta$, bino-stau \\cite{cdm}, bino-stops \\cite{boem, santoso1} or bino-chargino \\cite{darkn} coannihilations and/or A-pole effect \\cite{lah} can efficiently reduce $\\Omega_{\\tilde\\chi}h^2$. Also several kinds of non-universality in the Higgs \\cite{ellis3} and/or gaugino \\cite{edjo, nelson} and/or sfermionic sector \\cite{su5b} can help in the same direction, creating additional coannihilation effects. As is expected, a more or less tuning of the {\\sc SUSY} parameters is needed in these cases, without a simultaneous satisfaction of other phenomenological constraints to be always possible (see, e.g. Ref. \\cite{wmapl}). On the other hand, $\\Omega_{\\tilde\\chi}h^2$ turns out to be lower than the bound of Eq. (\\ref{cdmb}{\\sf a}) in other models, e.g., in the anomaly mediated {\\sc SUSY} breaking model ({\\sc AMSBM}) where $\\tilde\\chi$ is mostly wino \\cite{wells} or in models based on $SU(5)$ gaugino non-universality \\cite{roy}, where $\\tilde\\chi$ can be higgsino. Then we have to invoke another CDM candidate \\cite{axino}, in order for the range of Eq. (\\ref{cdmb}) to be fulfilled.  However, this picture can dramatically change, if the standard assumption (i) is lifted. Indeed, since there is no direct information for the history of the Cosmos before the epoch of nucleosynthesis (i.e., temperatures $T>1~{\\rm MeV}$) the decoupling of $\\tilde\\chi$ can occur not in the RD era. As was pointed out a lot years ago \\cite{McDonald} and, also, recently \\cite{riotto, fornengo}, the ${\\tilde\\chi}$ decoupling can be related to the decay of a massive scalar particle. The modern cosmo-particle theories are abundant in such fields, e.g. inflatons \\cite{chung, kubo, drees, dreesa}, dilatons or moduli \\cite{quevedo, moroi}, Polonyi field \\cite{yanagida}, $q$-balls \\cite{fujii} (see, also \\cite{rachel}). During their decay, these particles perform coherent oscillations ``reheating'' the universe \\cite{turner}. This phenomenon is not instantaneous \\cite{kolb, chung}. In particular, the maximum temperature during this period is much larger than the so-called reheat temperature, which can be better considered as the largest temperature of the RD era \\cite{riotto}. Consequently, the ``freeze out'' of $\\tilde\\chi$ could be realized before the completion of the reheating. The cosmological evolution during this phase is strongly modified as regards the standard one \\cite{turner}, with crucial consequences to $\\Omega_{\\tilde\\chi}h^2$ calculation \\cite{riotto, fornengo, moroi, shaaban}. Namely, two types of $\\tilde\\chi$-production emerge, in contrast with (ii): The chemically equilibrium (EP) and the non-equilibrium production (non-EP) (in both cases, kinetic equilibrium of $\\tilde\\chi$'s is assumed \\cite{riotto}). In this paper we extend the analysis in Ref. \\cite{riotto}, lifting also the assumption (iii) of the standard scenario: We include the possibility (which, naturally arises even without direct coupling \\cite{drees, dreesa}) that the decaying particle can decay to $\\tilde\\chi$. The problem has already been faced semi-quantitatively in Refs \\cite{moroi, shaaban, drees, dreesa} and numerically in Ref. \\cite{brazil} for significantly more massive $\\tilde\\chi$'s.  Our numerical and semi-analytical analyses are exposed in secs \\ref{sec:num} and \\ref{sec:neut}. The obtained results are compared with each other in sec. \\ref{appl}. There, we realize, also, a model independent application of our findings in the case of {\\sc SUSY} models inspired ${\\tilde\\chi}$ masses and cross sections. We find that comfortable satisfaction of Eq. (\\ref{cdmb}) can be achieved, by constraining the reheat temperature to rather low values, the mass of the decaying particle and the averaged number of the produced $\\tilde\\chi$'s, without any tuning of the {\\sc SUSY} parameters.  Throughout the text and the formulas, brackets are used by applying disjunctive correspondence and natural units ($\\hbar=c=k_B=1$) are assumed. ",
        "conclusions": "-O{\\ftn PEN} I{\\ftn SSUES}}\\label{con}  \\hspace{.562cm} We considered a deviation from the standard cosmological situation according to which the CDM candidate, $\\tilde\\chi$ decouples from the plasma (i) during the RD era (i.e. after reheating) (ii) being in equilibrium (iii) produced through thermal scatterings. On the contrary, we assumed that $\\tilde\\chi$ decoupling occurs (i$^\\prime$) during a decaying-massive-particle, $\\phi$, dominated era (and mainly before reheating) (ii$^\\prime$) being or not in chemical equilibrium with the thermal bath (iii$^\\prime$) produced by thermal scatterings and directly from the $\\phi$ decay.  We solved the problem (i) numerically, integrating the relevant system of the differential equations (ii) semi-analytically, producing approximate relations for the cosmological evolution before reheating and solving the properly re-formulated Boltzmann equations.   \\addtolength{\\textheight}{1.cm} \\newpage  \\begin{figure}[!ht]\\vspace*{-.15in} \\hspace*{-.71in} \\begin{minipage}{8in} \\epsfig{file=NTbn.eps,height=3.8in,angle=-90} \\hspace*{-1.37 cm} \\epsfig{file=NTwn.eps,height=3.8in,angle=-90} \\hfill \\end{minipage}\\vspace*{-.01in} \\hfill\\hspace*{-.71in} \\begin{minipage}{8in} \\epsfig{file=Nmbn.eps,height=3.8in,angle=-90} \\hspace*{-1.37 cm} \\epsfig{file=Nmwn.eps,height=3.8in,angle=-90} \\hfill \\end{minipage}\\vspace*{-.01in} \\hfill\\hspace*{-.71in} \\begin{minipage}{8in} \\epsfig{file=Tmbn.eps,height=3.8in,angle=-90} \\hspace*{-1.37 cm} \\epsfig{file=Tmwn.eps,height=3.8in,angle=-90} \\hfill \\end{minipage}\\vspace*{-.05in} \\hfill \\caption[]{\\sl Regions allowed by Eq. (\\ref{cdmb}) on the $T_{\\rm RH}-N_{\\tilde\\chi}$ plane ${\\sf (a_1, a_2)}$ for $m_\\phi=10^6~{\\rm GeV}$, $m_\\phi-N_{\\tilde\\chi}$ plane ${\\sf ( b_1 ~[b_2])}$ for $T_{\\rm RH}=0.05~[5]~{\\rm GeV}$ and $m_\\phi-T_{\\rm RH}$ plane ${\\sf (c_1~[c_2])}$ for $N_{\\tilde\\chi}=10^{-6[3]}$. We take $m_{\\tilde\\chi}=200~[500]~{\\rm GeV}$ (light grey [normal grey] areas) and $\\langle\\sigma v\\rangle=10^{-12}~[10^{-8}]~{\\rm GeV^{-2}}~{\\sf (a_1, b_1, c_1~[a_2, b_2, c_2])}$.} \\label{regions} \\end{figure}  \\addtolength{\\textheight}{-1.cm} \\newpage    The second way facilitates the understanding of the problem and gives, in most cases, sufficiently accurate results, provided a suitable $\\delta_{\\rm RH}$ is chosen. The variation of $\\Omega_{\\tilde\\chi}h^2$ w.r.t our free parameters $(m_\\phi, N_{\\tilde\\chi},T_{\\rm RH})$ was investigated and regions consistent with the present CDM bounds are constructed, using $m_{\\tilde\\chi}$'s and $\\langle\\sigma v\\rangle$'s commonly allowed in SUSY models.  These scenaria obviously let intact the SUSY parameter space but require rather low $T_{\\rm RH}$. This can be accommodated in AMSBM \\cite{moroi}, in models with intermediate scale unification \\cite{shaaban} and within the context of $q$-balls decay \\cite{fujii}. Also, low $T_{\\rm RH}$ naturally arises in models of thermal inflation  \\cite{lyth} which, as a bonus, may overcome the problem of unwanted relics (e.g., gravitino, moduli). Finally, restrictions on $T_{\\rm RH}$ arising from  baryogenesis and neutrino cosmology have been, also, studied in Refs \\cite{riotto1, neutrino}.  Our formalism can be easily extended to include coannihilations and pole effects. Therefore, it can become applicable for the calculation of $\\Omega_{\\tilde\\chi}h^2$ in the context of specific SUSY models. Also, these scenaria can assist us to the reduction of $\\Omega_{\\tilde\\chi}h^2$ in cases, where it turns out to be even more enhanced than in the standard scenario, as in the presence of Quintessence \\cite{salati}. Similar analysis of the $\\tilde\\chi$-decoupling during the extra dimensional cosmological evolution \\cite{extra} may be also, possible."
    },
    "0402/astro-ph0402056_arXiv.txt": {
        "abstract": "Gravitational waves from the coalescence of binary black holes carry away linear momentum, causing center of mass recoil.  This ``radiation rocket'' effect has important implications for systems with escape speeds of order the recoil velocity.  We revisit this problem using black hole perturbation theory, treating the binary as a test mass spiraling into a spinning hole.  For extreme mass ratios ($q\\equiv m_1/m_2 \\ll 1$) we compute the recoil for the slow inspiral epoch of binary coalescence very accurately; these results can be extrapolated to $q \\sim 0.4$ with modest accuracy.  Although the recoil from the final plunge contributes significantly to the final recoil, we are only able to make crude estimates of its magnitude.  We find that the recoil can easily reach $\\sim 100-200\\,{\\rm km/s}$, but most likely does not exceed $\\sim 500\\,{\\rm km/s}$.  Though much lower than previous estimates, this recoil is large enough to have important astrophysical consequences. These include the ejection of black holes from globular clusters, dwarf galaxies, and high-redshift dark matter halos. ",
        "introduction": "Along with energy and angular momentum, gravitational waves (GWs) carry {\\it linear} momentum away from a radiating source {\\citep{br61,peres62,bek73}}. Global conservation of momentum requires that the center of mass (COM) of the system recoil. This recoil is independent of the system's total mass.   {\\citet{f83}} first computed GW recoil for binaries.  He treated the members as non-spinning point masses ($m_1,m_2$), the gravitational force as Newtonian, and included only the lowest GW multipoles needed for momentum ejection.  For circular orbits Fitchett's recoil is \\begin{equation} V_F \\simeq 1480\\,\\mbox{km/s}\\,{f(q)\\over f_{\\rm max}}\\left({2 G M / c^2 \\over r_{\\rm term}}\\right)^4\\;, \\label{eq:vfitchett} \\end{equation} where $r_{\\rm term}$ is the orbital separation where GW emission terminates, $q = m_1/m_2 \\le 1$ is the mass ratio, and $M = m_1 + m_2$ is the total mass.  The function $f(q)=q^2(1-q)/(1+q)^5$ has a maximum $f_{\\rm max}$ at $q \\simeq 0.38$, is zero for $q = 1$, and has the limit $f(q) \\approx q^2$ for $q \\ll 1$.  Equation (\\ref{eq:vfitchett}) tells us that in the coalescence of binary black holes (BHs)---where $r_{\\rm term}$ can approach $GM/c^2$---the kick might reach thousands of km/s.  This is far greater than the escape velocity of many globular clusters (typically $\\sim 30$ km/s), and may even exceed galactic escape velocities ($\\sim 1000$ km/s).  Recoil could thus have important astrophysical implications {\\citep{rr89}} [some of which are explored in a companion paper (\\citealt{paperII}; Paper II)]. This has motivated us to revisit this problem.   Equation (\\ref{eq:vfitchett}) indicates that the recoil is strongest at small separations, when the relativistic effects neglected by Fitchett are most important. This issue has been addressed in restricted circumstances using perturbation theory {\\citep{nh83,fd84,nok87}}, post-Newtonian expansions {\\citep{agw92,kidder}}, and numerical relativity {\\citep{ap97,ABprl,BAprd,lp04}}. Unlike previous analyses, our treatment applies to the strong-gravity, fast-motion regime around spinning holes undergoing binary coalescence. Using BH perturbation theory we model the dynamics of the binary, the generation of GWs, and the backreaction of those waves on the system up to the inner-most stable circular orbit (ISCO). Our results are accurate only for extreme mass ratio inspirals ($q\\ll 1$), but we can extrapolate to $q \\sim 0.4$ with modest error. We model the GW emission from the final plunge more crudely. ",
        "conclusions": "The punchline of this analysis is simple: quasi-Newtonian estimates have significantly overestimated the kick velocity from anisotropic GW emission during binary coalescence.  The recoil is strongest when the smaller member is deep in the strong-field of the large black hole.  General relativistic effects, such as the gravitational redshift and spacetime curvature-scattering, act on the emitted GWs and reduce the recoil.  Though reduced, the recoil remains large enough to have important astrophysical consequences. Recoils with $V \\sim 10$--$100\\, {\\rm km/s}$ are likely; kicks of a few hundred km/s are not unexpected; and the largest possible recoils are probably $\\lesssim 500\\,{\\rm km/s}$.  These speeds are smaller than most galactic escape velocities, suggesting that BH mergers that follow galaxy mergers will remain within their host structures.  However, these recoils are similar to the escape speeds of dwarf galaxies; and they may be sufficient to escape from mergers in high redshift structures [$z\\gtrsim 5-10$; cf.\\ {\\citet{bl01}}, Fig.\\ 8]. Binary BH ejection from globular clusters is quite likely, with significant implications for the formation of intermediate mass black holes (IMBH) via hierarchical mergers {\\citep{mc03}}. Our recoil estimates will also be useful in simulations of supermassive and IMBH evolution in dark halos {\\citep{volonteri1,volonteri2}}.  Future papers will present the formalism used for this analysis, and will investigate the influence of orbital inclination on the recoil. More work in perturbation theory also remains in addressing the recoil from the plunge and final ringdown of the merging black holes.  Finally, \\citet{rr89} have speculated that spin-orbit misalignment could lead to recoil directed out of the orbital plane. This recoil might accumulate secularly rather than oscillate, and would be similar to the ``electromagnetic rocket'' in pulsars with off-centered magnetic dipole moments {\\citep{ht75,lcc01}}.  We suspect that this effect occurs but it is likely small compared to the recoil from the final plunge and merger. Firm estimates of the final kick velocity will rely on correctly modelling the final phase of BH coalescence.  For comparable mass binaries, full numerical relativity will ultimately be needed to accurately compute the GW recoil."
    },
    "0402/astro-ph0402260_arXiv.txt": {
        "abstract": "{A 1-D loop radiative hydrodynamic model that incorporates the effects of gravitational stratification, heat conduction, radiative losses, external heat input, presence of helium, and Braginskii viscosity is used to simulate elementary flare loops. The physical parameters for the input are taken from observations of the Bastille-Day flare of 2000 July 14. The present analysis shows that: (a) The obtained maximum values of the electron density can be considerably higher ($4.2\\times 10^{11}$ cm$^{-3}$ or more) in the case of footpoint heating than in the case of apex heating ($2.5\\times 10^{11}$ cm$^{-3}$). (b) The average cooling time after the flare peak takes less time in the case of footpoint heating than in the case of apex heating. (c) The peak apex temperatures are significantly lower (by about 10 MK) for the case of footpoint heating than for apex heating (for the same average loop temperature of about 30 MK). This characteristic allows us to discriminate between different heating positioning. (d) In both cases (of apex and footpoint heating), the maximum obtained apex temperature $T^{max}$ is practically independent of the heating duration $\\sigma_{t}$, but scales directly with the heating rate $E_{H0}$. (e) The maximum obtained densities at the loop apex, $n_e^{max}$, increase with the heating rate $E_{H0}$ and heating duration $\\sigma_{t}$ for both footpoint and apex heating. In Paper II we will use the outputs of these hydrodynamic simulations, which cover a wide range of the parameter space of heating rates and durations, as an input for forward-fitting of the multi-loop arcade of the Bastille-day flare. ",
        "introduction": "Solar flares are complex systems that involve many magnetic field lines and thus can rarely be represented by a single flux-tube. During the Bastille-Day 2000 July 14 flare for instance, which exhibits a classical double-ribbon flare configuration, an ensemble of over 200 individual post-flare loops has been identified (Aschwanden \\& Alexander 2001). Each of these individual loops has its own hydrodynamic evolution during a flare, occurring in magnetic flux systems that are thermodynamically isolated from each other and have their own independent timing and physical parameters. Hydrodynamic modeling of flare loops, however, have been performed for single flare loops \\citep{m87, m89}, but only few MHD simulation studies have been orchestrated in a multi-loop configuration \\citep{h97, h98}. Even the multi-loop simulations have been designed only in the simplest way, by assuming regular spacing and time intervals to merely mimic the superposition effect, but no detailed fitting of the observed spatial configurations and timing has ever been attempted. A rigorous hydrodynamic modeling effort for complex large flares would be extremely valuable to constrain the total energy budget, the heating functions, the fractal geometric structure, the plasma filling factors, and the spatio-temporal organization of {\\it unsteady} \\citep{p00} or {\\sl impulsive bursty magnetic reconnection} processes \\citep{l82}, which are likely to occur in large double-ribbon flares because of the large shear and resulting tearing mode instability \\citep{s66}.  In this series of papers we present a method of radiative hydrodynamic modeling of large, complex, multi-loop flares. In this Paper I we perform numerical simulations with a 1-D hydrodynamic code to obtain the temperature evolution $T_e^{max}(t), T_e^{avg}(t)$ and density evolution $n_e^{max}(t)$ in a large parameter space of heating functions. We vary the maximum heating rate, heating duration, and location (footpoint, apex) of the heating functions $E_H(t)$. In Paper II we parameterize the results, suitable to forward-fitting of a multi-loop system and fit the multi-wavelength data of the Bastille-day flare 2000 July 14 using TRACE, Yohkoh/SXT, HXT, and GOES data. In subsequent papers we plan to extend the hydrodynamic results of this flare to constrain magnetic modeling, magnetic reconnection geometries, and particle acceleration processes. ",
        "conclusions": "In summary, we have used a  radiative hydrodynamic numerical code to simulate flares. The physical parameters of the input were obtained from observations of the the Bastille-day flare \\citep{aa01}. Our simulations confirm the general picture of flare dynamics: Transient heat deposition either at loop footpoints (with maximum heat input at the bottom of the region connecting corona to the chromosphere or i.e. top of chromosphere) or at the apex leads to an average loop temperature of $T_{avg} \\approx 30$ MK first. Then, evaporation of material from the chromosphere and the region connecting corona to the chromosphere into corona ensues with up-flows in the order of a few hundreds of km s$^{-1}$. During the peak of the flare, the combined action of heat input and conductive and radiate loss yields an oscillatory flow pattern with typical amplitudes of up to few tens of km s$^{-1}$. Finally we enter a cooling phase, when down-flows in the order of hundred km s$^{-1}$ can be seen as the plasma drains out of the loop, ultimately reaching an equilibrium.  We have established the following: \\begin{enumerate} \\item In the case of footpoint heating, the obtained maximum values of the density are considerably higher ($4.2\\times 10^{11}$ cm$^{-3}$ or more) than in the case of apex heating ($2.5\\times 10^{11}$ cm$^{-3}$). This is due to the fact that footpoint heating is more efficient in evaporating material from the region connecting corona to the chromosphere and chromosphere itself, which yields denser loops during the flare. In the case of apex heating, which was used to model flares, insufficient downward heating conduction prevents significant material evaporation. \\item In the case of footpoint heating, as compared to the apex heating, on average cooling after the flare takes less time. This due to the fact that the time scale of conduction loss is proportional to the density, while the time scale of radiative loss is reciprocal to the density  \\citep{aa01}. Therefore, since radiative losses dominate over heat conduction losses for most of the time, it is clear that  the denser loops cool faster. \\item In principle, {\\it our simulations would allow to discriminate} between different heating functions of the loop during the flare, if  one would have temperature dynamics in a given point of the loop, such as at the apex. This is based on our observation that in the case of footpoint heating the peak apex temperatures ({\\it corresponding to the same, as in the case of apex heating, average temperature of about 30 MK}) are significantly lower (less by about 10 MK). \\item In the case of footpoint heating, up-flow velocities are higher (roughly up to  100 km s$^{-1}$) than in the case of apex heating due to more efficient evaporation. \\item In both cases (of apex and footpoint heating) the maximum obtained temperature $T^{max}$ at the loop apex is practically independent of the heating duration $\\sigma_t$, but it increases with higher heating rates $E_{H0}$. \\item  The maximum obtained densities at the loop apex increase with the increase of both the flare heating rate $E_{H0}$ and the heating duration $\\sigma_t$, in the case of apex as well as footpoint heating. \\item Varying the loop length (see Table 2) in the range of $L=(0.25,...,2.0)\\times L_0$ (with $L_0=55$ Mm), we find (1) that the mean loop temperature averaged over the loop length does not change dramatically, (2) that the loop apex temperature increases notably for longer loops only for the case of apex heating, but much less for footpoint heating, and (3) that the mean electron density decreases somewhat with longer loops, i.e., $n_e/10^{11}$ (cm$^{-3}$) $=10.45 L^{-0.362}$ for apex heating and $n_e/10^{11}$ (cm$^{-3}$) $=9.62 L^{-0.207}$ for footpoint heating. Here, $L$ is loop length in Mm.  \\end{enumerate}  In practically all of our numerical runs we have detected quasi-periodic oscillations in all physical quantities. In fact, such oscillations are frequently seen during solar flares (e.g. \\citet{terekhov02}) as well as stellar flares (e.g. \\citet{mathio}). Our preliminary analysis shows that quasi-periodic oscillations seen in our numerical simulations bear many similar features as the observed ones. The key point is that the traditional explanation of these oscillations in the observations involves {\\it MHD waves}. In the numerical simulations presented here, however, they are likely to be produced by standing sound waves caused by impulsive and localized heating. Therefore, our explanation of these oscillations is purely hydrodynamic -- they are related to the standing slow mode acoustic waves, similar to the observed by SUMER \\citep{wang02}. A detailed study of these quasi-periodic oscillations will be presented elsewhere.  In a next step we plan to use the outputs of this parametric study of hydrodynamic simulations, which cover a wide parameter range of heating rates and heating time scales, as input for forward-fitting to the observed physical parameters (densities and temperatures) of the multi-loop flare on Bastille-day 2000 July 14."
    },
    "0402/astro-ph0402110_arXiv.txt": {
        "abstract": "We describe new optical images and spectra of POX 52, a dwarf galaxy with an active nucleus that was originally detected in the POX objective-prism survey.  While POX 52 was originally thought to be a Seyfert 2 galaxy, the new data reveal an emission-line spectrum very similar to that of the dwarf Seyfert 1 galaxy NGC 4395, with broad components to the permitted line profiles, and we classify POX 52 as a Seyfert 1 galaxy.  The host galaxy appears to be a dwarf elliptical, and its brightness profile is best fit by a S\\'ersic model with an index of $3.6\\pm0.2$ and a total magnitude of $M_V = -17.6$.  Applying mass-luminosity-linewidth scaling relations to estimate the black hole mass from the broad H$\\beta$ linewidth and nonstellar continuum luminosity, we find $M_{\\mathrm{BH}} \\approx 1.6\\times10^5 ~M_{\\sun}$. The stellar velocity dispersion in the host galaxy, measured from the \\ion{Ca}{2} $\\lambda8498, 8542$ \\AA\\ lines, is $36\\pm5$ km s$^{-1}$, also suggestive of a black hole mass of order $10^5 ~M_{\\sun}$. Further searches for active nuclei in dwarf galaxies can provide unique constraints on the demographics of black holes in the mass range below $10^6 ~M_{\\sun}$. ",
        "introduction": "\\begin{figure*} \\begin{center} \\rotatebox{-90}{\\scalebox{0.66}{\\includegraphics{fullspec.ps}}} \\end{center} \\caption{Keck ESI spectrum of POX 52. \\label{fullspec}} \\end{figure*}  Studies of the dynamics of stars and gas in the nuclei of nearby galaxies have detected the signature of supermassive black holes in an ever-increasing number of galaxies \\citep[for recent reviews, see][]{kor04, bar04}, and it is now widely accepted that all galaxies with a massive bulge component contain a central black hole.  While existing surveys have made great progress in the black hole census for masses in the range $\\sim2\\times10^6 - 3\\times10^9$ \\msun, very little is known about the population of black holes with masses below $10^6$ \\msun.  Stellar-dynamical and gas-dynamical searches do not have sufficient sensitivity to detect black holes of mass $\\lesssim10^6$ \\msun\\ for galaxies much beyond the Local Group. ``Intermediate-mass'' black holes with masses of $\\sim10^3 - 10^6$ \\msun\\ might be present in the centers of very late-type spiral galaxies, in dwarf elliptical and dwarf spheroidal galaxies, and possibly even in massive globular clusters, but thus far dynamical studies have only been carried out for a few very nearby objects, sometimes with ambiguous results \\citep{vdm04}.  At present, there is little hope of obtaining a systematic census of black holes with $M < 10^6$ \\msun\\ by using traditional dynamical detection techniques.  The demographics of intermediate-mass black holes are of particular importance for future gravitational-wave studies, since it is expected that the most common massive black hole merger events detected by the space-based \\emph{LISA} interferometer will be in the $10^5-10^6$ \\msun\\ range \\citep{hug01}.  Intermediate-mass black holes that would be undetectable by dynamical measurements might still reveal their presence by their accretion luminosity.  The detection of a low-luminosity active galactic nucleus (AGN) in a dwarf galaxy would be a good indication that a black hole is present, and indirect methods could be used to estimate the black hole mass even if a stellar-dynamical mass measurement is impossible. However, only two AGNs in dwarf galaxies have been identified previously.  The better-known case is the nearby ($D \\approx 4$ Mpc), late-type, dwarf spiral galaxy NGC 4395.  \\citet{fs89} found that NGC 4395 contains the least luminous known Seyfert 1 nucleus.  Despite the fact that this galaxy has no discernible bulge, and thus might not be expected to contain a central black hole, its nucleus exhibits all of the characteristics of Seyfert activity, including broad emission lines \\citep{fs89}, a compact, unresolved nonstellar continuum source in the optical and ultraviolet \\citep{fhs93}, unresolved radio emission with a high brightness temperature \\citep{wfh01}, and rapid X-ray variability \\citep{iwa00, sif03, mor04}.  The central stellar velocity dispersion is $\\sigmastar < 30$ \\kms, which combined with the radius of the central star cluster (3.9 pc) gives a firm upper limit of $M < 6\\times10^6$ \\msun\\ for the combined mass of the central star cluster plus black hole \\citep{fh03}.  The black hole itself has a likely mass of $\\sim10^4-10^5$ \\msun, based on the X-ray variability timescale \\citep{sif03} and extrapolation of the linewidth-luminosity-mass correlations for Seyfert galaxies \\citep{fh03}.  Finding more active galaxies like NGC 4395 would be a valuable step toward understanding the demographics of black holes with $M < 10^6$ \\msun.  POX 52 (also known as PGC 038055 or G1200--2038) is a dwarf galaxy at $cz = 6533$ \\kms.  It was discovered in the course of an objective-prism search for emission-line objects performed by \\citet{ksk81}.  They listed the fluxes for several emission lines in POX 52 and noted that it appears ``star-like'' in Palomar Sky Survey images.  Follow-up spectroscopy and imaging data were presented by \\citet[][hereafter KSB]{ksb87}.  KSB called attention to POX 52 as an unusual example of a dwarf galaxy having an AGN spectrum.  The object was classified as a Seyfert 2 based on the narrow-line ratios, although they noted that the \\hbeta\\ line had a weak broad component with full-width at half-maximum (FWHM) $\\approx840$ \\kms.  They also found a weak broad component to the \\ion{He}{2} $\\lambda4686$ emission line, which they tentatively ascribed to emission from Wolf-Rayet stars. The host galaxy was found to have an exponential scale length of only 1 kpc and an absolute magnitude of $M_V=-16.9$.  This is a surprisingly small host galaxy for an AGN.  However, a literature search reveals that no additional optical observations of POX 52 have been published since the initial study by KSB.  Motivated by the possibility that POX 52 might contain an intermediate-mass black hole similar to the one in NGC 4395, we obtained new spectra and images of POX 52 at the Keck and Las Campanas observatories.\\footnote{At the time of this writing, the SIMBAD database contains incorrect coordinates for POX 52.  The correct J2000 coordinates are listed in NED as $\\alpha=12^\\mathrm{h}02^\\mathrm{m}56\\fs9$, $\\delta= -20\\arcdeg56\\arcmin03\\arcsec$.}  Throughout this paper we assume a Hubble constant of $H_0 = 70$ km s\\per\\ Mpc\\per. ",
        "conclusions": "New spectra confirm that POX 52 is a genuine Seyfert galaxy, and the broad components seen on the permitted line profiles (first detected by KSB) demonstrate that POX 52 should be classified as a Seyfert 1. The host galaxy morphology is most similar to that of a dwarf elliptical.  This is only the second known example of a Seyfert nucleus in a dwarf galaxy, and the first in a dwarf elliptical.  The black hole mass, estimated from the \\hbeta\\ linewidth-luminosity-mass correlations for Seyferts, is $1.6\\times10^5$ \\msun, and for this value of \\mbh\\ POX 52 falls very close to the extrapolated \\msigma\\ relation of nearby galaxies.  If this black hole mass estimate is accurate, then the AGN in POX 52 is likely to be radiating at nearly its maximal rate, with $\\lbol\\sim(0.5-1)\\times\\ledd$.  With only two known AGNs in dwarf galaxies, little can be said about the statistics or demographics of this class of objects.  The small sample can be at least partly attributed to the fact that no systematic searches for AGNs in dwarf galaxies have been carried out. Additional examples could potentially be detected in objective-prism surveys for emission-line galaxies such as the KISS survey \\citep{sal00}.  Currently, the Sloan Digital Sky Survey presents perhaps the best opportunity to find more members of this apparently rare class of AGNs.  A search of the Sloan data archives to select additional Seyfert nuclei in dwarf galaxy hosts is currently underway \\citep{gh04}, as a first step toward filling out the census of massive black holes with $M < 10^6$ \\msun."
    },
    "0402/astro-ph0402326_arXiv.txt": {
        "abstract": "We have used optical interferometry to obtain multi-wavelength visibility curves for eight red giants over the wavelength range 650--1000 nm. The observations consist of wavelength-dispersed fringes recorded with MAPPIT (Masked APerture-Plane Interference Telescope) at the 3.9-m Anglo-Australian Telescope. We present results for four Miras (R~Car, $o$~Cet, R~Hya, R~Leo) and four semi-regular variables (R~Dor, W~Hya, L$_2$~Pup, $\\gamma$~Cru).  All stars except $\\gamma$~Cru show strong variations of angular size with wavelength.  A uniform-disk model was found to be a poor fit in most cases, with Gaussian (or other more tapered profiles) preferred.  This, together with the fact that most stars showed a systematic increase in apparent size toward the blue and a larger-than-expected linear size, even in the red, all point toward significant scattering by dust in the inner circumstellar environment. Some stars showed evidence for asymmetric brightness profiles, while L$_2$~Pup required a two-component model, indicating an asymmetrical circumstellar dust shell. ",
        "introduction": "Late-type giants have particularly extended atmospheres, which makes it difficult to define a particular value for the stellar diameter \\citep{Baschek91,Scholz01}. Even within the narrower context of observed intensity diameters, their complicated center-to-limb brightness profiles and strong dependence on bandpass from absorption in atmospheric layers leaves no simple ``continuum diameter'' to be measured.  To truly characterise the brightness distribution of a single late-type giant, one would thus like the complete stellar intensity profile at all wavelengths. Since all of these stars pulsate and some show long-term cycle-to-cycle variations in diameter (eg \\citealt{Tuthill95}), as expected from models \\citep{Hofmann98}, this information must be obtained simultaneously and the observations then repeated at many different pulsation phases.  There is a significant history of multi-wavelength angular diameter measurements of red giants, beginning with \\citet{Bonneau73}. Unfortunately, difficulties in calibration and the availability of only a few selected bandpasses has meant that spectral coverage of apparent diameter measurements has been limited. In this paper, we present simultaneous multi-wavelength measurements of the apparent sizes of 8 giants using aperture masking interferometry. With a maximum baseline of 3.89\\,m, our visibility data (related to the object intensity profile by the Fourier transform) were able to resolve all the stars in this sample. ",
        "conclusions": "Simultaneous spatial and spectral data have been recorded for a sample of eight red giants using the novel cross-dispersed interferometry technique developed by the MAPPIT project. With spectral resolving power $\\lambda/\\delta\\lambda< 100$ and spatial resolutions sufficient to measure structures $> 10$\\,mas in size, it has been possible to recover the angular size distribution as a function of wavelength spanning the R and I-bands for our sample stars. A combination of wavelength and baseline bootstrapping enabled coherent averaging of many data sets, enabling a large increase in the signal-to-noise ratio. With the exception of $\\gamma$~Cru, which had a much earlier spectral type (M3.5) than the rest of the sample, diameters were found to exhibit dramatic changes with wavelength. These diameter excursions occurred both as broad trends with wavelength across the entire band, and in narrow spectral windows. In the latter case, the changes were manifest as enlargements across spectral regions associated with strong TiO absorption, although these were mixed with nearby quasi-continuum layers at the spectral resolutions achieved here, arguing for further work with higher dispersion.  For the Mira variables in the sample, Gaussian-like (or other tapered) radial profiles were found to give a better fit to the data than uniform disk profiles. The prevalence of these Gaussian-like profiles, the increase in apparent size toward the blue as separate from molecular effects and the larger than expected apparent sizes even at 920\\,nm all point toward the significance of scattering by dust in the inner circumstellar environment affecting interferometric observations at these wavelengths. Non-centro-symmetric elements were detected for 3 stars, which may be explained as thermal or opacity inhomogeneities in the stellar atmosphere or inner circumstellar regions.  The star L$_{2}$~Pup was found to be something of a special case, with the visibility data betraying the presence of two resolved components, which we interpret as a stellar disk and a dusty circumstellar envelope. Asymmetries detected in this star raise the interesting possibility of resolved highly clumpy structure as close as the dust condensation radius.  These results are all in accord with the theoretical and observational picture of pulsating late-M giants which emphasizes the potential for the extended molecular atmosphere to dramatically affect the observed properties of the stars."
    },
    "0402/hep-ph0402049_arXiv.txt": {
        "abstract": "We combine the most recent observations of large-scale structure (2dF and SDSS galaxy surveys) and cosmic microwave anisotropies (WMAP and ACBAR) to put constraints on flat cosmological models where the number of massive neutrinos and of massless relativistic relics are both left arbitrary.  We discuss the impact of each dataset and of various priors on our bounds. For the standard case of three thermalized neutrinos, we find $\\sum m_{\\nu} < 1.0 \\, ({\\rm resp.} \\, 0.6)$ eV (at 2$\\sigma$), using only CMB and LSS data (resp. including priors from supernovae data and the HST Key Project), a bound that is quite insensitive to the splitting of the total mass between the three species. When the total number of neutrinos or relativistic relics $N_{\\rm eff}$ is left free, the upper bound on $\\sum m_{\\nu}$ (at 2$\\sigma$, including all priors) ranges from $1.0$ to $1.5$ eV depending on the mass splitting. We provide an explanation of the parameter degeneracy that allows larger values of the masses when $N_{\\rm eff}$ increases. Finally, we show that the limit on the total neutrino mass is not significantly modified in the presence of primordial gravitational waves, because current data provide a clear distinction between the corresponding effects. ",
        "introduction": "Neutrino properties are among the most difficult to be probed experimentally, due to the weakness of their interactions.  {}Data from particle accelerators tell us that there are only three flavor neutrinos, while neutrino oscillation experiments show evidence for non-zero neutrino masses (for a recent review, see e.g.\\ \\cite{Maltoni:2004ex}). Recent results strongly constrain the mass differences of the individual neutrino masses (actually masses squared, $\\Delta m^2$) and mixing angles, but no definite conclusion can be drawn neither on the absolute scale of neutrino masses, nor on the existence of weakly coupled sterile neutrinos.  Fortunately, cosmology is quite sensitive to the neutrino sector (see \\cite{Dolgov:2002wy} for a review), and can shed light on these questions, as well as other interesting issues regarding the Early Universe, such as the process of neutrino decoupling from the primordial plasma.  Currently, the most popular cosmological model is the flat adiabatic $\\Lambda$CDM scenario, in which the present density of the Universe is shared between baryons, Cold Dark Matter (CDM) and a cosmological constant $\\Lambda$. This model makes rather simplistic assumptions concerning the neutrino sector, consisting only of three ultra-relativistic neutrinos. It turns out that with a more refined description of the neutrino sector, one finds that only small corrections to the standard picture are allowed after comparing with current data on Cosmic Microwave Background (CMB) anisotropies and Large Scale Structure (LSS).  However, these small corrections carry enough interesting physical implications to justify an active research effort, in particular after the first releases of the WMAP and SDSS data. The results of this effort are not only new cosmological bounds on neutrino properties but also a better understanding of how the errors depend on (i) the experimental CMB and LSS data, (ii) external priors on the cosmological parameters, (iii) intrinsic parameter degeneracies in the theory of cosmological perturbations, (iv) assumptions concerning the underlying cosmological model and parameter space.  We here perform a new analysis using the most recent LSS (2dF, SDSS) and CMB (WMAP, ACBAR) data and an extended cosmological model with an arbitrary number of massive neutrinos and additional relativistic particles, parametrized via an effective number of neutrinos ($N_{\\rm eff}$). We extend the recent work of \\cite{Barger:2003vs} and those that appeared after the release of WMAP data \\cite{Spergel:2003cb}-\\cite{Hannestad:2003ye}. In particular, our underlying model is identical to that of ref.\\ \\cite{Hannestad:2003ye}, but our analysis differs since we include an extended set of data (such as the SDSS results and a more updated version of the 2dF ones), a new prior on the matter density from SN-Ia \\cite{Tonry:2003zg} and non-linear corrections to the LSS power spectrum on scales $0.1 \\, h~{\\rm Mpc}^{-1} < k < 0.2 \\, h~{\\rm Mpc}^{-1}$.  Furthermore, we increase the number of free parameters to ten, in order to analyze the bounds in the presence of primordial tensor perturbations.  The rest of the paper is organized as follows. After a short summary of the effects of neutrino masses and additional relativistic particles in Sec.\\ \\ref{effects}, we describe our analysis method and dataset in Sec.\\ \\ref{method_data}. We discuss our results and compare with previous works in Sec.\\ \\ref{results}. Finally, we conclude in Sec.\\ \\ref{conc}. ",
        "conclusions": "\\label{conc}   We have calculated cosmological bounds on neutrino masses and relativistic relics ($N_{\\rm eff}$) using the latest data on CMB (WMAP and ACBAR) and LSS (SDSS and 2dF galaxy surveys) in the framework of an extended flat $\\Lambda$CDM. In the cases in which a comparison is possible, our results are in good agreement with those of previous analyses \\cite{Barger:2003vs}-\\cite{Hannestad:2003ye}. In the well motivated case of three flavor neutrino with degenerate masses, we found an upper limit on the total masses of $M < 1.0 ~({\\rm resp.}~ 0.6)$ eV using only CMB and LSS data and priors (resp. including priors on $h$ and $\\Omega_{\\Lambda}$).  The bound for four thermalized neutrinos with only one of them carrying a significant mass is $M < 0.8-1.2$ eV, depending on the priors used. Therefore, the 4-neutrino solution to the LSND results is not completely ruled out, but some tension with cosmological data exists, especially if the strong SN03 prior is taken into account.  In the case of arbitrary $N_{\\rm eff}$, our results are summarized in Fig.\\ \\ref{fig1abcd} and listed in Table \\ref{table1}. They clearly show the existence of a parameter degeneracy between the total neutrino mass and $N_{\\rm eff}$, a trend already observed in previous works \\cite{Lesgourgues:2001he,Elgaroy:2003yh,Hannestad:2003xv,Hannestad:2003ye} that we have explained in Sec.\\ \\ref{results}. External priors on $h$ and $\\Omega_{\\Lambda}$ are found to be of particular importance for constraining respectively $N_{\\rm eff}$ and $M$.  Since the standard $\\Lambda$CDM model (with its three effectively massless thermal relics) sits within the 1$\\sigma$ preferred region, we find no evidence for exotic physics such as out-of-equilibrium neutrino decoupling, non-standard nucleosynthesis, extra relativistic relics, a significant amount of hot dark matter, etc. However, large deviations from the standard case are still compatible with observations: for instance, a model with one neutrino of mass $M=1.5$ eV and eight relativistic degrees of freedom is allowed by CMB and LSS data, even when all priors are included (see Fig.\\ \\ref{fig_split}). In order to exclude this model, it is necessary to take into account the prediction of standard BBN, which gives stronger limits on $N_{\\rm eff}$.  The bounds obtained in this paper are based on the observation of cosmological perturbations (CMB and LSS), combined with constraints on the current expansion and acceleration rates of the Universe (HST and SN priors). Therefore they are completely independent from the predictions of primordial nucleosynthesis.  It is remarkable that in the space of the two standard BBN free parameters $(\\omega_b, N_{\\rm eff})$, the preferred regions deduced from cosmological perturbations and from primordial abundances are perfectly compatible with each other, and more or less orthogonal: indeed, our analysis (with the most restrictive priors) gives $0.0215 < \\omega_b < 0.0235$ and $1.6 < N_{\\rm eff} < 8.5$, while standard BBN favors $0.017 < \\omega_b < 0.026$ and $1.6 < N_{\\rm eff} < 3.6$ \\cite{Cuoco:2003cu} (all these bounds are at the 2$\\sigma$ level).  In order to test their robustness, we have also calculated the bounds on $M$ and $N_{\\rm eff}$ in the presence of primordial tensor perturbations. Our results show that the bounds are practically unchanged, because current cosmological data is able to distinguish between the respective effects of tensors and neutrinos.  Finally, we have considered the impact of a different splitting of the total neutrino mass among the individual states, an analysis also recently performed in ref.\\ \\cite{Hannestad:2003ye}. We compared the case of complete mass degeneracy (all neutrinos with the same mass) with that where one neutrino state effectively possesses the whole mass. We found that the bounds on the degenerate case are more restrictive due to its more efficient free-streaming, in particular for larger values if $N_{\\rm eff}$. However, for three or four neutrinos the differences between the two cases are not significant.  Our bounds are a clear indication that present cosmological data provide interesting bounds on the neutrino sector, complementary to those from terrestrial experiments. These include tritium beta decay experiments, which provide a current upper bound on the total neutrino mass of $6.6$ eV at 2$\\sigma$ \\cite{Bonn:tw}, while the KATRIN experiment \\cite{Osipowicz:2001sq} is planned to have an accuracy of the order $0.35$ eV. Sub-eV sensitivity to neutrino masses is also expected for experiments measuring neutrinoless double beta decays \\cite{Elliott:2002xe}, but only for Majorana neutrinos and with a dependence on the details of the mixing matrix.  However, the cosmological bounds should be taken with care, due to their dependence on the data (or priors) used, and also on the assumption of a particular underlying model. Examples are given by the works \\cite{Blanchard:2003du,Allen:2003pt}, where non-zero neutrino masses are preferred. This warning should not prevent us to be confident on the power of future cosmological experiments to limit (and eventually detect) neutrino masses and other neutrino properties. For instance, forecast analyses have shown that with future data there will be potential sensitivities to $\\Delta N_{\\rm eff}\\sim 0.2$ \\cite{Bowen:2001in,Bashinsky:2003tk} (eventually improving BBN results) and neutrino masses of the order $0.1-0.2$ eV with Planck and final SDSS data \\cite{Eisenstein:1998hr,Lesgourgues:1999ej,Hannestad:2002cn}, or with galaxy and CMB lensing \\cite{Kaplinghat:2003bh,Abazajian:2002ck}."
    },
    "0402/astro-ph0402657_arXiv.txt": {
        "abstract": "We study the strong gravitational lensing properties of galaxy clusters obtained from N-body simulations with different kind of Dark Energy (DE). We consider both dynamical DE, due to a scalar field self--interacting through Ratra--Peebles (RP) or SUGRA potentials, and DE with constant negative $w=p/\\rho= -1$ (\\LCDM). We have 12 high resolution lensing systems for each cosmological model with a mass greater than $5.0 \\times 10^{14}$\\Msunh. Using a Ray Shooting technique we make a detailed analysis of the lensing properties of these clusters with particular attention to the number of arcs and their properties (magnification, length and width). We found that the number of giant arcs produced by galaxy clusters changes in a considerable way from \\LCDM~ models to Dynamical Dark Energy models with a RP or SUGRA potentials. These differences originate from the different epochs of cluster formation and from the non-linearity of the strong lensing effect. We suggest the Strong lensing is one of the best tool to discriminate among different kind of Dark Energy. ",
        "introduction": "The mounting observational evidence for the existence of Dark Energy (DE), which probably accounts for $\\sim 70\\%$ of the critical density of the Universe \\cite{Perlmutter,Riess,Tegmark01,Netterfield,Pogosian03,Efstathiou2,Percival,spergel2003}, rises a number of questions concerning galaxy formation. The nature of DE is suitably described by the parameter $w=p/\\rho$, which characterizes its equation of state. The \\LCDM~model ($w=-1$) was extensively studied during the last decade. Recently much more attention was given to physically motivated models with variable $w$ \\cite{Mainini03a}, for which a number of N--body simulations have been performed (Klypin et al 2003, KMMB03 hereafter, Dolag et al. 2003, Linder \\& Jenkins 2003, Macci\\`o et al. 2004). One of the main results of KMMB03 was that dynamical DE halos are denser than those with the standard \\LCDM~ one. In this work we want to analyze the impact of this higher concentration on the strong lensing properties of the cluster size halos.  Was first noted by Bartelmann et al. (1998) (for OCDM, SCDM and \\LCDM~ cosmology) that the predicted number of giant arcs varies by orders of magnitude among different cosmological models. The agreement between data and \\LCDM~ simulation was tested by many authors (see Meneghetti et al 2000, Dalal et al. 2003, Wambsganss et al. 2004) but the situation is still unclear. A direct comparison of arcs statistic with observational data is out of the scope of this work, what we want to point out is the capability of Strong Lensing to discriminate between different kinds of Dark Energy (a similar paper but for a different choice of the dynamical DE parameters was recently submitted by Meneghetti et al. (2004)).  Here, using a Ray Shooting technique, we make a lensing analysis of dark matter halos extracted from N-body simulations of cosmological models with varying $w$ arising from physically motivated potentials which admit tracker solutions. In particular, we focus on the two most popular variants of dynamical DE \\cite{wett1,RP,wett2}. The first model was proposed by Ratra \\& Peebles (1984, RP hereafter) and it yields a rather slow evolution of $w$. The second model \\cite{BraxMartin99,BraxMartinRiazuelo,BraxMartin00} is based on potentials found in supergravity (SUGRA) and it results in a much faster evolving $w$. Hence, RP and SUGRA potentials cover a large spectrum of evolving $w$. These potentials are written as \\begin{eqnarray} V(\\phi) &=& \\frac{\\Lambda^{4+\\alpha}} {\\phi^\\alpha} \\qquad RP, \\\\ V(\\phi) &=& \\frac{\\Lambda^{4+\\alpha}}{\\phi^\\alpha} \\exp (4\\pi G \\phi^2)~~~ SUGRA. \\end{eqnarray} Here $\\Lambda$ is an energy scale, currently set in the range $10^2$--$10^{10}\\, $GeV, relevant for the physics of fundamental interactions. The potentials depend also on the exponent $\\alpha$. The parameters $\\Lambda$ and $\\alpha$ define the DE density parameter $\\Omega_{DE}$. However, we prefer to use $\\Lambda$ and $\\Omega_{DE}$ as independent parameters.  Figure~10 in Mainini et al. (2003b) gives examples of $w$ evolution for RP and SUGRA models.  The SUGRA model considered in this paper has $\\Lambda=10^3$ GeV this implies $w=-0.85$ at $z=0$, but $w$ drastically changes with redshift: $w\\approx -0.4$ at $z=5$. The first RP model (RP$_1$) has the same value for $\\Lambda$ of the SUGRA model. At redshift $z=0$ it has $w=-0.5$; then value of $w$ gradually changes with the redshift: at $z=5$ it is close to $-0.4$. Although the $w$ interval spanned by this RP model covers values significantly above -0.8 (not favored by observations), this case is still important both as a limiting reference case and to emphasize that models with constant $w$ and models with variable $w$ produce different results even if average values of $w$ are not so different. For the second RP model (RP$_2$) we have chosen $\\Lambda=10^{-8}$ GeV, in this case the value of the state parameter at redshift $z=0$ is the same of SUGRA: $w(z=0,\\Lambda=10^{-8}$ GeV$)=-0.85$. This model is certainly better in agreement with CMB constrains but it loses most of its interest from a theoretical point of view: such a small value of $\\Lambda$ has not any clear connection with the physics of fundamental interactions and so it has exactly the same ``fine tuning'' problem of the \\LCDM~ model.  We have normalized all the models in order to have today the same value of the $rms$ density fluctuation on a scale of 8 \\Mpch~, that has been chosen as $\\sigma_8=0.8$. ",
        "conclusions": "Models with dynamical DE are in an infant state. We do not know the nature of DE. Thus, the  state parameter $w(t)$ is still uncertain. In view of this functional indetermination, at first sight, it could seem that the situation is hopeless.  In spite of that, we can outline some general trends that result from our analysis: in dynamical DE models, halos tend to collapse earlier than in a \\LCDM~ model with the same normalization at $z=0$. As the result, halos are more concentrated and denser in their inner parts (KMMB03). Starting from this finding we have explored the consequences of this higher concentration, on strong lensing properties of dark matter halos, in SUGRA and RP cosmologies.   We found that RP$_1$ halos (obtained assuming the cluster abundance of the power spectrum and a value for the energy scale $\\Lambda$ in the range suggested by the physics of fundamental interactions) produce a higher number of arcs with a $L/W>10$ if compared to the standard \\LCDM~ model. This model (RP$_1$) is marginal consistent with observations and its purpose is mainly to illustrate the principal effect of a dynamical dark energy component on arcs statistic.  The second model we analyzed based on RP potential (RP$_2$) is more realistic from an observational point of view but less motivated by theoretical arguments. This model produce about 50\\% more arcs with $L/W>10$ than the \\LCDM~ one for $z_l=0.3$ and $z_s=1$ but it is marginally distinguishable from \\LCDM~ for lensing system at moderate high redshift ($z_l>0.35$, fig \\ref{fig:arcsz}) or for high redshift sources/arcs ($z_l=0.3$ and $z_s=2$).   The SUGRA model is always in between the \\LCDM~ and the RP$_1$ models and it produces about 70-80\\% more arcs than \\LCDM~. This difference is almost constant both changing the sources and the lens redshift and it tends to disappear for $z_l>0.6$ (for $z_s=1.0$) where all the lensing systems considered in this paper ($M_l \\approx 5 \\times 10^{14}$ ) are unable to produce highly distorted images (except in the test model RP$_1$). We also noted that part of the stronger lensing signal due to the higher concentration of halos in dynamical DE models is partially canceled by geometrical effects that increase the critical surface density in such models (fig. \\ref{fig:sigma} and fig. \\ref{fig:arcs2}).    As final remark we would like to stress that arc statistic is a powerful tool to investigate the nature of the Dark Energy. The forthcoming observational surveys(i.e. CFHT Legacy Survey, SDSS and others) will improve the statistic of giant arcs on the sky (for example the RCS-2 Survey (Gladders et al. 2003) will cover an area of 830 deg$^2$ and is expected to produce 50-100 new arcs). Such an observational material will provide a discrimination between DE cosmologies possibly allowing to constrain the $\\Lambda$ scale of the SUGRA and RP potentials."
    },
    "0402/astro-ph0402182_arXiv.txt": {
        "abstract": "We are conducting deep searches for radio pulsations at L-band ($\\sim 20$\\,cm) towards more than 30 globular clusters (GCs) using the 305\\,m Arecibo telescope in Puerto Rico and the 100\\,m Green Bank Telescope in West Virginia. With roughly three quarters of our search data analyzed, we have discovered 12 new millisecond pulsars (MSPs), 11 of which are in binary systems, and at least three of which eclipse.  We have timing solutions for several of these systems. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0402/astro-ph0402461_arXiv.txt": {
        "abstract": "We address whether possible scale-dependent deviations from gaussianity in the primordial density field that are consistent with the cosmic microwave background  observations  could explain the apparent excess of early cluster formation at high redshift.  Using two phenomenological non-gaussian models we find that at fixed normalisation to the observed local abundance of massive clusters, the protoclusters observed at z$\\sim$4 are significantly more likely to develop in strongly non-gaussian models than in the gaussian paradigm. We compute the relative $z<1$ evolution of X-ray cluster counts in the non-gaussian case with respect to the gaussian expectation, and the relative excess contribution to the CMB power spectrum due to the integrated thermal SZ effect.  We find that both the observed hints of an  unexpectedly slow evolution in the X-ray counts and the excess power at high $\\ell$ that may have been observed by CMB interferometers  can also be reproduced in our non-gaussian simulations. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0402/astro-ph0402527_arXiv.txt": {
        "abstract": "We present a family of spherical models for elliptical galaxies and bulges consisting of a stellar component and a central black hole. All models in this family share the same stellar density profile, which has a steep central cusp. The gravitational potential of each models is a linear combination of the potential generated selfconsistently by the stars and the potential of a central black hole. The relative importance of these two contributions is a free parameter in the models. Assuming an isotropic dynamical structure, almost all kinematical properties of these models can be calculated analytically. In particular, they form the first simple dynamical models for galaxies with a central black hole where the distribution function and differential energy distribution can be written completely in terms of elementary functions only. We also present various extensions of this family to models with anisotropic orbital structures. Also for these dynamical models, the distribution function and its moments can be expressed completely in terms of elementary functions.  This family is useful for a large range of applications, in particular to generate initial conditions for N-body and hydrodynamical simulations to model galactic nuclei with a central black hole. ",
        "introduction": "During the past few years, various numerical dynamical modelling techniques have been developed to construct accurate dynamical models for galaxies (Dejonghe~1989; Emsellem, Monet \\& Bacon~1994; Rix et al.~1997; van der Marel et al.~1998; Gerhard et al.~1998; Cretton et al.~1999; Gebhardt et al.~2000a; Verolme \\& de Zeeuw~2002). The state-of-the-art techniques can handle various degrees of complexity such as the capability to include non-trivial geometries and higher-order kinematical information. Nevertheless, analytical models remain interesting and important for various reasons. They can provide a simple testbed where various physical processes can be investigated, or where new modelling or data analysis techniques can be explored. Analytical models are particularly useful to generate the initial conditions for N-body, SPH or Monte Carlo simulations. Most attention has been devoted to the construction of spherical selfconsistent models, i.e.\\ models in which the stars move in a potential generated completely by the stars themselves. Famous examples include the Plummer model (Plummer~1911; Dejonghe~1987), the isochrone sphere~(H\\'enon~1959, 1960) and the Jaffe sphere (Jaffe~1983) and the Hernquist model (Hernquist~1990; Baes \\& Dejonghe~2002). Such models often serve as a template model for elliptical galaxies, globular clusters or the bulges of spiral galaxies.  Recent observational evidence indicates, however, that the central regions of these stellar systems cannot always be faithfully represented by selfconsistent models. The existence of supermassive black holes in the nuclei of galaxies has been suspected for a long time, as accretion onto massive compact objects was regarded as the only reasonable explanation for the existence of active galaxies and quasars (Salpeter~1964; So{\\l}tan~1982). During the last decade, quiescent supermassive black holes have been detected in the centre of the Milky Way (Ghez et al.~2000; Sch\\\"odel et al.~2002) and in the nuclei of virtually all external galaxies which are nearby enough to resolve the sphere of influence of the black hole (e.g.~see list in Tremaine et al.~2002). The masses of these black holes are roughly between a million and a few billion solar masses and are tightly coupled to structural parameters of the host galaxies (Kormendy \\& Richstone~1995; Ferrarese \\& Merritt~2000; Gebhardt et al.~2000b; Graham et al.~2001; McLure \\& Dunlop~2002; Ferrarese~2002; Baes et al.~2003; Marconi \\& Hunt~2003). Recently, evidence for the presence of intermediate mass black holes in the centre of globular clusters has been reported (Gebhardt, Rich \\& Ho~2002; Gerssen et al.~2002, 2003), although this evidence is still controversial (Baumgardt et al.~2003a,b; McNamara, Harrison \\& Anderson~2003). These observations clearly indicate that there is a need for simple but realistic analytical dynamical models with a central black hole.  The first requirement for the construction of models with a black hole is a potential-density pair with a cuspy central density profile. Tremaine et al.~(1994) demonstrate that spherical isotropic systems need a central density profile that diverges at least as $r^{-1/2}$ to be able to support a central black hole. The need for such cuspy models also follows from observations: HST photometry has shown that ellipticals and the bulges of spiral galaxies generally have central density cusps that diverge as $r^{-\\gamma}$ at small radii with $0\\leq\\gamma\\leq2.5$ (Lauer et al.~1995; Gebhardt et al.~1996; Ravindranath et al.~2002; Seigar et al.~2002). Also for the Milky Way there is evidence for a central density cusp in the innermost regions (Genzel et al.~2003). A very useful one-parameter family of spherical models with this characteristic, which we refer to as the $\\gamma$-models, was introduced independently by Dehnen~(1993) and Tremaine et al.~(1994). This family has a simple analytical density profile which diverges as $r^{-\\gamma}$ in the central regions $(0\\leq\\gamma<3)$, and includes the Hernquist and Jaffe models as special cases. Zhao~(1996) generalized this family further to a very general three-parameter family of models, the so-called $(\\alpha,\\beta,\\gamma)$-models. For many of these models, interesting dynamical properties such as the intrinsic and projected velocity dispersions, the distribution function and the differential energy distribution can be calculated analytically if one assumes an isotropic selfconsistent dynamical structure.  The second step in the construction of dynamical models with a cuspy core is to add to the potential of the selfconsistent model an extra contribution from the black hole, and re-calculate the dynamical properties with this new potential. The calculation of the (intrinsic and/or projected) velocity dispersions in the presence of a black hole is not very difficult, and can usually be performed analytically for those models where the dispersions can be calculated analytically in the selfconsistent case. For example, for the sets of potential-density pairs considered by Tremaine et al.~(1994) and Zhao~(1996), the addition of a black hole was not a problem for the calculation of the velocity dispersion profile. The reason is that the intrinsic and projected velocity dispersions are just linear functions of the potential, and therefore linear functions of the black hole mass. Many other interesting kinematical properties, in particular the phase space distribution function, depend on the potential in a strongly non-linear way, however. The construction of dynamical models in which these more complicated kinematical quantities can be expressed analytically in the presence of a black hole proves to be more difficult. Apart from the asymptotic behaviour, these characteristics usually need to be calculated numerically (Tremaine et al.~1994; Dejonghe et al.~1999).  The work by Ciotti~(1996) provides a first attempt to construct completely analytical dynamical models for galaxies with a central black hole. In an effort to construct realistic dynamical models for elliptical galaxies embedded in massive cuspy dark matter haloes, he considers a set of two-component Hernquist models. They consist of a stellar component and a halo component, both of which are modelled as a Hernquist profile, but with different masses and core radii. The interesting aspect of his work for our goal is that the masses and core radii of each component can be taken arbitrarily. Setting the core radius of the halo component to zero, his two-component model degenerates into a Hernquist model with a central black hole. Ciotti demonstrates that the distribution function and differential energy distribution of Hernquist models with a central black hole can be expressed analytically, albeit as rather cumbersome expressions. For example, the distribution function can be written as the derivative of a complicated function involving incomplete elliptic integrals and Jacobian functions. Nevertheless, his work presents the first model for galaxies with a central black hole where most of the kinematical properties can be calculated analytically.  In this paper, we present a detailed kinematical analysis of a family of spherical models for elliptical galaxies and bulges consisting of a stellar component and a central black hole. In Section~2 we define our family of models. All models in this one-parameter family share the same stellar density profile, which corresponds to one of $\\gamma$-models introduced by Dehnen~(1993) and Tremaine et al.~(1994). The potential of the models is a linear combination of the potential generated by the stars and the potential of a central black hole. The importance of both components to the total potential can be set arbitrarily by varying the parameter $\\mu$, which represents the mass of the central black hole relative to the total mass. The reason why we chose this particular family of models is that most of the interesting kinematical properties can be expressed completely analytically for all values of $\\mu$. In Section~3 we describe some intrinsic kinematical properties of this family of models, such as the distribution function, the differential energy distribution and the moments of the distribution function. In particular, we give closed analytical expressions for these quantities, completely in terms of elementary functions. In Section~4 we describe some observed kinematical properties as they are projected on the plane of the sky. Most of these quantities can be expressed analytically involving simple incomplete elliptic integrals. In Section~5 we present a number of generalizations to models with a anisotropic orbital structure. In particular, we discuss models with a constant anisotropy profile, models with a distribution function of the Osipkov-Merritt type and models with a completely radial orbital structure. Also for these models, the distribution function and its moments can be calculated completely analytically for all values of the central black hole mass. Finally, a discussion is given in Section~5. ",
        "conclusions": "We have described a one-parameter family of spherical models for elliptical galaxies and bulges consisting of a stellar component and a central black hole. All models in this family share the same stellar density profile, which has a steep cusp with a slope of $-\\tfrac{5}{2}$ in the centre. The gravitational potential for the models is a linear combination of the potential generated selfconsistently by the stars and the potential of a central supermassive black hole. The parameter $\\mu$, representing the mass of the black hole relative to the total mass of the galaxy, can assume any value between 0 and 1. The family therefore contains models ranging from a selfconsistent model without black hole to a model where the potential is entirely dominated by the black hole. We have done an extensive study of the internal and projected kinematics for this family of models. With the assumption of isotropy, we have calculated the intrinsic velocity dispersions, the distribution function, the differential energy distribution, the moments of the distribution function, the surface brightness and the projected velocity dispersions. All of these quantities have been expressed completely analytically, for all values of $\\mu$.  We have also described some extensions of the models to anisotropic orbital structures. In particular, we have considered models with a constant anisotropy, models with distribution functions of the Osipkov-Merritt type and models with a completely radial orbital structure. Also for these families of models, the distribution function and its moments can be calculated completely analytically for all values of the central black hole mass $\\mu$.  We are well aware that the stellar component we have considered does not completely represent the detailed structure of the centre of galaxies. Firstly, it has a fairly steep density cusp, which is at the edge of the observed range in real elliptical galaxies (Lauer et al.~1995; Gebhardt et al.~1996). With such a steep cusp, much of the stellar mass is very centrally concentrated, resulting in an infinitely deep stellar potential well. The addition of a black hole potential to this stellar potential does not drastically change the global mass distribution, such that the effects of a black hole on the kinematical properties are probably quite conservative. For example, the distribution function is substantially, but not very drastically affected by the presence of a black hole, whereas this effect is much stronger for $\\gamma$-models with a less steep density cusp (see figure~6 in Tremaine et al.~1994). The reason why we focused on the density profile~(\\ref{density}), and not on one of the other members of the family of $\\gamma$-models with a shallower central density cusp for example, is that the $\\gamma=\\tfrac{5}{2}$ model is unique in the way that it allows a relatively simple expression of the augmented density in the presence of a central black hole. It is the only member of the family of $\\gamma$-models where the distribution function, the differential energy distribution and the moments can completely be expressed in terms of elementary functions. We are currently undertaking a more general study of the properties of the $\\gamma$-models with a central black hole, using both analytical and numerical means (Baes, Buyle \\& Dejonghe~2004).  Secondly, the models presented here are spherically symmetric and isotropic, whereas few elliptical galaxies are thought to be perfectly spherical. Observational studies suggest that a substantial fraction of elliptical galaxies are at least moderately triaxial (Franx, van Gorkum \\& de Zeeuw~1991; Tremblay \\& Merritt~1995; Bak \\& Statler~2000). The central regions of galaxies with a central black hole are generally believed to be roughly axisymmetric, however. Indeed, supermassive black holes drive the shape of galaxies in a triaxial haloes toward axisymmetry by stochastic diffusion, either globally (Gerhard~\\& Binney~1985; Norman, May \\& van Albada~1985; Merritt \\& Quinlan~1997; Wachlin \\& Ferratt-Mello~1998; Valluri \\& Merritt~1998), or at least the central regions (Holley-Bockelmann et al.~2001, 2002). Still, axisymmetric systems have a much larger freedom in orbital structure than spherical systems. Using detailed axisymmetric dynamical modelling, Gebhardt et al.~(2003) found that the central regions of the massive early-type galaxies generally have a significant tangential anisotropy, whereas less massive galaxies have a range of anisotropies.  In spite of these critical notes, the family of dynamical models presented in this paper has the huge advantage over more complicated numerical models that all kinematical properties can be calculated completely in terms of elementary functions, for any value of the parameter $\\mu$. This is not only the case for rather simple kinematical properties such as the velocity dispersions, but also for more complicated kinematical properties which depend in a strongly non-linear way on the potential. In particular, the distribution function and the differential energy distribution can be written in a compact form and in terms of elementary functions only. As a result, this family of models is useful for a large set of applications. In particular, it can be used to easily generate the initial conditions for N-body or hydrodynamical simulations, which are needed to investigate how black holes interact with the stellar, gaseous and dark matter components of their host galaxies."
    },
    "0402/astro-ph0402241_arXiv.txt": {
        "abstract": "We examine the ability of photoevaporative disk winds to explain the low-velocity components observed in the forbidden line spectra of low-mass T Tauri stars. Using the analytic model of Shu, Johnstone, \\& Hollenbach (1993) and Hollenbach et al. (1994) as a basis, we examine the characteristics of photoevaporative outflows with hydrodynamic simulations.  General results from the simulations agree well with the analytic predictions, although some small differences are present. Most importantly, the flow of material from the disk surface develops at smaller radii than in the analytic approximations and the flow-velocity from the disk surface is only one-third the sound speed.  A detailed presentation of observational consequences of the model is given, including predicted line widths, blue-shifts, and integrated luminosities of observable sulfur and nitrogen emission lines.  We demonstrate that these predictions are in agreement with current observational data on the low-velocity forbidden line emission of ionized species from T Tauri stars.  This is in contrast with magnetic wind models, which systematically under-predict these forbidden line luminosities.  However, the present model cannot easily account for the luminosities of neutral oxygen lines in T Tauri stars. ",
        "introduction": "It is now generally accepted that most young low-mass stars are born with circumstellar disks and that the evolution of these disks holds a key to planet formation.  Observational evidence for circumstellar disks takes many forms, from direct imaging (e.g., Jayawardhana et al. 2002) and imaging of disk silhouettes (e.g., Bally, O'Dell, \\& McCaughrean 2000) through indirect measurements such as infrared excess (e.g., Lada et al. 2000). These data provide an estimate for the lifetime of disks around low-mass stars, $\\tau_{\\rm disk} \\sim 6 \\times 10^6\\,$yr (Haisch, Lada, \\& Lada 2001), requiring the presence of an efficient disk dispersal mechanism.  Hollenbach, Yorke, \\& Johnstone (2000) considered a variety of disk dispersal mechanisms and concluded that viscous accretion of the disk onto the central star (e.g. Hartmann et al. 1998), although dominating erosion of the inner disk, is incapable of removing the entire disk mass in the required time. Alternatively, photoevaporation of the disk, while providing an effective mechanism for disk removal at large radii $\\gtrsim  10\\,$AU, cannot remove the inner material [Shu, Johnstone, \\& Hollenbach 1993 (hereafter SJH93); Hollenbach et al. 1994 (hereafter, HJLS94); Johnstone, Hollenbach, \\& Bally 1998].  Acting together viscous accretion and photoevaporation were predicted to efficiently remove the entire disk. Numerical calculations by Clarke, Gendrin, \\& Sotomayor (2001) and Matsuyama, Johnstone, \\& Hartmann (2003a) have shown that the combined effects of photoevaporation and viscous accretion are capable of dispersing disks on $10^5 - 10^7\\,$yr timescales.  Additionally, the formation of gaps within the gaseous disk during the dispersal era may place constraints on the evolution of planetary orbits (Matsuyama, Johnstone, \\& Murray 2003b).  The photoevaporation model for disk dispersal has been shown to fit the observational data well in the case of external heating via nearby massive stars (Bally et al. 1998; Johnstone, Hollenbach, \\& Bally 1998; St\\\"orzer \\& Hollenbach 1998).  However, with the exception of a few cases (e.g., MWC349A), evidence for disk photoevaporation due to heating from the central star (SJH93, HJLS94) is largely circumstantial.  Low-velocity $\\sim 10\\,$km$\\,$s$^{-1}$, blue-shifted forbidden lines observed in the spectra of T Tauri stars [Hartigan, Edwards, \\& Ghandour 1995  (hereafter HEG95)] may provide evidence for the $10^4\\,$K thermal disk winds expected in the photoevaporation model. The theoretical models of SJH93 and HJLS94 do not, however, include detailed hydrodynamic calculations and thus it has not been possible to test the model directly against these observations.  In this paper we take a first look at the hydrodynamic flow of thermally-driven disk winds powered by photoevaporative heating in order to determine the observational consequences of the photoevaporation model. We begin in \\S 2 with a brief description of how photoevaporation works, concentrating on photoevaporation due to heating from a star at the center of the disk. In \\S 3 we describe how the hydrodynamical simulations are performed and their results.  The predicted observable properties of the simulations are presented in \\S 4 and are compared with existing data.  Further discussion and conclusions are presented in \\S 5. ",
        "conclusions": "The evidence for self-produced photoevaporative disk winds around low-mass stars is largely circumstantial.  The primary reason for this has been the lack of detailed calculations on the outflow properties of these winds to compare with observations.  Photoevaporative disk wind models have, however, been shown to fit the observational data of externally-heated low-mass stars in the vicinity of massive stars (e.g., Johnstone et al. 1998; Bally et al. 1998; St\\\"orzer \\& Hollenbach 1998) because the disk wind geometry and flow characteristics are much simpler.  In this paper we use hydrodynamic simulations to compute the properties of the photoevaporative disk wind model and, subsequently, the predicted observational signatures.  We show that the hydrodynamic simulations converge to steady-state solutions and that a number of general properties (e.g., scale height distribution, integrated disk wind mass-loss rate) agree well with the earlier analytic results of SJH93 and HJLS94.  The outflows become nearly spherically symmetric and match the analytic solution to the Parker solar wind problem at large radii, as expected.  We further demonstrate that the photoevaporative winds produce blue-shifted ionized forbidden lines with typical offset velocities of $\\sim -10$ km s$^{-1}$, widths of $\\sim 30$ km s$^{-1}$, and integrated luminosities of $10^{-7}-10^{-4} L_{\\odot}$.  These values compare favorably with the observations of HEG95 (see Figs. 9 and 10).  The strength of the observed OI emission is difficult to account for with the present model, however.  The photoevaporative disk winds examined here are intuitive and physically-simplistic.  They rely on the availability of $\\Phi_* = 10^{40-42}$ ionizing photons per second to be emitted by the central star disk accretion shock.  Theoretical arguments support such high levels of ionizing emission (e.g., Matsuyama et al. 2003a).  Observationally, it is extremely difficult to directly pin down the value of $\\Phi_*$ for low-mass T Tauri stars.  HEG95 measured the mass accretion rates onto the surface of the T Tauri stars, however, and these, in principle, can be used to estimate the value of $\\Phi_*$. Given the mass accretion rate and the stellar radius, we calculate the accretion shock luminosity of the gas as it falls onto the star's surface via  \\begin{equation} L_{\\rm acc} = \\frac{G M_* \\dot{M}_{\\rm acc}}{2 R_*} \\end{equation}  \\noindent where $\\dot{M}_{acc}$ is the mass accretion rate and $R_*$ is the radius of the T Tauri star (see also Matsuyama et al. 2003a).  Using the inferred values of $L_{acc}$ and assuming black-body radiation with a characteristic temperature of 10,000 K (Johns-Krull et al. 2000; Gullbring et al. 2000), we indeed estimate $\\Phi_* \\sim 10^{40-42}$ s$^{-1}$.  This lends credence to the photoevaporative wind hypothesis.   Utilizing these ultraviolet photons, however, may be difficult. As mentioned above, both the protostellar jet (Shang et al. 2002) and the accretion column (Alexander et al. 2004) are likely to be optically thick to ionizing radiation, protecting the disk.  Nevertheless, energy released within the accretion disk, and in the fast-moving jet, may produce significant ionizing radiation. Alternatively, the disk may be heated by another source such as FUV photons (Johnstone et al. 1998, St\\\"orzer \\& Hollenbach 1998).  The calculations presented here reveal that heating the disk surface to $10^4\\,$K will produce a thermal wind with flow characteristics similar to those observed around T Tauri stars. Additionally, if the mass-loss rate is high enough, comparable to the evaporating disk model, the forbidden line profiles of SII and NII can be explained. Curiously, if the $10^4\\,$K wind is not ionized by EUV photons the resultant OI forbidden line radiation also would be comparable to that observed. Thus, the observations of HEG95 seem to require a partially ionized $10^4\\,$K wind.\"  Alternative models have also been proposed to explain the observed low-velocity component.  In fact, the prevailing view is that T Tauri winds are magnetically-driven (e.g., Anderson et al. 2003).  This is largely because strong magnetic fields seem to be the only way to power the collimated, high velocity jets that are responsible for the high-velocity component observed in T Tauri spectra.  In addition, magnetic wind models are able to reproduce a number of the observed features of the high-velocity component (e.g., Shang, Shu, \\& Glassgold 1998; Garcia et al. 2001a, 2001b).  It is less clear, however, if the magnetic wind models can account for the low-velocity component, which was the subject of investigation in \\S 4.  Below, we briefly discuss how the predicted low-velocity components of magnetic wind models compare with their observational counterparts.  We argue that the photoevaporative disk wind model proposed in the present study matches the current suite of observations at least as well magnetic wind models.  Generally speaking, magnetic wind models may be divided into two classes: `X-winds' and `disk winds'.  The primary difference between these two classes is the physical extent of the magnetic field and, therefore, the location of the outflow launching radius.  In X-wind models, the magnetic field is contained entirely within the very inner regions of the disk immediately next to the star. Disk wind models, on the other hand, assume that there is a large scale magnetic field present which extends to several AU.  Unfortunately, it is extremely difficult to work out the exact observational predictions of such models in a self-consistent manner (Garcia et al. 2001a, 2001b).  This is one advantage the physically-simplistic photoevaporative disk wind model has over the magnetic wind models.  While it is difficult to calculate the precise observational characteristics of magnetic wind models, there have been a number of basic predictions made which are testable with current observational data.  In terms of the X-wind models, the predicted launching radius for the disk outflow is extremely close to the star.  However, recent high resolution observations suggest that the low-velocity component probably originates at `large' disk radii (e.g., Bacciotti et al. 2002; Anderson et al. 2003), suggesting that some other mechanism is responsible for the observed low-velocity component.  In fact, irrespective of the model, one should expect the outflow velocity to be of order the escape velocity at the region where the wind originated.  This is a strong argument that the low-velocity component must originate at moderate radii ($\\sim 1-10$ AU).  Magnetic `disk wind' models, on the other hand, are able to produce both high and low-velocity components (e.g., Cabrit et al. 1999; Garcia et al. 2001a, 2001b), with the low-velocity component being launched from reasonably large disk radii.  Furthermore, such models predict that the high velocity component is spatially extended while the low-velocity component is confined within the central (projected) regions of the system, which is in good agreement with observations.  However, the problem with `disk wind' models is that the level of ionization at large disk radii (i.e., at the launching radius) is too low to allow for strong magnetic coupling (e.g., Gammie 1996).  Hence, such models tend to under-predict the luminosity of the low-velocity component by a substantial margin (e.g., Garcia et al. 2001b).  As we have shown, however, the PDW model presented here (which contains no magnetic field) can naturally produce an outflow that matches a number of observed trends seen in low-mass T Tauri stars, so long as the photons can penetrate from the central star out to large radii (or if an external mechanism for producing EUV photons is present).   It is possible, however, that {\\it both} photoevaporation and a large scale magnetic field could be contributing to the low-velocity component, especially since the higher level of ionization due to photoevaporation should increase the effectiveness of the magnetic disk wind.  This two component model could possibly explain the relatively large amount of scatter in the forbidden line profiles and luminosities and why some stars apparently show multiple low-velocity components (e.g., UY Aur, see Fig. 9).  Additionally, it is likely that an interface between the fast jet-like wind arising from the inner disk will modify these results.  HJLS94 modeled the effect of a strong stellar wind on the ionized atmosphere and concluded that only the inner regions would be affected. This suggests that the results of the calculations in this paper would need to be modified only in the central densest regions, resulting in somewhat lower line fluxes. This would predominantly affect transitions having high critical densities such as the nitrogen line. Moreover, Matsuyama et al. (2003a) found that gaps could form in the inner part of the disk, near $r_g$, due to the erosive power of photoevaporation coupled to the viscous evolution of the disk.  Although the Matsuyama models should be recalculated, taking into account the results of this paper, these gaps would remove the inner densest part of the ionized atmosphere produced in this paper. An evolutionary sequence might be visible in the low-velocity spectra from T Tauri stars dependent upon the state of the underlying disk. Detailed calculations, taking into account each of these factors, should be undertaken.  \\vskip0.1in \\noindent The authors wish to thank the referee, David Hollenbach, for suggesting significant improvements to the paper. Thanks also to G. Mellema for providing his software package for calculating forbidden line emissivities.  I.G.M. thanks A. Babul for useful discussions.  A.F. acknowledges financial support from J.F. Navarro. I.G.M. is supported by a postgraduate scholarship from the Natural Sciences and Engineering Research Council of Canada (NSERC). The research of D.J. has been supported by an NSERC Discovery Grant. D.R.B. thanks NSERC for financial support."
    },
    "0402/astro-ph0402288_arXiv.txt": {
        "abstract": "{We present evolutionary models of helium accreting carbon-oxygen white dwarfs in which we include the effects of the spin-up of the accreting star induced by angular momentum accretion, rotationally induced chemical mixing and rotational energy dissipation. Initial masses of 0.6 \\Msun{} and 0.8 \\Msun{} and constant accretion rates of a few times $10^{-8}{\\rm M_{\\odot}/yr}$ of helium rich matter have been  considered, which is typical for the sub-Chandrasekhar mass progenitor scenario for Type Ia supernovae. It is found that the helium envelope in an accreting white dwarf is heated efficiently by friction in the differentially rotating spun-up layers. As a result, helium ignites much earlier and under much less degenerate conditions compared to the corresponding non-rotating case. Consequently, a helium detonation may be avoided, which questions the sub-Chandrasekhar mass progenitor scenario for Type Ia supernovae. We discuss implications of our results for the evolution of helium star plus white dwarf binary systems as possible progenitors of recurrent helium novae. ",
        "introduction": "Type Ia supernovae are of key importance for the chemical evolution of galaxies, as they are a major producer of iron group elements (e.g. Nomoto et al.~\\cite{Nomoto84}; Renzini~\\cite{Renzini99} ). They were also found to be excellent distance indicator and have become an indispensable tool in cosmology (Phillips~\\cite{Phillips93}; Hamuy et al.~\\cite{Hamuy96}; Branch~\\cite{Branch98}). The recent suggestion of a non-zero cosmological constant is partly based on observations of SNe Ia at high redshift (Leibundgut~\\cite{Leibundgut01}). Given that distance determinations at high redshift through SNe~Ia depend on the assumption of the homogeneity of SNe Ia light curves throughout the ages, an understanding of the possible diversity of their progenitors is crucial to evaluate this approach. Nevertheless, the progenitors of Type Ia supernovae have not been identified yet, and the debate on their exact nature continues (e.g., Livio~\\cite{Livio01}).  One possibility to obtain a SN~Ia is the detonation of the degenerate helium layer accumulated on top of a CO white dwarf due to mass transfer from its low mass helium star companion in a close binary system, which triggers a carbon detonation in the white dwarf core. This is the so called double detonation or sub-Chandrasekhar mass scenario for SNe~Ia, as it may allow to explode white dwarfs with masses well below the Chandrasekhar mass (e.g. Nomoto~\\cite{Nomoto82b}; Fujimoto~\\cite{Fujimoto82c}; Limongi \\& Tornamb\\'e~\\cite{Limongi91}; Livne~\\cite{Livne90}; Woosley \\& Weaver~\\cite{Woosley94}; Livne \\& Arnett~\\cite{Livne95}). While the capability of the helium detonation to ignite the CO core is still debated (e.g., Livio~\\cite{Livio01}), the helium detonation by itself would produce an explosion of supernova scale.  Currently, the sub-Chandrasekhar mass scenario is not favored as a major source of SNe Ia mainly because the light curves and spectra obtained from this model are not in good agreement with observations (e.g. H\\\"oflich \\& Khokhlov~\\cite{Hoeflich96}; Nugent et al.~\\cite{Nugent97}). Especially, the predicted presence of high velocity Ni and He is most stringently criticized (e.g. Hillebrandt \\& Niemeyer~\\cite{Hillebrandt00}; Livio~\\cite{Livio01}).   On the other hand, stellar and binary evolution theory predicts a realization frequency of binary systems such as helium star cataclysmics --- which might produce double detonating sub-Chandrasekhar mass white dwarfs --- which amounts to a few times $10^{-3}~{\\rm yr^{-1}}$ per galaxy (e.g. Iben \\& Tutukov~\\cite{Iben91}; Reg\\\"os et al.~\\cite{Regos02}) which is comparable to the expected total SN~Ia rate in the Milky Way. This raises the question why such explosions are practically never observed.  We note that the sub-Chandrasekhar mass SN progenitor models which have been constructed so far neglected the effects of rotation, which can be one of the primary factors determining the evolution of stars, in particular of massive stars (Langer~\\cite{Langer98};  Maeder \\& Meynet~\\cite{Maeder00}). Iben \\& Tutukov (\\cite{Iben91}) pointed out that rotation may indeed be important in helium star cataclysmic systems. Yoon \\& Langer (\\cite{Yoon02}, \\cite{Yoon04a}) and Yoon et al. (\\cite{Yoon04d}) showed that effects of rotation might be essential for the evolution of accreting white dwarfs when the accreted matter contains a high specific angular momentum. The induced spin-up was found to change the white dwarf structure significantly and to produce rotationally induced chemical mixing.  In this paper, we suggest that rotation could play a key role in helium accreting white dwarfs such that in model which would produce a helium detonation this phenomenon is completely avoided when the white dwarf spin-up is considered. After explaining the numerical method and physical assumptions of the present study in Sect.~\\ref{sect:method}, we investigate the evolution of helium accreting carbon-oxygen white dwarfs with accretion rates of $\\sim 10^{-8}$~\\msyr{}, with the effects of rotation considered, in Sect.~\\ref{sect:results}. Implications of our results for helium novae and neutron capture nucleosynthesis are discussed in Sect.~\\ref{sect:discussion}. Our main conclusions are summarized in Sect.~\\ref{sect:conclusion} ",
        "conclusions": "\\label{sect:conclusion}   We have shown that the effects of rotation in helium accreting white dwarfs may be incompatible with the scenario of double detonations in sub-Chandrasekhar mass CO white dwarfs as possible SNe Ia progenitors. In helium accreting white dwarfs, we find the thermal evolution to be affected by viscous heating due to differential rotation in the spun-up layers, such that helium ignition is induced at too low densities to develop a detonation. This may give a plausible solution to the long standing problem of the missing observational counterparts of sub-Chandrasekhar explosions, which are predicted to occur with a frequency comparable to the observed SN Ia rate.  We discussed that binary systems consisting of a CO white dwarf and a less massive helium star may be possible progenitors of recurrent helium novae (Iben \\& Tutukov~\\cite{Iben91}), which may be analogous to V445 Puppis (Ashok \\& Banerjee~\\cite{Ashok03}; Kato \\& Hachisu~\\cite{Kato03}). After the first strong helium nova flash in such binary systems, rather mild nova outbursts are expected to occur recurrently with a period of $\\sim 10^5$ yr. The realization frequency of such a helium nova may be as high as $\\sim 0.1~{\\rm yr^{-1}}$ in our Galaxy.  Rotation induces chemical mixing of \\Ne{22} and \\He{4} at the bottom of the helium envelope, which may provide interesting conditions for neutron capture nucleosynthesis to occur during the helium nova flashes.  Finally, we note that other important mechanisms for the angular momentum redistribution in white dwarfs may exist than those considered in the present study. In particular, we neglected the possible role of magnetic fields, which may increase the efficiency of the angular momentum transport significantly (cf. Heger et al.~\\cite{Heger03}; Maeder \\& Meynet~\\cite{Maeder03}). If the spin-up time scale is shorter than considered in the present study, the resulting shear strength will be weaker and the effect of rotational energy dissipation may not be as important as shown here. However, studies of magnetic effects in accreting white dwarfs have to be left for future investigations."
    },
    "0402/astro-ph0402594_arXiv.txt": {
        "abstract": "{  In  a  previous paper we  reported  on a discontinuity in  the extreme horizontal  branch (EHB)  of the  Galactic   globular cluster NGC6752, which we called the second $U$-jump.  This feature was attributed to a combination of post zero-age horizontal branch evolution and diffusion effects.  In this follow-up study  we  analyze other EHB clusters  and show that the  second $U$-jump is a  common feature among EHB clusters reaching  $T_{\\rm eff}\\ge  23,000$K, and that   its onset in different clusters converges around $T_{\\rm eff}\\sim 21,000\\pm3,000$K. We also present  near-ultraviolet diagrams of $\\omega$Cen and NGC2808, the only two   objects with spectroscopically confirmed  ``blue hook'' stars ($T_{\\rm eff}\\ge 35,000$K). We confirm predictions of a photometric discontinuity separating late  from   early-helium flashers.  Moreover,   we present empirical evidence   that  the second  $U$-jump   population might be mainly composed by early-helium flashers. Lastly, we revisit the discussion on the ubiquitous nature of the gaps and jumps  so  far  identified in  the blue   HB tails, suggesting   a possible discrete nature of the  distribution in temperature of the HB stars. ",
        "introduction": "Over the last decades, both observational and theoretical efforts have been devoted  to the analysis   of the observed  distribution of stars along the Horizontal Branch (HB) of  Galactic globular clusters.  This notwithstanding,  our understanding of  several observational features is still  incomplete.  In fact,  even if theory and observations agree that    the   HB morphology    is   governed  by    metallicity   (the first-parameter),  since late sixties it  has  become clearer that the color distribution of HB stars in Galactic globular  clusters is not a unique function of metallicity (Sandage \\& Wildey \\cite{sand67}). NGC362 and NGC288 are a classical example of how two clusters  sharing similar  metallicities can show   remarkably different  HB morphology.  Hence,  other parameters (e.g. age, cluster environment,    Helium abundance,  mass    loss and rotation)  besides metallicity   affect  the evolution  of   HB   stars (second-parameter debate).    \\begin{table*}[t] \\begin{center} \\caption{Ground-based observations log along with the HST archival data} \\begin{tabular}{llllll} \\hline\\hline \\noalign{\\smallskip} Object & Instrument & Prog. ID & Date & Filters & Seeing \\\\ \\hline NGC1904 & WFI@2.2   & 64.L-0255 &  1999 & $U,B,V$   & 0\\farcs8--1\\farcs2 \\\\ NGC6752 & WFI@2.2   & 65.L-0561 &  2000 & $U,B,V$   & 0\\farcs6--0\\farcs9 \\\\ NGC7099 & WFI@2.2   & 65.L-0561 &  2000 & $U,B,V$   & 0\\farcs6--0\\farcs9 \\\\ NGC6273 & WFI@2.2   & 65.L-0561 &  2000 & $U,B,V$   & 0\\farcs7--1\\farcs1 \\\\ NGC7089 & WFI@2.2   & 69.D-0582 &  2002 & $U,B,V,I$ & 0\\farcs7--1\\farcs3 \\\\ NGC5139 & WFI@2.2   & 69.D-0582 &  2002 & $U,B,V,I$ & 0\\farcs6--1\\farcs5 \\\\ NGC5986 & SUSI2@NTT & 71.D-0175 &  2003 & $U,V$     & 0\\farcs9--1\\farcs4 \\\\ NGC6656 & SUSI2@NTT & 71.D-0175 &  2003 & $U,V$     & 0\\farcs9--1\\farcs4 \\\\ NGC6715 & SUSI2@NTT & 71.D-0175 &  2003 & $U,V$     & 0\\farcs9--1\\farcs4 \\\\ \\hline NGC2808                & WFPC2@HST & GO8655 & 2001 & $F450W$        &  \\\\ % NGC2808                & WFPC2@HST & GO6804 & 1998 & $F336W$ &  \\\\ % NGC6093\t\t      & WFPC2@HST  & GO8655 & 2000 & $F450W$        & \\\\ NGC6093\t\t      & WFPC2@HST  & GO6460 & 1997 & $F336W$        & \\\\ NGC6205\t\t      & WFPC2@HST  & GO5903 & 1996 & $F336W$,$F450W$,$F555W$ & \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\label{t_log} \\end{center} \\end{table*}  The second parameter  effect is not the only  puzzling feature  in the evolution  of HB stars.  In particular,  photometric  studies of stars hotter than the RR Lyrae instability strip showed the presence of: (a) gaps along the blue tail (Ferraro et al.\\ \\cite{ferr98}; Piotto et al. \\cite{piot99};  (b)  a  jump  around  $T_{\\rm eff}\\sim11,500$K in the Str\\\"omgren $u$, $u-y$ (Grundahl et al.\\  \\cite{grun99}, hereafter G99) and Johnson $U$, $U-V$ (Bedin et al.\\ \\cite{bedin00}) color-magnitude diagrams (CMDs); (c) hot HB stars reaching temperatures of $T_{\\rm eff}\\simeq30,000$K or more in metal-poor (D'Cruz et al \\cite{cruz96}, Brown et al.\\  \\cite{brow01}) and metal-rich (Rich et al.\\ \\cite{rich97})   clusters;  and  (d) the   still   unexplained presence of fast HB rotators (Behr et al.\\ \\cite{behr00}; Recio-Blanco et  al.\\  \\cite{ale02}).   On the other   hand,  spectroscopic studies showed  the presence  of  abundance (Behr et  al.\\ \\cite{behr99}), and gravity anomalies (Moehler et al.\\ \\cite{moeh00}) in stars hotter than $T_{\\rm eff}\\sim11,500$K.  Horizontal branch stars  hotter than   $T_{\\rm eff}\\sim20,000$ K   are usually referred to  as extreme HB  (EHB) stars.  It is  believed that EHBs experience high mass-loss  during their red giant phase, reducing their H-rich envelope down to  $\\le 0.05$M$_{\\odot}$, to the point  of being  unable to  sustain  H-shell burning.   However,  it  is hard to explain  why such an enhanced  mass  loss occurs  along the red  giant branch. Near-UV  CMDs of NGC6752 (Momany  et al.\\ \\cite{moma02}) have revealed another  interesting feature  along  the EHB.    In the  $U$ {\\it vs.} $(U-V)$  plane,  the   HB showed  a discontinuity   at $U-V\\simeq-1.0$ (corresponding $T_{\\rm  eff}\\sim23,000$  K).  Given  the  (1) apparent photometric similarities with  the  ``cooler'' G99 jump, and   (2) the clear difference in temperature with respect  to blue hook stars (i.e. blue hook stars are  generally hotter than $T_{\\rm  eff}\\sim35,000$ K) we called    this feature  the ``second-$U$ jump'',    and tentatively attributed it to  a combination of post  ZAHB evolution  and diffusion effects.  Most of  these puzzling  features  remain unsolved.  In particular, we lack of a global  view on the  origin and  internal properties of  EHB stars. This is not a problem confined to the final stages of evolution of globular cluster   stars.  Indeed, the  nature of  EHBs has a  more general relevance in astrophysics as  these are considered responsible of  the  UV excess observed    in the spectra  of elliptical  galaxies ($UV$-upturn galaxies, Greggio \\& Renzini \\cite{gregg90}).  In  this paper we  present new near-UV  CMDs  for a selected sample of Galactic globular  clusters, characterized by  an HB with  an extended blue tail, with the aim to investigate the observational properties of the EHB stars in these clusters.  This  paper is organized as follows: in the following section we  discuss  the observational data-base  and briefly  outline the  main  reduction  and calibration  procedures; in Section~\\ref{s_diagrams} we show   that the  second $U$-jump,  already identified   in  NGC6752, is also  present  in  other EHB clusters; in Section~\\ref{s_bh}   we  suggest a  link  between  the second $U$-jump feature and the He flash induced mixing scenario discussed by Brown et al.\\ (\\cite{brow01}).  A summary will close the paper.  \\begin{table*}[t] \\begin{center} \\caption{Properties of cluster sample. Columns 2 to  5 are from the Harris on-line-catalog:    http://physun.physics.mcmaster.ca/~harris/mwgc.dat (\\cite{harris96}) as updated  on February  2003.  Columns 7  and 8 are from Rosenberg et al.\\ (\\cite{rosen99})} \\begin{tabular}{llllllclll} \\hline\\hline \\noalign{\\smallskip} Object    & $E_{B-V}$ &  $(m-M)_V$  & [Fe/H] &   c$^{\\mathrm{1}}$  & M$_V$ & $\\Delta V_{TO}^{HB}$ & Normalized relative age & second $U$-jump & blue hook\\\\ \\hline NGC6752   & 0.04    & 13.13 & $-$1.56  & c                     &  -7.73   & 3.55 & 1.03 &     y  & n  \\\\ NGC6656   & 0.34    & 13.60 & $-$1.64  & 1.31                  &  -8.50   & 3.55 & 1.04 &     y  & n  \\\\ NGC5139   & 0.12    & 13.97 & $-$1.62$^{\\mathrm{2}}$ & 1.61    &  -10.29  & ---  & ---  &     y  & y  \\\\ NGC6205   & 0.02    & 14.48 & $-$1.54  & 1.51                  &  -8.70   & 3.55 & 1.02 &   y & y$^{\\mathrm{3}}$ \\\\ NGC7099   & 0.03    & 14.62 & $-$2.12  & c                     &  -7.43   & ---  & ---  &      n  & n  \\\\ NGC288    & 0.03    & 14.83 & $-$1.24  &  0.96                 &-6.74   & 3.55 & 0.97 & n & n \\\\ NGC7089   & 0.06    & 15.49 & $-$1.62  & 1.80                  &   -9.02  & ---  & ---  &     y  & n  \\\\ NGC6093   & 0.18    & 15.56 & $-$1.75  & 1.95                  &   -8.23  & 3.55 & 1.04 &     y  & n  \\\\ NGC2808   & 0.22    & 15.59 & $-$1.15  & 1.77                  &   -9.39  & 3.30 & 0.81 &     n  & y  \\\\ NGC1904   & 0.01    & 15.59 & $-$1.57  & 1.72                  &   -7.86  & 3.50 & 1.00 &     n  & n  \\\\ NGC6273   & 0.41$^{\\mathrm{4}}$   & 15.95 & $-$1.68  & 1.53    &   -9.18  & ---  & ---  &      y  & n  \\\\ NGC5986   & 0.28    & 15.96 & $-$1.58  & 1.22                  &   -8.44  & ---  & ---  &      y  & n  \\\\ NGC6715   & 0.15    & 17.61 & $-$1.58$^{\\mathrm{2}}$ & 1.84                  &   -10.01  & ---  & --- &      y  & y$^{\\mathrm{3}}$  \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\label{t_clusters} \\end{center} \\begin{list}{}{} \\item[$^{\\mathrm{1}}$] {Central concentration, c = log(r$_t$/r$_c$); a ``c'' denotes a core-collapsed cluster} \\item[$^{\\mathrm{2}}$] {Refers to the metallicity of the dominant population} \\item[$^{\\mathrm{3}}$] {We present evidence of the presence of blue hook stars} \\item[$^{\\mathrm{4}}$] {A cluster suffering severe differential reddening}  \\end{list}  \\end{table*} ",
        "conclusions": "\\label{s_summary}  We have presented near-UV diagrams for a  dozen of clusters.  Our main result is that the   previously reported discontinuity in NGC6752 (around $T_{\\rm eff}\\simeq 23,000$K)  is also present in other extreme HB  clusters.  The  onset  of  this second  $U$-jump   in the examined clusters seems also to coincide  at a  temperature of $\\sim  21,000\\pm 3,000$K.   Both  these facts strengthen  the idea that the second $U$-jump is an indicator of a physical process, acting in  all   extreme HB clusters,  that  has yet  to be   fully understood.  We have clearly shown that the second $U$-jump population is photometrically well disentangled from the  hotter blue hook stars, and brought  evidence that the  two, chemically different, populations can co-exist in objects like $\\omega$Cen and M54.  In the flash mixing scenario, early helium flashers are expected to pile  up at the end of the extreme  HB, and this is  exactly what  the second jump population seems to   show.  Hence, we suggest that    the second jump population contains a significant number of early helium flashers.  \\subsection{Final remarks and future work}  One of  the main difficulties in  studying the HB star distribution is providing  a   satisfactory,   quantitative,    description  of    its morphology. To this aim, many HB morphology parameters were introduced (Fusi Pecci  et  al.\\ \\cite{fusi93},  Buonanno et  al.\\ \\cite{buon97}, Catelan et    al.\\ \\cite{cata98}, Piotto   et  al.\\  \\cite{piot99} and references  therein) to describe the    distribution  of HB stars   in color/temperature, measure  their maximum extension, and  reveal peaks and gaps.  It remains however, that the  majority of these studies was based on $BV$ CMDs, clearly not an ideal plane for detailed studies of the HB (Ferraro et al.\\ \\cite{ferr98}).  In light of  the growing number of  $UV$ CMDs,  and recent photometric findings  (the  G99  and second   $U$-jumps, and blue  hook stars)  we revisit the  HB  morphology in  the $UV$ plane.   {\\em  In  brief, the analysis of  the HBs in 7 clusters  lead us to suggest  that the HB in $UV$ CMDs can be envisaged as the sum of  discrete segments}.  This is not a  new idea  (Buonanno  et al.\\  \\cite{buon85}).  However  the new observations  in  the $UV$  plane and the  larger  sample  of HB stars unveil a number of discrete branches in the HB previously undisclosed. To better explain  this fact we rely  on Fig.~\\ref{f_twins}, showing a combination of $UV$ CMDs.  A more  detailed and complete analysis will be subject  of   a future  work.   Figure~\\ref{f_twins} allows  us  to suggest the following:  \\begin{enumerate} \\item The discontinuities (gaps or jumps) at $T_{\\rm eff}\\sim10,000$K and $11,500$K  occur  in {\\em all}    blue HB   clusters,  in particular, the G99   jump {\\em coincides} with  the  endpoint of many blue HB clusters and dwarf galaxies; \\item In addition, there are two other {\\em main} discontinuities at $T_{\\rm eff}\\sim16,000$K and  $21,000$K,  marking the endpoints of HB reaching these temperature. \\end{enumerate}   Point (1) basically confirms the  already known ubiquitous nature of the G99 jump,  and  suggests  a similar ubiquity  of the gap at $T_{\\rm eff}\\sim10,000$K ({\\it the gap at $B-V$  about zero} in the Caloi \\cite{caloi99} terminology). Similarly, Ferraro et al.\\  (\\cite{ferr98}) also suggested that the gaps at $T_{\\rm eff}\\sim10,000$K  and $11,500$K (G0 and G1 in their terminology)  occur in  {\\em  many but not all  clusters}. Hence, the novelty in point (1) lies in suggesting the ubiquity of the gap at $T_{\\rm eff}\\sim10,000$K, and this mainly relies on panels (B), (C)  and  (D) of Fig.~\\ref{f_twins}.    The three panels propose three ways in which  clusters populate the blue  HB (i.e., the  HB extending from the bluest boundary of the  RR instability strip  to the onset of the G99 jump) around the $T_{\\rm eff}\\sim10,000$K gap.  The {\\it first case}  (panel A overplotting NGC7099 on NGC6752) is the most frequent, that is, a more or  less uniform distribution on the two parts surrounding the gap at $T_{\\rm eff}\\sim10,000$K. Relying  on the   optical CMDs of   the  HST snapshot (Piotto   et al. \\cite{piot02}), and   our $UV$ diagrams, we   count 34 clusters  with a metallicity range between  [Fe/H]$=-1.27$ (NGC5904) and [Fe/H]$=-2.29$ (NGC5053)  which uniformly  populate {\\em  only}  this part of  the HB (i.e.  having no red HB   stars and no  HB  stars hotter  than the G99 jump).   Most  of these clusters   (e.g.   NGC7099) show  a clear gap at $T_{\\rm eff}\\sim10,000$K. On the other hand, we note that Local Group dwarf spheroidals are also representative of a uniformly populated blue HB.  The diagrams of Ursa Minor (Carrera et al.\\ \\cite{car02}) and Sculptor (Hurley-Keller et al. \\cite{hur99}) are perfect examples of blue HB  extending only in this temperature  range,   possibly  showing   the   gap at    $T_{\\rm eff}\\sim10,000$K gap.  The  {\\it second and  third  case} (panels B  and  C) is when clusters preferentially populate one side  of the $T_{\\rm eff}\\sim10,000$K gap;  either  the hotter or   the cooler side.  NGC7078\\footnote{Note that the group of stars distributed on  the right side  of the blue HB clump are probably variables stars (see Zheleznyak \\& Kravtosov \\cite{zhel03}) caught at random phase.} and NGC5466 (see  CMD in Buonanno et  al.\\ \\cite{buon85})  are  examples of blue HB extending  from the RR    instability strip and  stopping  at  $T_{\\rm eff}\\sim10,000$K; i.e. populating only the right side  of the gap. On  the  other hand   NGC288 (panel  C)   shows an  opposite behavior; populating  the  part between  $T_{\\rm eff}\\sim10,000$K and $11,500$K; i.e.  populating the left side  of  the $T_{\\rm eff}\\sim10,000$K  gap. Overall, the 3 proposed distribution modalities might explain previous difficulties in ascertaining the ubiquity of  the gap at $\\sim10,000$K (see discussions in Ferraro et al.\\ (1998) on  the {\\em varying width} of the $T_{\\rm eff}\\sim10,000$K gap from cluster to cluster).   We did not find any cluster with HB extending beyond the G99 jump which do not possess blue HB stars with  $T_{\\rm eff}\\le11,500$K.  In  other words, all the clusters with hot HB stars  ($T_{\\rm eff}\\ge11,500$K) do have stars in the region between the RR Lyrae instability strip and the G99 jump.  We do not have an explanation for such a complex behavior in this part of the HB.    However, the presented   evidence implies that    the discontinuities  at  $T_{\\rm  eff}\\sim10,000$K  and $11,500$K are  presumably present in  {\\em all} clusters.   Whereas the G99 jump can    be   easily discerned   in    UV  CMDS,   the   gap at  $T_{\\rm eff}\\sim10,000$K might be  hampered by photometric  errors and post-HB evolution. Moreover, given the high frequency of globular clusters and dwarf galaxies populating only this specific range of  the blue HB (RR instability  strip---onset  of the  G99 jump),  this group  of objects might    represent the ``standard'' HB   morphology  in the metal-poor regime.  As discussed   in Buonanno et  al.\\ (\\cite{buon85}),  one way around the second parameter debate can be to isolate certain groups of clusters (with   a substantial  similarity   in some   of their  basic properties) and  then explore the effects of  any other difference the clusters in the group  have.  Clusters/dwarf galaxies showing only the blue  HB (between the RR instability  strip and  the  onset of the G99 jump) are most probably a separate group,  within which one can search for ''similarities''.  {\\it What is so special in the group of clusters (e.g. NGC7099) ending their  blue HB  exactly   at the  onset of the    G99 jump at  $T_{\\rm eff}\\sim11,500$}K ?  The occurrence of  the G99 jump was explained  as the  aftermath of  radiative   levitation  that causes  a  substantial increase  in  the metal content   of  the outermost layers.  Radiative levitation  is  possible   after the  disappearance    of the envelope convective layers  located across the  H and HeI ionization regions at $T_{\\rm eff}\\sim10,000$K and $11,000$K respectively (Caloi \\cite{caloi99}, Sweigart  \\cite{sweigart00}).  Hence, both the  gap at $T_{\\rm eff}\\sim10,000$K and  the G99  jump  at $\\sim 11,500$K can  be attributed to atmospheric effects. If we adopt  this explanation for  the two  discontinuities, and given their omni-presence in  different environments, then systems with only blue HBs ending just before the  onset of the G99 jump  can be seen as clusters in which a ``standard'' mass  loss mechanism takes place. The net product of this ``standard'' mass loss on the  red giant branch is {\\em always} HB with an envelope massive enough to possess an extended convective region.  Similarly,    for   clusters  with     HB  extending  beyond   $T_{\\rm eff}\\sim11,500$K,   we have identified  two  truncations points in the blue tail:    one at  at  $T_{\\rm  eff}\\sim16,000$K  and  a  second at $\\sim21,000$K.     Panel (D) shows   that NGC288 ends at   $T_{\\rm  eff}\\sim16,000$K.  As shown in the color distribution  of the HB  in NGC6752 (panel A), this temperature corresponds to a marked  decrease in the stellar counts in NGC6752. On the other hand, this is also the  temperature at which the first of  the   NGC2808 gaps occurs   (Bedin  et al.\\  \\cite{bedin00}). Hence,   besides the G99 jump,  $T_{\\rm eff}\\sim16,000$K seems to mark    another endpoint.  Panels    (E) and (F)    indicate another important   endpoint: the onset  of the second  $U$-jump. Obviously one cannot rely on few stars to mark the end  of the  HB blue  tail of NGC1904 and  NGC4833, however  the coincidence with the onset of the second jump is rather tempting. In  this regards,  it is of   great interest to  further investigate a possible  relation between the   second jump population and the  early helium flashers.    Just as a  working  hypothesis,  if  {\\em all} the second jump  stars were to  be early helium  flashers,  it would imply that stars   hotter  than $T_{\\rm  eff}\\sim21,000$K  have  a different physical origin (flash-mixing  scenario?)  with respect  to ``cooler'' HB stars (produced by standard mass-loss mechanism).  This possibility would have  significant   implications on   our understanding of   the complicated second parameter problem.  In conclusion,  the overall picture of  the HB in  $UV$ diagrams seems rather segmented.   The endpoints which define these segments may be acting  like markers, highlighting the  signature of  different physical   processes  working in HB  stars. The  origin of the $T_{\\rm eff}\\sim10,000$K and $\\sim  11,500$K discontinuities seems  to be related to  the   disappearance of the  convective  envelope layers located  across   the H     and     HeI ionization  regions     (Caloi \\cite{caloi99}).   In this paper we  have shown that the discontinuity at $T_{\\rm  eff}\\sim21,000$K can be  related to the presence  of early helium flashers. For the   endpoint at $T_{\\rm eff}\\sim16,000$K we  do not yet have an explanation."
    },
    "0402/astro-ph0402307_arXiv.txt": {
        "abstract": "We present two adjacent jet candidates with a length of $\\sim9\\degr$ each -- 10$\\times$ longer than the largest known jets -- discovered by us on 60 $\\mu$m and 100 $\\mu$m IRAS maps, but not observed at any other wavelength. They are extremely collimated (length-to-width ratios 20--50), curved, knotty, and end in prominent bubbles. Their dust temperatures are 25 $\\pm$ 3 K and 30 $\\pm$ 4 K, respectively. Both harbour faint stars, one having a high proper motion ($0\\farcs23$ yr$^{-1}$) and being very red, suggesting a distance of $\\sim$ 60 pc. At this distance, the combined mass of both jets (assuming a gas-to-dust ratio of 200) totals $\\sim$ 1 M$_{\\sun}$. We suspect that these gigantic ($\\sim9$ pc length) jets have a common origin, due to the decay of a system of evolved stars. They are the first examples of jets radiating in the far IR and might be the closest non-diffuse nebulae to the solar system. ",
        "introduction": "Collimated jets are one of the most fascinating but poorly understood phenomena in astronomy. They represent ubiquitous features and are found in quasars, active galactic nuclei, young stellar objects, symbiotic systems, planetary nebulae and pulsars. Their acceleration and collimation mechanisms might be the same in all the classes of objects (Livio \\cite{Livio}; Price et al. {\\cite{Price}). Although an agreement on the processes that drive all these jets has not yet been achieved, there is mounting evidence that bipolar ejection is powered by accretion and that magnetic fields play a crucial role in accelerating and collimating the gas (Anderson et al. \\cite{Ander}). The exact launching mechanism remains to be identified (Cabrit et al. \\cite{Cabrit}). The main reason is the rather large distance of even the nearest Galactic jets which preclude detailed observations of regions close enough to the star ($<$ several AU) where the jets are generated. The distances ($\\gg100$ pc) also impede measurements of various important features like the cross-section structure of jets, internal knots and other microstructures. In this paper we present preliminary results on a pair of jets which obviously is unique in several respects and appears to be close to the solar system. ",
        "conclusions": "Two gigantic adjacent jet candidates, found in the far infrared, have been discovered. They represent the first examples of jets radiating in this wavelength range. Due to their vicinity on the sky and their comparable length, shape, high degree of collimation and brightness, and a proper motion of one of the candidate ejecting stars which obviously links both jets, we suggest a common origin. The sources responsible for the ejection can hardly be young objects, since the jet pair is far away from any star forming region. The physical size of the jets and the candidate ejecting star of one of the jets with its very high proper motion suggest a close distance, less than 100 pc, i.e. within the LB. No clear conclusion can be drawn on the origin of the jets, but we suggest that they stem from a decayed stellar system, where a massive star was a mass donor and has generated accretion disks around two low-mass stellar members. The jet pair appears to contain cooled-down material, i.e. it is of fossil type. Other fossil jets may exist, but will not be easily detectable because of the ubiquitous presence of interstellar dust.  A wealth of observations will be necessary to unravel details of these enigmatic adjacent jets. Primary goals should be to carry out high-resolution imaging and spectroscopic observations of the candidate ejection sources, and optical spectroscopy and CO observations of several portions of the jets. First and foremost, the derivation of a reliable distance should be undertaken. Apart from understanding its very nature, this pair might be important for future studies of the acceleration and collimation processes in astrophysical jets due to its closeness."
    },
    "0402/astro-ph0402131_arXiv.txt": {
        "abstract": "We present a detailed reanalysis of {\\it Chandra X-ray Observatory} data for the galaxy cluster Abell~4059 and its central radio galaxy, PKS2354--35.  We also present new 1.4\\,GHz and 4.7\\,GHz CnB-array radio data from the {\\it Very Large Array}\\footnote{The VLA is part of the National Radio Astronomy Observatory which is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.}, as well as a short archival WFPC2 image from the {\\it Hubble Space Telescope}.  The presence of a strong interaction between this radio galaxy and the intracluster medium (ICM) was suggested by Huang \\& Sarazin (1998) on the basis of a short observation by the {\\it High Resolution Imager} on {\\it ROSAT}, and confirmed in our preliminary analysis of the {\\it Chandra}/ACIS-S data.  In particular, X-ray imaging clearly shows two cavities within the ICM that are {\\it approximately} aligned with the radio-galaxy axis.  However, using our new radio maps (which are at lower frequencies and better matched to searching for $\\sim 1\\arcmin$ structures than the previous high-quality maps) we fail to find a detailed correspondence between the $\\sim 1\\arcmin$ scale radio-lobes and the ICM cavities.  This suggests that the cavities are ``ghosts'' of a previous burst of powerful activity by PKS~2354-35.  This is supported by detailed, spatially-resolved, X-ray spectroscopy which fails to find any shock-heated ICM, suggesting that the cavities are evolving subsonically.  We also examine the nature of the central asymmetric ridge (or bar) of X-ray emission extending for $\\sim30\\kpc$ south-west (SW) of the cluster center that has been noted in these previous analyzes.  We find the ridge to be denser and cooler than, but probably in pressure balance with, its surroundings.  The thermal evolution of this structure seems to be dominated by radiative cooling, possibly enhanced by the radio-galaxy ICM interaction.  We discuss several possible models for the formation of this SW ridge and find none of them to be entirely satisfactory.  In our preferred model, the SW ridge is due to radiative cooling induced by an interaction between a radio-galaxy driven disturbance and a pre-existing bulk ICM flow. The presence of such a bulk flow (with a velocity of $\\sim 500\\kmps$ projected on the plane of the sky) is suggested by the off-center nature of the pair of X-ray cavities.  Such a bulk flow can be created during a cluster/sub-cluster merger --- the presence of a prominent dust-lane in the cD galaxy of Abell~4059, ESO~349-G010, is circumstantial evidence for just such a merger. ",
        "introduction": "There is an increasing realization that the core regions of clusters of galaxies are complex and dynamic environments.  For some time now, it has been argued on the basis of data from imaging X-ray telescopes that the hot intracluster medium (ICM) of the core regions of rich clusters is radiatively-cooling on timescales shorter than the age of the cluster.  This gives rise to the phenomenon known as a cooling flow.  Fabian (1994) gives an extensive review of cooling flows up to and including constraints from the {\\it ROSAT} observatory.  Prior to the launches of the {\\it Chandra X-ray Observatory} and {\\it XMM-Newton}, the X-ray data strongly argued for the inhomogeneous cooling flow model in which gas in cluster cores cools from X-ray emitting temperatures down to unobservable temperatures as part of a multi-phase ICM over a spatially distributed region of the cluster core.  An obvious mystery, and a strong hint that the real situation is more complex, was the lack of cool gas (including significant star formation) observed at other wavebands.  Not surprisingly, the X-ray view of cluster cores became appreciably more complicated with the launch of {\\it Chandra} and {\\it XMM-Newton}.  With the very high dispersions possible using the reflection grating spectrometer (RGS) on {\\it XMM-Newton}, detailed emission line spectroscopy of cluster cores became possible for the first time.  Using these techniques, observations of Abell~1795 (Tamura et al. 2001) and Abell~1835 (Peterson et al. 2001) both revealed clear evidence for gas cooling from the virial temperature $kT>4\\keV$ down to 1--2\\,keV.  In particular, one could isolate and identify the L-shell emission lines of iron corresponding to gas spanning this temperature range.  However, very tight upper limits were set on the amount of gas below 1--2\\,keV which were in strong disagreement with the standard cooling flow model.  In other words, there is evidence for gas cooling from the virial temperature down to 1--2\\,keV, whence it disappears.  This result has been generalized to a sample of clusters by Peterson et al. (2003).  The explanation for the temperature floor is still far from clear.  Strong (i.e., order of magnitude) metallicity inhomogeneities will skew the apparent cooling function such that gas below 1\\,keV cools extremely rapidly, thereby eluding detection (Fabian et al. 2001).  However, it is not known how such inhomogeneities will be formed or maintained.  Thermal conduction and the action of a central radio galaxy may also be important in producing these temperature floors (Fabian, Voigt \\& Morris 2002; Voigt et al. 2002; Ruszkowski \\& Begelman 2002).  In addition to spectral complexity, {\\it Chandra} has revealed that many clusters possess morphological complexities that are thought to arise due to the interaction of the ICM with a central radio galaxy. In some cases, the association is clear.  For example, Perseus~A (Fabian et al. 2000), Hydra~A (McNamara et al. 2000; David et al. 2001; Nulsen et al. 2002), Abell~2052 (Blanton et al. 2001), and Cygnus~A (Smith et al. 2001) all show well defined cavities in the X-ray emitting gas which are coincident with the current radio lobes of the central radio galaxy.  In these sources, it is clear that the radio lobes have displaced the X-ray emitting gas producing the observed X-ray/radio anti-coincidence.  {\\it Chandra} has also revealed the presence of ``ghost'' cavities, i.e., X-ray cavities that are {\\it not} coincident with the active radio lobes.  Examples include the outer cavity of Perseus~A (Fabian et al. 2000), Abell~2597 (McNamara et al. 2001), NGC~4636 (Jones et al. 2002), and Abell~4059 (Heinz et al. 2002).  In these sources, it is believed that the cavities are associated with old radio lobes (related to previous cycles of AGN activity).  The low-frequency (74\\,MHz) synchrotron radio emission expected within this scenario from these old radio lobes has been observed from the ghost cavity of Perseus-A (Fabian et al. 2002).  Collectively, these observations give rise to several questions.  Most hydrodynamic models for the formation of these cavities (e.g., Clarke et al. 1997; Heinz, Reynolds \\& Begelman 1998; Reynolds, Heinz \\& Begelman 2001; Reynolds, Heinz \\& Begelman 2002, and references therein) involve the pressure-driven growth of a shock-bounded cocoon. However, in almost all cases (with NGC~4636 being a notable exception; Jones et al. 2002), the X-ray shells that bound the observed cavities are cooler than the ambient ICM, seemingly at odds with the shock scenario.  It is plausible that the cool shell arises due to the ``lifting'' of lower-entropy material from the cluster core by the radio galaxy activity (B\\\"ohringer et al. 1995; Reynolds et al. 2001; Nulsen et al. 2002), but we would still expect to see some fraction of the sources in the shock-bounded phase.  More generally, we need to assess the implications of such data for models of radio galaxy evolution.  To achieve this goal requires the detailed analysis of more {\\it Chandra} data together with directed numerical simulations.  Abell~4059 was one of the first clusters known to possess X-ray cavities on the basis of data from the {\\it ROSAT} high-resolution imager (Huang \\& Sarazin 1998).  These cavities were approximately coincident with the radio lobes of the FRI radio galaxy PKS~2354--35, which is hosted by the cD galaxy at the center of the cluster.  Huang \\& Sarazin (1998) also noted an interesting bar-like feature in the central regions of the cluster perpendicular to the radio axis.  The {\\it Chandra} Advanced CCD Imaging Camera (ACIS) observation of this cluster has been previously described by us in Heinz et al. (2002). In that paper, it was shown that the coincidence between the radio lobes and the X-ray cavities is not exact, leading to the conclusion that these are actually ``ghost'' cavities.  It was also suggested that the complex X-ray morphology (including the central bar) arises from an interaction of a radio-galaxy driven expanding cocoon and a pre-existing bulk ICM flow.  Such an ICM flow may result from the accretion of a galaxy group by the cluster.  In this paper, we present a detailed reanalysis of the {\\it Chandra}-ACIS observation of the core regions of Abell~4059.  We present a spatially-resolved spectral study of the core regions of this cluster.  We also present new 1.4\\,GHz and 4.7\\,GHz radio data from the Very Large Array (VLA) taken with the CnB configuration, thereby providing a better match providing a better match to the typical spatial scales characterizing the X-ray cavities.  We confirm that the arcmin scale radio lobes indeed do not coincide precisely with the X-ray cavities, especially to the south-east of the center. We also find that the ridge of emission to the SW of the center is cooler and denser, but probably in pressure equilibrium, with the surrounding ICM.  Furthermore, it is determined that the thermal evolution of this structure must be dominated by radiative cooling. We discuss various models for the SW ridge, but prefer an explanation in which it corresponds to shock/compression induced cooling of ICM caused by interaction of the radio-galaxy driven disturbance with a bulk ICM flow --- however, such a model may suffer fine tuning problems.  Finally, we also present an archival {\\it Hubble Space Telescope} (HST) Wide Field Planetary Camera 2 (WFPC2) image of the cD galaxy ESO~349--G010.  The presence of a significant dust lane in this elliptical galaxy suggests that it has accreted a gas rich companion galaxy within the past $\\sim 10^8$\\,yrs.  This provides further circumstantial evidence for the putative cluster/group merger required to produce the bulk ICM flow.  Section~2 details the {\\it Chandra}, VLA and HST data reduction, and Section~3 describes our imaging and spectroscopy investigations.  The observational results are summarized, and possible models discussed, in Section~4.  Throughout this paper we assume $H_{0}=65$ km s$^{-1}$ Mpc$^{-1}$ and $q_{0}=0.5$.  Given a redshift of $z=0.049$, this cosmology places PKS~2354--35 and Abell~4059 at a luminosity distance of 226\\,Mpc. ",
        "conclusions": "\\subsection{Summary of observational results}  There is clear evidence of a vigorous and complex radio-galaxy/cluster interaction between PKS2354--35 and A4059.  Prior to the analysis presented in this paper, the known facts relevant to this interaction were: \\begin{enumerate} \\item There are two large ICM cavities approximately aligned with the axis of the radio galaxy.  Huang \\& Sarazin (1998) and, later Heinz et al. (2002), showed that the radio source, as defined in the A- and B- array 4.8\\,GHz and 8.5\\,GHz VLA observations of Taylor et al. (1994) extends into the NW cavity, but does not extend to (or even point at) the SE cavity. \\item There is an offset between the center of the axis connecting the two cavities and the galactic nucleus.  One is given the impression that, assuming the cavities were created symmetrically by the radio galaxy, they have subsequently ``drifted'' in a NE direction. \\item There is a bright ridge of emission extending from the center of the cluster in the SW direction.  This ridge terminates about 25\\,kpc to the SW of the center in an abrupt edge. \\end{enumerate} To this, we can now add the following informations: \\begin{enumerate} \\item VLA/CnB-array data taken at 1.4\\,GHz, which is much better matched to detecting arcmin-scale structures than the previous radio data, still fails to detect any radio emission associated with the SE X-ray cavity. \\item There is no indication that the gas around the X-ray cavities is any hotter or higher entropy than the ambient gas.  In other words, there is no evidence for a strong (or even moderately weak) shock surrounding the X-ray cavities. \\item The SW ridge appears to be in approximate pressure balance with the ambient material and is X-ray bright because of its lower temperature and higher density.  The radiative cooling time in this structure is much shorter than that of the surrounding ICM, becoming as short as 100\\,Myr (compared with a general ``core'' cooling time of greater than 500\\,Myr). \\item There is a robust metallicity gradient within the cluster, with high metallicity (approaching solar) in the cluster center and then declining by a factor of 2 beyond 50\\,kpc.  This is reproduced in both the annular (i.e. projected) and deprojected spectral study. The presence of a central depression in the metallicity profile is suggested by single temperature fits to either the projected or deprojected spectra.  However, the reality of this feature is unclear (see above for details). \\item HST/WFPC-2 imaging reveals that the cD galaxy and host of PKS2354--35, ESO349--G010, displays a prominent 5\\,kpc dust lane oriented roughly perpendicular to the radio-axis.  This suggests that it has accreted a dust rich companion galaxy in the past $10^8\\yr$ or so. \\end{enumerate} In this section, we discuss the constraints that these observations place on the nature of the interaction.  \\subsection{Inflating the cavities}  As discussed in Heinz et al. (2002), the current radio source is likely too weak to produce notable cavities, and it is likely that the observed ICM cavities are ``ghosts'' of a previous and more powerful period of activity.  In this picture, the cavities are in a passive phase of evolution (see Reynolds, Heinz \\& Begelman 2002).  The X-ray cavities, which were created by a past phase of supersonic lobe expansion, have decelerated to sub-sonic velocities. Any shocks once bounding the lobes have weakened into mere compression waves.  The fact that this activity produces an expanding shell of ICM implies that gas from the core regions will be lifted to higher points in the cluster, thereby adiabatically cooling as it de-pressurizes.  This cooling effect can largely offset the heating from the ICM compression and (certainly to within the accuracy of our data) mask any remaining signs of compressional heating.  This explains the lack of hot gas in or around the cavities.  In this evolutionary phase, the cavities will buoyantly rise within the cluster potential on a timescale a factor of a few longer than the sound crossing time of $\\sim 2 \\times 10^7\\,{\\rm yrs}$. As they rise buoyantly and expand, the relativistic electron population will undergo synchrotron, inverse Compton, and adiabatic energy losses. The synchrotron and inverse Compton losses result in a high-frequency cut-off that gradually marches to lower and lower radio frequencies. Using the standard formulae for synchrotron losses (e.g., Rybicki \\& Lightman 1979), it is readily shown that, assuming an isotropic relativistic electron distribution evolving in a constant or decreasing strength magnetic field, the high-frequency cut-off of the synchrotron spectrum will obey \\begin{equation} \\nu_{\\rm cut} \\lesssim 26\\,\\left(\\frac{B}{60\\,\\mu{\\rm G}}\\right)^{-3}\\left(\\frac{t}{20\\,{\\rm Myr}}\\right)^{-2}\\,{\\rm MHz}, \\end{equation} where approximate equality corresponds to the case where the magnetic field and the particle pressure are constant in time.  This expression assumes no fresh injection or acceleration of relativistic electrons (which would turn the cut-off into a spectral break), and hence only applies once the radio-lobes are no longer supplied by active jets (i.e., after the radio-source ``dies''). The ICM pressure at the location of the ghost-cavities is measured to be approximately $p \\approx 10^{-10}\\,{\\rm ergs\\,cm^{-3}}$. If we assume that the synchrotron emitting plasma is in pressure equilibrium with the surrounding ICM (which is very likely to be true for the ghost cavities) and furthermore, that the magnetic field in the plasma has approximately equipartition strength and is tangled on scales small compared to the cavity size, this pressure gives us a field strength of $B \\approx 60\\,\\mu{\\rm G}$.  Thus, assuming ICM/cavity pressure balance and equipartition magnetic fields, we can see from eqn.(1) that the cavities will fade out of the 1.4 GHz band only 4 Myr or so after the outburst of the radio-galaxy activity has ceased. Since we believe the ghost cavities to be approximately 20 Myr old (Heinz et al.~2002), we see that there has been ample time for the plasma filling the cavities to fade out of the higher frequency radio bands {\\it if the magnetic field posesses roughly equipartition strength}.  Studies with ROSAT, {\\em Chandra}, and {\\em XMM-Newton} have allowed the magnetic field strengths of several radio lobes to be estimated through the direct detection of the X-rays thought to be produced by inverse Compoton scattering of the Cosmic Microwave Background (CMB) by the relativistic electrons (Leahy \\& Gizani 2001; Hardcastle et al.  2002; Grandi et al. 2003; see also Wilson, Young, \\& Shopbell 2001 for related arguments in the hot spots of Cygnus A). In these studies, it is typically found that the magnetic field is at least a factor of two lower than the equipartition value.  Even if the magnetic field has half of the equipartition field strength, there is sufficient time for the 1.4\\,GHz emission from the ghost cavities to fade.  Having put forward a fairly traditional hypothesis for the formation and evolution of the X-ray cavities, we now proceed to consider the complexities special to Abell~4059.  \\subsection{Possible formation mechanisms for the SW ridge}  One of the most striking feature in the X-ray morphology of A~4059 is the bright and cool SW ridge.  The SW edge of this ridge appears to be surface across which the temperature and entropy of the gas change significantly with little or no change in pressure.  In many ways, this is similar to the ``cold fronts'' that have been observed in many clusters (Markevitch et al. 2000; Vikhlinin, Markevitch \\& Murray 2001).  Here, we will discuss four possible formation mechanisms for this structure.  \\subsubsection{A cool disk associated with a rotating cooling flow}  Huang \\& Sarazin (1998), who were the first to note the SW ridge using {\\it ROSAT} HRI data, suggested that it might be the rotationally-supported disk of cooled gas expected to form at the center of a rapidly-rotating cooling flow.  The notion that such a disk-like structure can form in high angular momentum cooling flows has gained support from axisymmetric hydrodynamically simulations (Garasi et al. 1998), although there are still unresolved questions as to the effect that turbulent angular momentum transport may have on the formation and stability of such disks (Nulsen, Stewart, \\& Fabian 1984).  However, it is clear from the high-resolution {\\it Chandra}-ACIS data that the SW ridge does {\\it not} extend NE of the cluster center, i.e., it is one-sided.  This can be seen in both the total intensity map (Fig.~\\ref{fig:acis_image}) and, more clearly, in the temperature map (Fig.~\\ref{fig:maps}b).  This runs counter to the idea that the SW ridge is part of a large ($\\sim 20$\\,kpc) disk at the center of the cluster.  Thus, just on the basis of morphology, we can reject the hypothesis that this structure is part of a disk associated with a rotating cooling flow.  \\subsubsection{Cool wakes of buoyantly rising radio plumes}  Numerical simulations of the buoyant phase of a radio-galaxies evolution show that appreciable amounts of ICM from the cluster core can become entrained in the ``wake'' of a buoyantly rising plume of radio plasma (Br\\\"uggen et al. 2002; Reynolds, Heinz, \\& Begelman 2002).  This material adiabatically decompresses and cools as it is dragged upwards in the cluster potential, and would appear as distinct filaments of cold and dense material strung out along the path of the buoyant plume.  As discussed by Young, Wilson \\& Mundell (2002), these wakes of cold gas are probably responsible for the arc-like feature seen in {\\it ROSAT}-HRI and {\\it Chandra}-ACIS observations of M87 and the core of the Virgo cluster.  This structure is composed of narrow filaments or columns of cold gas (with $kT\\sim 1\\keV$, compared with $kT\\sim 3\\keV$ for the surrounding ICM), probably in pressure equilibrium with their surroundings, that extend for 2--3\\,arcmins East and South-West of M87.  They are coincident with, but more more narrowly confined than, the 90\\,cm radio arc observed by Owen, Eilek \\& Kassim (2000). This supports the idea that the filament has been entrained and pulled out of the central parts of M87/Virgo by a buoyantly rising plume.  However, it seems unlikely that such a model can explain the SW ridge of A~4059.  There is no indication of any radio-lobe (even a very old one) in the SW direction, i.e., there is no radio emission and no ICM cavity in that quadrant of the cluster.  Furthermore, the SW ridge does not take on the form of a narrow filament reaching out from the cD galaxy, as would be expected for wake material on the basis of both the numerical simulations and the Young et al. (2002) observations of Virgo.  Instead, the SW ridge is a rather broad and flaring feature extending from the cD galaxy.  On the basis of these two observations, we reject the hypothesis that the SW ridge corresponds to cool material that has been entrained in the wake of a buoyantly rising plume of radio-plasma.  \\subsubsection{The accreted core of a cooler sub-cluster}  The discovery of a large dust lane in the HST-WFPC2 image of ESO349-G010 suggests the accretion of a dust and gas rich companion galaxy within the past $10^8$\\,yrs.  It is possible that this galactic-merger event was actually the late stages of the merger of a smaller galaxy cluster or group with A~4059.  In this case, one may attempt to identify the SW ridge with the remnant ICM core of the minor cluster.  The well known correlation between the X-ray luminosities and temperatures of galaxy clusters and groups, $L\\sim T^3$, means that the accreted minor system is likely to possess an ICM that is significantly cooler than the ambient temperature of A~4059 ($kT\\sim 4\\,keV$).  We note that similar ideas have been proposed to explain the cold fronts observed in other clusters (Markevitch et al. 2000; Bialek, Evrard \\& Mohr 2002; Nagai \\& Kratsov 2002).  However, this scenario may be problematic for the case of A~4059 (although it may well explain some of the classical cold fronts seen in other clusters).  While it is true that the ICM {\\it temperature} of the accreted subcluster may initially be cooler than that of A~4059, it will compress and heat as it enters the higher pressure environment of the richer cluster.  The relevant thermodynamic quantity to consider is the {\\it entropy} of the ICM cores of A~4059 and the accreted cluster.  In fact, for clusters of the mass of A~4059 and smaller, the entropy of the ICM core is almost constant from one cluster to another (Lloyd-Davies, Ponman \\& Cannon 2000; Mushotzky et al. 2003).  Thus, even if it evolved adiabatically, the ICM-core of the accreted group would be compressionally heated to approximately the ambient temperature of A~4059.  Any departure from adiabatic evolution (e.g., the effects of shocks) will only increase the entropy and temperature of the accreting ICM.  In order to produce a colder region, radiative cooling needs to dominate the evolution of the accreted core.  While this may be true, either fine tuning or significant feedback (either via conduction or radio-galaxy heating) is needed to prevent the core from cooling completely.   \\subsubsection{A compression front associated with bulk ICM motion}  The apparent displacement of the center of the cavities from the cluster center and the different position angle between the extended radio emission and the cavities suggests bulk motion of the ICM flow. Noting that the line connecting the centers of the two cavities misses the radio-galaxy core by approximately 10\\,arcsec, corresponding to 10\\,kpc, we can use estimates of the age of the radio source ($\\sim 20$\\,Myr; Heinz et al. 2002) to estimate that the ICM is flowing past the cD galaxy at a velocity of $500\\kmps$ projected onto the plane of the sky.  The different position angle between the axis connecting the ghost cavities and the current extended radio emission demands either a change in the radio-axis itself, or some rotation in the ICM flow. Both numerical simulations (e.g., Roettiger, Loken \\& Burns 1997) and {\\it Chandra} observations (e.g., Markevitch et al. 2003) suggest this kind of large scale ICM ``sloshing'' can readily occur after a major cluster merger.  In this picture, the bright and cool SW ridge is located at the position where we expect the radio-galaxy induced expanding ICM shell to be maximally compressed by the ICM flow.  The sharp SW edge of this feature is readily interpreted at the interface between the ambient ICM (which we suppose is flowing in a NE direction) and the expanding ICM shell formed by the same period of radio-galaxy activity that formed the X-ray cavities.  One might expect that such compression would heat this material, contrary to observations.  However, the fact that the cooling time of the ridge material is small demands that we consider radiative cooling effects.  In the simplest case of adiabatic compression in the bremsstrahlung regime, the cooling time is proportional to $n^{-2/3}$. Hence, a weak shock will slightly reduce the cooling timescale. Radiative cooling can be further aided by the kinematics of the radio-galaxy/cluster interaction, which keeps this material in the high pressure regions of the cluster core for longer.  Even given this, there appears to be a fine tuning problem; it is difficult to explain the cooling of the SW ridge unless it was on the verge of undergoing dramatic radiative cooling anyways.  A detailed exploration of these hydrodynamical and radiative questions will be deferred until future publications."
    },
    "0402/astro-ph0402125_arXiv.txt": {
        "abstract": "{I present brief status reports on three large observational projects that are designed to test our current understanding of the evolution of cataclysmic variables (CVs): The spectroscopic selection of new CVs in the Hamburg Quasar Survey, the search for pre-CVs based on Sloan colours and UK Schmidt/6dF multiobject spectroscopy, and the identification of CVs that descended from supersoft X-ray binaries using a \\textit{HST}/STIS far-ultraviolet spectroscopic survey.}  \\resumen{I present brief status reports on three large observational projects that are designed to test our current understanding of the evolution of cataclysmic variables (CVs): The spectroscopic selection of new CVs in the Hamburg Quasar Survey, the search for pre-CVs based on Sloan colours and UK Schmidt/6dF multiobject spectroscopy, and the identification of CVs that descended from supersoft X-ray binaries using a \\textit{HST}/STIS far-ultraviolet spectroscopic survey.}  \\addkeyword{Stars: Cataclysmic Variables} \\addkeyword{Stars: Evolution}  \\begin{document} ",
        "introduction": "The essential idea of the standard model of cataclysmic variable (CV) evolution (disrupted magnetic braking, King 1988) is that CVs evolves towards shorter periods due to a combination of angular momentum losses: magnetic braking (dominating in systems with orbital periods $P_{\\rm orb}\\ga 3$\\,h) and the less efficient gravitational radiation (dominating in systems with orbital periods $P_{\\rm orb}\\la 2$\\,h). The standard paradigm of CV evolution successfully explains the 2--3\\,h gap in the observed CV period distribution. However, most other predictions made by this model are in strong contrast with the properties of the known CV population.  Recently, a number of far-reaching modifications for the standard scenario have been proposed (see Spruit \\& Taam 2001, King \\& Schenker 2002, Schenker \\& King 2002 and Andronov et al. 2003). Unfortunately, none of them has been completely successful in tuning the predictions so that they fully agree with the observations. It is apparent that the disturbing disagreements between theory and observations have a common denominator: the possible impact of selection effects on the currently known population of CVs. In order to quantitatively test any theory of CV evolution it is necessary to establish a large and unbiased sample of CVs as well as of their progenitors, (pre-CVs, detached white dwarf/late type main sequence stars). Such an observational data base will also serve for future improvements of the theory of CV evolution. ",
        "conclusions": ""
    },
    "0402/astro-ph0402313_arXiv.txt": {
        "abstract": "{ The photometric evolution of M31-RV has been investigated on 1447 plates of the Andromeda galaxy obtained over half a century with the Asiago telescopes. M31-RV is a gigantic stellar explosion that occurred during 1988 in the Bulge of M31 and that was characterized by the appearance for a few months of an M supergiant reaching $M_{bol}=-$10. The 1988 outburst has been positively detected on Asiago plates, and it has been the only such event recorded over the period covered by the plates (1942-1993). In particular, an alleged previous outburst in 1967 is excluded by the more numerous and deeper Asiago plates, with relevant implication for the interpretative models of this unique event. We outline a close analogy in spectral and photometric evolution with those of V838~Mon which exploded in our Galaxy in 2002. The analogy is found to extend also to the closely similar absolute magnitude at the time of the sudden drop in photospheric temperature that both M31-RV and V838~Mon exhibited. These similarities, in spite of the greatly differing metallicity, age and mass of the two objects, suggest that the same, universal and not yet identified process was at work in both cases. ",
        "introduction": "Rich et al. (1989) discovered in 1988 a highly unusual stellar outburst in the Bulge of the Andromeda galaxy (M31), known since then as M31-RV (for ``red variable''). The event peaked at M$_{bol} \\approx -10$~mag and its spectrum closely resembled that of M supergiants, evolving from M0~I at discovery (Sept 5, 1988) to $>$M7~I about 58 days later when the brightness in the $V$ band had dropped by at least 4 mag (Rich 1990). Two similar events have been later identified in our Galaxy, V4332~Sgr that exploded in 1994 (Martini et al. 1999) and V838~Mon that erupted in 2002 (Munari et al. 2002a, Bond et al. 2003, and references therein).  The M31-RV event has been characterized by radiative luminosities in-between those of classical novae and supernovae. The mass of the ejected envelope (optically thick during the whole observed evolution) is uncertain but it is certainly larger than in typical novae and much less than in supernovae. The radiative and kinetic energetics place therefore M31-RV, and by analogy also V4332~Sgr and V838~Mon, in the gap between classical novae and supernovae, making them stars of special interest. So far, few theoretical attempts to explain their highly peculiar nature have been pursued. Soker and Tylenda (2003), to explain the energetics and multi-maxima behaviour of V838~Mon, have suggested the merging of two main sequence stars of masses 0.1$-$0.5~M$_\\odot$ and 1.5~M$_\\odot$, with the second one expanding to large radii, low temperature and high luminosity in response to the frictional energy dissipation of the cannibalized less massive companion. A similar scenario has been proposed by Retter and Marom (2003). They postulated the multi-maximum eruption of V838~Mon as the result of the swallowing of massive planets in close orbit around a parent star expanding while on the RGB (red giant branch) or AGB (asymptotic giant branch). A thermonuclear runaway (TNR) model was instead developed by Iben and Tutukov (1992) to explain M31-RV. The model envisages a binary system, composed of a WD and a low mass companion, that evolves to orbital periods shorter than 2 hours by loss of angular momentum via gravitational waves, without experiencing classical nova eruptions on the way to. The accretion at very low rates ($\\sim 10^{-11}$~M$_\\odot$~yr$^{-1}$) occurring onto a cold white dwarf (WD) can lead to the accumulation of a massive H-rich envelope of the order of $\\sim$0.05~M$_\\odot$ before this is expelled in a gigantic hydrogen shell flash (some 10$^3$ times the mass expelled in a typical nova eruption). Friction energy dissipation of the binary revolving within such a massive and dense common envelope can raise the drag luminosity to 10$^7$~L$_\\odot$, with as much as 10$^6$~L$_\\odot$ ($\\gg$~L$_{\\rm Eddington}$) coming out in the form of radiation.  \\begin{table*}[!t] \\caption{Plates of Andromeda galaxy from the Asiago archive inspected in the M31-RV search.  The last two columns detail the number of plates obtained in 1988, when the outburst of M31-RV occurred, and in 1967, when an alleged previous outburst should have taken place.} \\centering \\begin{tabular}{ccrrrrrcc} \\hline &&\\\\ \\multicolumn{2}{c}{telescope}& \\multicolumn{1}{c}{focal} & \\multicolumn{1}{c}{first} & \\multicolumn{1}{c}{last}  & \\multicolumn{1}{c}{N$_{tot}$}& \\multicolumn{1}{c}{main} & \\multicolumn{1}{c}{plates in} & \\multicolumn{1}{c}{plates in} \\\\ &&\\multicolumn{1}{c}{length}&\\multicolumn{1}{c}{plate}&\\multicolumn{1}{c}{plate} &\\multicolumn{1}{c}{plates}&\\multicolumn{1}{c}{band} &\\multicolumn{1}{c}{1967}&\\multicolumn{1}{c}{1988}\\\\ &&\\\\ \\hline &&\\\\ 1.22 m   & Newton     & 6.0 m & 29 Oct 1942 & 26 Aug 1992 & 831 & 795 in $B$ band & 25 & 2 \\\\ 1.22 m   & Cassegrain &19.1 m & 31 Oct 1961 & 27 Nov 1972 & 94  &  93 in $B$ band & 4  &   \\\\ 1.82 m   & Cassegrain &16.4 m & 05 Aug 1973 & 09 Dec 1988 & 194 & 177 in $B$ band &    & 2 \\\\ 67/92 cm & Schmidt    & 2.2 m & 02 Oct 1965 & 17 Dec 1993 & 291 & 264 in $B$ band & 2  & 8 \\\\ 40/50 cm & Schmidt    & 1.0 m & 14 Oct 1958 & 05 Mar 1986 &  37 &  28 in $B$ band &    &   \\\\ &&\\\\ \\hline \\end{tabular} \\end{table*}  Common to all models above is the uniqueness of the event: the progenitor can experience a single such outburst in its life. In the Soker and Tylenda (2003) and Retter and Marom (2003) approaches, it is the result of a merger event that obviously cannot be repeated. In the Iben and Tutukov (1992) model an extremely long time (a sizable fraction of a Hubble time) is required to accrete at very low rates $10^{-2}$~M$_\\odot$ on a WD that had to cool to low temperatures.  \\begin{figure} \\centerline{\\psfig{file=0716fig1.ps,width=8.7cm}} \\caption[]{Finding chart for the {\\em BVR$_{\\rm C}$I$_{\\rm C}$} comparison sequence listed in Table~2. The ``{\\em X}'' marks the location of M31-RV.} \\end{figure}  Therefore, the report by Sharov (1990) about a second outburst of M31-RV in 1967, 20 years before the main one, is something that deserves careful scrutiny and independent verification. If confirmed, it would have profound consequences on the theoretical modeling of M31-RV, V4332~Sgr and V838~Mon, perhaps even more than the discovery of a massive and young B3~V companion to the latter (Munari et al. 2002b).  The recent eruption of V838~Mon has considerably revitalized the interest on this class of objects. In anticipation of a growing modeling effort by the community, we decided to take advantage of the Asiago plate archive to evaluate the reality of a second outburst of M31-RV and to investigate its long term photometric evolution. ",
        "conclusions": ""
    },
    "0402/astro-ph0402580_arXiv.txt": {
        "abstract": "We describe the design and calibration of an external cryogenic blackbody calibrator used for the first two flights of the Absolute Radiometer for Cosmology, Astrophysics, and Diffuse Emission (ARCADE) instrument. The calibrator consists of a microwave absorber weakly coupled to a superfluid liquid helium bath. Half-wave corrugations viewed 30\\deg ~off axis reduce the return loss below -35 dB. Ruthenium oxide resistive thermometers embedded within the absorber monitor the temperature across the face of the calibrator. The thermal calibration transfers the calibration of a reference thermometer to the flight thermometers using the flight thermometer readout system. Data taken near the superfluid transition in 8 independent calibrations 4 years apart agree within 0.3 mK, providing an independent verification of the thermometer calibration at temperatures near that of the cosmic microwave background. ",
        "introduction": "The Absolute Radiometer for Cosmology, Astrophysics, and Diffuse Emission (ARCADE) is a balloon-borne instrument designed to measure the temperature of the cosmic microwave background at centimeter wavelengths \\cite{kogut/etal:2004}. ARCADE uses a set of narrow-band cryogenic radiometers to compare the sky to an external blackbody calibration target, in order to detect or limit deviations from a blackbody spectrum. At centimeter wavelengths, raw sensitivity is not an important design criterion; the instrument is designed instead to reduce or eliminate major sources of systematic uncertainty. The instrument is fully cryogenic; all major components are independently temperature-controlled to remain near 2.7 K, isothermal with the signal from deep space. Boiloff helium vapor, vented through the aperture plane, forms a barrier between the instrument and the atmosphere at 35 km altitude; there are no windows or warm optics to correct. An independently controlled blackbody calibrator located on the antenna aperture plane rotates to cover each antenna in turn, so that each antenna alternately views the sky or a known blackbody. The cryogenic design minimizes the effect of internal reflection or absorption: any residual instrumental signals cancel to first order in the sky--calibrator comparison. All radiometers view the same calibrator in turn, eliminating cross-calibration uncertainties when comparing the sky temperature between different frequency channels.  The external calibrator is critical to the ARCADE experiment. It establishes a common blackbody reference for all frequency channels, and establishes an absolute temperature reference for comparison with other instruments. Assuring the electromagnetic performance of the calibrator, though non-trivial, is straightforward. The ability to achieve ARCADE's science goals can thus be traced to precision and accuracy with which the physical temperature distribution within the calibrator can be monitored. ",
        "conclusions": ""
    },
    "0402/astro-ph0402549_arXiv.txt": {
        "abstract": "We discuss identifications for 18 sources from our MAMBO 1.2mm survey of the region surrounding the NTT Deep Field. We have obtained accurate positions from Very Large Array 1.4\\,GHz interferometry and in a few cases IRAM mm interferometry, and have also made deep BVRIzJK imaging at ESO.  We find thirteen 1.2mm sources associated with optical/near-infrared objects in the magnitude range K=19.0 to 22.5, while five are blank fields at K$>$22.  We argue from a comparison of optical/near-infrared photometric redshifts and radio/mm redshift estimates that two of the thirteen optical/near-infrared objects are likely foreground objects distinct from the dust sources, one of them possibly lensing the mm source.  The median redshift of the radio-identified mm sources is $\\sim$2.6 from the radio/mm estimator, and the median optical/near-infrared photometric redshifts for the objects with counterparts $\\sim$2.1.  This suggests that those radio-identified mm sources without optical/near-infrared counterparts tend to lie at higher redshifts than those with optical/near-infrared counterparts. Compared to published identifications of objects from 850$\\mu$m surveys of similar depth, the median K and I magnitudes of our counterparts are roughly two magnitudes fainter and the dispersion of I-K colors is less. Real differences in the median redshifts, residual mis-identifications with bright objects, cosmic variance, and small number statistics are likely to contribute to this significant difference, which also affects redshift measurement strategies. Some of the counterparts are red in J-K ($\\gtrsim$20\\%), but the contribution of such mm objects to the recently studied population of near-infrared selected (J$_s$-K$_s$$>$2.3) high redshift galaxies is only of the order a few percent. The recovery rate of MAMBO sources by pre-selection of optically faint radio sources is relatively low ($\\sim$25\\%), in contrast to some claims of a higher rate for SCUBA sources ($\\sim$70\\%).  In addition to this difference, the MAMBO sources also appear significantly fainter ($\\sim$1.5 magnitudes in the I-band) than radio pre-selected SCUBA sources.  We discuss basic properties of the near-infrared/(sub)mm/radio spectral energy distributions of our galaxies and of interferometrically identified submm sources from the literature. From a comparison with submm objects with CO-confirmed spectroscopic redshifts we argue that roughly two thirds of the (sub)mm galaxies are at z$\\gtrsim$2.5. This fraction is probably larger when including sources without radio counterparts. ",
        "introduction": "\\label{Introduction}  With the discovery of distant submm and mm galaxies in blank field surveys and cluster lens assisted surveys \\citep[see][for a review and references]{blain02}, it was immediately recognized that this population holds important clues for the understanding of the formation and evolution of galaxies. A significant part of the cosmic submm background is produced by these objects that must be extremely luminous (L$_{IR}\\sim 10^{12-13}$L$_\\odot$) distant (z$\\gtrsim$1) infrared galaxies powered by intense star formation and/or powerful AGN. X-ray data argue in most cases against (Compton-thin) AGN and in favor of intense star formation dominating their luminosity \\citep[e.g.,][]{alexander03}. These high star formation rates ($\\approx$100-1000\\,M$_\\odot$/yr) and the similarity of co-moving space densities of submm sources and local ellipticals suggest that they are likely indicating the formation of massive spheroids. This opens a direct route to locating the formation of spheroids between the two extremes of an early formation similar to the classical `monolithic collapse', and a late formation in hierarchical merging. More specifically, properties and mass functions for these high redshift objects are robust tests for current hierarchical models of galaxy formation \\citep{guiderdoni98,kauffmann99,somerville01,baugh03}. Indeed there is evidence \\citep[e.g.,][]{genzel03, neri03} that these models need some modification to reproduce the space densities of the massive high redshift submm galaxies and their quiescent phases that must exist since their star formation rates and gas content suggest a duty cycle of the bright (sub)mm phase well below one.  Much of this promise will only fully come to fruition after a difficult process of identification and spectroscopic redshift determination. Based on optical spectroscopy, significant progress in the identification work was done by \\citet{chapman03a} who presented 10 redshifts of radio identified counterparts of SCUBA galaxies.  But still, less than a dozen optical/near-infrared redshifts for suggested counterparts have been confirmed by CO observations as the true redshift of the submm source \\citep[][Greve et al. in prep, and SMM02399-0314, Kneib in prep.]{frayer98, frayer99, neri03}.  Accurate positions from radio or mm interferometry are available for just a few dozen non-radio-preselected submm galaxies \\citep[e.g.,][]{downes99, smail00, eales00, gear00, lutz01, ivison02, ledlow02, webb03a,webb03b} and very few mm-selected galaxies \\citep{bertoldi00, dannerb02}.  These accurate positions are indispensable for reliable optical/near-infrared identifications since several possible optical/near-infrared counterparts are usually found in the several arc second radius error circles of the (sub)mm surveys. This identification step is the subject of the present work. Of the two interferometric identification methods in use for (sub)mm sources, mm interferometry has the advantage of directly and unambiguously locating the dust emission, but it is very time consuming with current instruments, with tens of hours typically needed for a single object in the small field of view.  Because of the tight radio/far-infrared relation for star-forming galaxies \\citep[e.g.,][]{dejong85, helou85, condon92}, radio interferometry is the next best option, with the advantage of the large VLA primary beam covering one of the currently typical (sub)mm survey fields entirely. Of the brighter (sub)mm sources from present surveys, deep 1.4GHz VLA maps will detect all but the highest redshift ones \\citep{carilli99,carilli00a,barger00}, and the risk of false associations of (sub)mm and radio sources is modest for the 1.4GHz source counts at the relevant flux levels of tens of $\\mu$Jy \\citep[e.g.,][]{richards00}.  We are building on our mm survey (Bertoldi et al. 2003, in preparation) that uses the MAMBO array \\citep{kreysa98} at the IRAM 30m telescope to cover three fields, the Lockman Hole, Abell 2125, and a region centered on, but larger than the NTT Deep Field \\citep[NDF,][]{arnouts99}. This paper focuses on the NDF region.  \\citet[hereafter Paper~I]{dannerb02} have presented results for three of the brightest NDF mm sources using the IRAM Plateau de Bure mm interferometer (PdBI) and BVRIK imaging. We now increase to 18 the number of NDF mm sources with accurate positions by using our VLA interferometry, and introduce additional z- and J-band optical/near-infrared imaging obtained since publishing Paper~I. We thus identify counterparts and start assembling optical to radio spectral energy distributions (SEDs) of the objects constraining their nature and redshift.  The improved statistics is the basis for a discussion of the properties of mm galaxies and their relation to, or differences from, the submm population. Throughout the paper we adopt the cosmological parameters $H_0$=70\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_M$=0.3, $\\Omega_\\Lambda$=0.7. ",
        "conclusions": "We have discussed optical/near-infrared identifications for those 18 of 42 sources in our MAMBO 1.2mm map of the NTT Deep Field region for which interferometric positions are available through a VLA 1.4GHz map and in three cases IRAM PdBI mm interferometry. In addition to being the basis for identifications, our deep BVRIzJK imaging allows the derivation of optical/near-infrared photometric redshifts for the counterparts and for nearby objects. In comparison with radio/submm redshift estimates, these photometric redshifts suggest that two of the optical/near-infrared sources close to interferometric positions are in the foreground, in one case likely lensing the background mm source. One strongly lensed object in this sample is consistent with expectations for the (sub)mm population \\citep{blain99,chapman02c}. This leaves us with eleven detections of counterparts at magnitudes of K=19 to 22.5, and seven limits or blank fields with most limits at K$>$22. The I-K and J-K colors of the counterparts are consistent with redshifted SEDs similar to local ultra-luminous infrared galaxies, and likely with a similarly large spread of the rest frame UV/optical SED properties. The counterparts of mm sources contribute to the recently discussed population of J-K$>$2.3, K$<$22.5 high redshift galaxies, but only at the few percent level for the current mm survey depths.  At least $\\sim$20\\% of the MAMBO sources with radio counterparts have J-K$>$2.3 and K$<$22.5.  The counterparts to NDF mm sources are on average about 2 magnitudes fainter than counterparts presented for the 8mJy 850$\\mu$m survey which is of similar depth and are similarly faint compared to radio pre-selected submm sources. Remaining mis-identifications, redshift or temperature differences between the two populations, and small number statistics/cosmic variance all may contribute to this difference, at levels that are hard to quantify from currently available data. Our result reinforces the view that direct (e.g., wide band CO) spectroscopic redshifts may be necessary for a substantial fraction of the (sub)mm population which is very faint in the optical/near-infrared. From a comparison of near-infrared/submm/radio spectral indices with those of submm sources with CO-confirmed redshifts we suggest that the fraction of (sub)mm galaxies at z$>$2.5 is about two thirds for the interferometrically located ones and larger when adding the radio-undetected part of the population."
    },
    "0402/astro-ph0402019_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:1} At a distance of 8 kpc only, the center of the Galaxy provides a unique testbed for studies about the physics of the interstellar medium and star formation in the nuclei of galaxies. Since it shares many properties of other, more spectacular nuclei, a detailed investigation of the Galactic Center is a necessary step to gain better understanding of the physical processes governing galactic nuclei in general. \\\\ We review the  characteristics of the central molecular cloud layer, the physical properties of these clouds, and their interplay with the magnetic field. ",
        "conclusions": ""
    },
    "0402/astro-ph0402533_arXiv.txt": {
        "abstract": "{ We present XMM-Newton results on the spatially resolved temperature profiles of eight massive galaxy clusters of a volume-limited sample at redshifts $z\\sim0.3$ (REFLEX-DXL sample) and an additional luminous cluster at $z=0.2578$, selected from the REFLEX survey. Useful temperature measurements could be performed out to radii with overdensity 500 ($r_{500}$). The scaled temperature distributions show good similarities. We discovered diversities in the temperature gradients at the outer cluster radii with examples of both flat and strongly decreasing profiles which call for different physical interpretations. We found an indication of the 'warm-hot' gas existing in or around the hot clusters. Using RXCJ0307.0$-$2840 we demonstrate that the errors on the mass estimates are within 25\\% up to $r_{500}$. ",
        "introduction": "The most massive clusters are especially important in tracing large scale structure (LSS) evolution since they are expected to show the largest evolutionary effects. In hierarchical modeling the structure of the X-ray emitting plasma in the most massive clusters is essentially determined by gravitational effects and shock heating. With decreasing cluster mass and intracluster medium (ICM) temperature, non-gravitational effects play an important role before and after the shock heating (Voit \\& Bryan 2001; Voit et al. 2002; Zhang \\& Wu 2003; Ponman et al. 2003).  Therefore, the most massive clusters provide the cleanest results in comparing theory with observations.  In this project we are analysing an almost volume-complete sample of thirteen X-ray luminous ($L_{X}~\\geq~10^{45}~{\\rm erg~s^{-1}}$ for 0.1--2.4~keV) clusters selected from the ROSAT-ESO Flux-Limited X-ray (REFLEX) galaxy cluster survey (B\\\"ohringer et al. 2001) in the redshift interval $z=0.27$ to $0.31$ (see Fig.~\\ref{f:reflex1}). \\begin{figure} \\includegraphics% [width=5.6cm]{plots/f1.ps} \\caption{X-ray luminosity-redshift distribution of the REFLEX clusters. The box shows the selection of the 13 REFLEX-DXL clusters (see B\\\"ohringer et al. 2003).} \\label{f:reflex1} \\end{figure} \\begin{figure} \\includegraphics[width=6cm]{plots/f2.ps} \\caption{A preliminary REFLEX-DXL temperature function (crosses) is compared to the $<z>=0.05$ (solid circles) and $<z>=0.38$ (open circles) temperature functions from Henry (2000). See B\\\"ohringer et al. 2003.} \\label{f:reflex2} \\end{figure} There is only a very small correction to the volume completeness with a well known selection function for $L_{X}~\\geq~10^{45}~{\\rm erg~s^{-1}}$ at the higher redshift as described in B\\\"ohringer et al. (2003). With this REFLEX-DXL (Distant X-ray Luminous) sample, we obtain reliable ICM temperatures to measure the cluster masses based on the high resolution observations from XMM-Newton (Zhang et al 2004). Since peculiarities in the cluster structure introduce a scatter in the mass-temperature relation and since in particular on-going cluster mergers can lead to a temporary increase in the cluster temperature and X-ray luminosity (Randall et al. 2002), we aim for a detailed study of the deep XMM-Newton observations described here. The clusters are also scheduled for a detailed spectroscopic study of the cluster dynamics with the ESO-VLT-VIMOS instrument. One prime goal is to study the temperature function evolution (see Fig.~\\ref{f:reflex2}, B\\\"ohringer et al. 2003) by comparing our sample with more nearby and more distant clusters in literature. The selection of the REFLEX-DXL sample and its properties are described in detail in B\\\"ohringer et al. (2003). The method is well described in Zhang et al. (2004), which is established for a reliable determination of the spatially resolved temperature profiles for the REFLEX-DXL clusters. XMM-Newton with its superior sensitivity combined with its good spatial resolution provides the best means for such studies (Arnaud et al. 2002). Previously, large data sets on cluster temperature profiles have been compiled from ASCA (e.g. Markevitch et al. 1998; White 2000; Finoguenov et al. 2001a; Finoguenov et al. 2002; Sanderson et al. 2003) and BeppoSAX observations (Molendi \\& De Grandi 1999; Ettori et al. 2002).  In this proceeding we contribute the temperature profile measurements, discuss physics of diversity, describe an indication for soft excess from warm InterGalactic Medium (IGM), and present the mass estimates. We adopt a flat $\\Lambda$CDM cosmology with the density parameter $\\Omega_{\\rm m}=0.3$ and the Hubble constant $H_{\\rm 0}=70$~km~s$^{-1}$~Mpc$^{-1}$. Error bars correspond to the 68\\% confidence level, unless explicitly stated otherwise.  \\section {Temperature distributions} \\label{s:method}  We developed a reliable double background subtraction method for the XMM-Newton data reduction, in which we use the XMM-Newton observations of the Chandra Deep Field South (CDFS) background and model the difference of the target and CDFS backgrounds with the data from the outer Field of view (FOV). The details are available in Zhang et al (2004).  Comparing the spectral results, we have noted that the differences between the global temperatures of the regions covering radii of $0.5<r<4^\\prime$ and $r<8^\\prime$, respectively, are caused by systematic differences in the temperature gradients. For a more detailed study of the temperature profiles we divide the cluster regions into the five annuli 0--0.5$^\\prime$, 0.5--1$^\\prime$, 1--2$^\\prime$, 2--4$^\\prime$, and 4--8$^\\prime$.  We use $>1$~keV band except for RXCJ0658.5$-$5556. We apply the 2--12~keV band for this high temperature cluster.  The temperature profile of each cluster is shown in Zhang et al. (2004). We detect the gradients in the spatially resolved temperature profiles with an accuracy of better than 10 to 20\\% in the $r<4^\\prime$ region. The temperatures vary as a function of the radius by a factor of 1.5 to 2. To some degree the difference of the central structure might reveal the effect of non-gravitational processes and radiative cooling. No significant cooling gas lower than 2 keV is found in the center.  In Fig.~\\ref{f:ktcomp2}, we present the scaled temperatures of eight REFLEX-DXL clusters together with an additional cluster at slightly lower redshift $z=0.2578$. The radii are scaled by the virial radii obtained from the M-T relation in Bryan and Norman (1998). The temperatures are scaled by the temperatures of the regions covering radii of $0.5<r<4^\\prime$, which cover almost the whole X-ray emission regions of the clusters. The grey shadow shows the self-similar temperature profile of Markevitch et al. (1998). The temperature profiles are on average flattened and the high accuracy of our data shows that the decrease in our temperature profiles happens at further out radii. There are also diversities in the temperature profiles which were not observed before. \\begin{figure*} \\centering \\includegraphics[angle=-90,width=12.cm] {plots/f3.ps} \\caption{Scaled temperature profiles fitted in the 1--10~keV band, noted the temperature profile of RXCJ0658.5$-$5556 was fitted in the 2--12~keV band.} \\label{f:ktcomp2}% \\end{figure*} ",
        "conclusions": "\\label{s:conclusion}  We obtain spatially resolved, consistant X-ray temperature profiles from three instruments for 9 clusters with an accuracy of better than 10--20\\% in the $<4^\\prime$ region. We found good similarities in the scaled temperature distributions. While the high accuracy of our data shows that the decrease in our data happens at further out radii comparing to the self-similar temperature profile of Markevitch et al. (1998). Temperature structures are easily disturbed by some physical processes and are thus complicated. Differences of the central structures reveal the effect of non-gravitational process. No significant cooling gas ($<2$~keV) is found in the center. The variations of the temperature distributions in the outer region call for different formation histories (e.g. Finiguenov et al 2001b). Among the sample, six clusters show regular temperature structures, while three do not. Combination of the regular temperature profiles and symmetric 2-D maps suggest a relaxed dynamical state. Fluctuations noted in 2-D maps might be diluted in the temperature profiles. The uncertainty in the temperature measurements is a key point to the accurate determinations of mass, gas mass fraction, and thus affects the L-T and M-T scaling relations."
    },
    "0402/astro-ph0402475_arXiv.txt": {
        "abstract": "We have compiled $L'$ (3.4--$4.1~\\mu$m) and $M'$ (4.6--$4.8~\\mu$m) photometry of 63 single and binary M, L, and T dwarfs obtained at the United Kingdom Infrared Telescope using the Mauna Kea Observatory (MKO) filter set.  This compilation includes new $L'$ measurements of 8 L dwarfs and 13 T dwarfs and new $M'$ measurements of 7 L dwarfs, 5 T dwarfs, and the M1 dwarf Gl~229A.  These new data increase by factors of 0.6 and 1.6, respectively, the numbers of ultracool dwarfs ($T_{\\rm eff} \\lesssim 2400$~K) for which $L'$ and $M'$ measurements have been reported.  We compute $L_{\\rm bol}$, BC$_K$, and $T_{\\rm eff}$ for 42 dwarfs whose flux-calibrated $JHK$ spectra, $L'$ photometry, and trigonometric parallaxes are available, and we estimate these quantities for 9 other dwarfs whose parallaxes and flux-calibrated spectra have been obtained.  BC$_K$ is a well-behaved function of near-infrared spectral type with a dispersion of $\\sim 0.1$~mag for types M6--T5; it is significantly more scattered for types T5--T9.  $T_{\\rm eff}$ declines steeply and monotonically for types M6--L7 and T4--T9, but is nearly constant at $\\sim 1450$~K for types L7--T4 with assumed ages of $\\sim 3$~Gyr.  This constant $T_{\\rm eff}$ is evidenced by nearly unchanging values of $L'$--$M'$ between types L6 and T3.  It also supports recent models that attribute the changing near-infrared luminosities and spectral features across the L--T transition to the rapid migration, disruption, and/or thinning of condensate clouds over a narrow range of $T_{\\rm eff}$.  The $L'$ and $M'$ luminosities of early-T dwarfs do not exhibit the pronounced humps or inflections previously noted in the $I$ through $K$ bands, but insufficient data exist for types L6--T5 to assert that $M_{L'}$ and $M_{M'}$ are strictly monotonic within this range of types.  We compare the observed $K$, $L'$, and $M'$ luminosities of L and T dwarfs in our sample with those predicted by precipitating-cloud and cloud-free models for varying surface gravities and sedimentation efficiencies.  The models indicate that the L3--T4.5 dwarfs generally have higher gravities (log~$g = 5.0$--5.5) than the T6--T9 dwarfs (log~$g = 4.5$--5.0).  The predicted $M'$ luminosities of late-T dwarfs are 1.5--2.5 times larger than those derived empirically for the late-T dwarfs in our sample.  This discrepancy is attributed to absorption at 4.5--$4.9~\\mu$m by CO, which is not expected under the condition of thermochemical equilibrium assumed in the models.  Our photometry and bolometric calculations indicate that the L3 dwarf Kelu-1 and the T0 dwarf SDSS~J042348.57--041403.5 are probable binary systems.  We compute log$(L_{\\rm bol}/L_{\\odot}) = -5.73 \\pm 0.05$ and $T_{\\rm eff} = 600$--750~K for the T9 dwarf 2MASSI~J0415195--093506, which supplants Gl~570D as the least luminous and coolest brown dwarf presently known. ",
        "introduction": "The number of known ultracool dwarfs -- dwarfs with effective temperatures $T_{\\rm eff} \\lesssim 2400$~K -- has grown dramatically over the last seven years, primarily because of the sizes and depths of the DEep Near-Infrared Survey of the Southern Sky \\citep[DENIS;][]{epc97}, the Two-Micron All Sky Survey \\citep[2MASS;][]{skr97}, and the Sloan Digital Sky Survey \\citep[SDSS;][]{yor00}.  These surveys, and others of lesser scope, have revealed ultracool dwarfs in numbers sufficient to populate a distribution of temperatures ranging from the coolest hydrogen-burning stars \\citep[$1550 \\lesssim T_{\\rm eff} \\lesssim 1750$~K;][]{bur93,cha00} to the coolest known brown dwarf ($600 \\lesssim T_{\\rm eff} \\lesssim 750$~K; this paper).  Consequently, two new spectral types, L and T, have been defined in order to classify dwarfs cooler than spectral type M \\citep[hereafter G02] {kir99b,mar99b,bur02a,geb02}.  The spectra of L dwarfs are characterized by absorption from neutral alkali metals (e.g., K, Na, Cs, and Rb) and metallic hydrides (e.g., FeH and CrH) at red wavelengths and by absorption from CO and H$_2$O at near-infrared wavelengths.  In contrast, the optical spectra of T dwarfs are dominated by pressure-broadened K~{\\small I} and Na~{\\small I} absorption lines, and their near-infrared spectra are sculpted by broad absorption bands of CH$_4$ and H$_2$O and collisionally-induced absorption (CIA) by H$_2$.  Most observational studies of ultracool dwarfs concern the spectral region 0.6--2.5~$\\mu$m.  This region is favored because the flux distributions of these objects peak around 1.2~$\\mu$m and because the spectral sensitivities of modern photoelectronic detectors coincide with the relatively dark and transparent atmospheric windows in this region.  Recently, attention has been given to the photometry of ultracool dwarfs longward of the $K$ bandpass. This attention has been motivated partly by the need to better constrain the bolometric luminosities of ultracool dwarfs and by the prospects for observing even cooler brown dwarfs or planets at these wavelengths with space-based infrared telescopes.  The intrinsic faintness of L and T dwarfs and the increasing brightness and variability of the sky make ground-based observations of these objects in the $L$ (2.5--4.0~$\\mu$m) and $M$ (4.1--5.5~$\\mu$m) bandpasses difficult and time-consuming.  Although $L$-band photometry of L and T dwarfs is extensive \\citep[hereafter L02]{jon96,leg98,tok99,ste01,rei02,leg02a}, $M$-band photometry has been published for only six of these objects (\\citealt{mat96,rei02}; L02).  Spectroscopic studies of brown dwarfs in the $L$ or $M$ bandpasses have so far been limited to three L dwarfs and the archetypal T dwarf, Gl~229B.  The $L$ bandpass contains the Q-branch of the fundamental absorption band of CH$_4$, which is situated near 3.3~$\\mu$m.  This absorption band appears as early as spectral type L5 \\citep[$T_{\\rm eff} \\approx 1700$~K;][]{nol00} and is deep and broad in the spectrum of the T6 dwarf Gl~229B \\citep[$T_{\\rm eff} \\approx 900$~K;][]{opp98}.  The $M$-band spectrum of Gl~229B features a broad but shallow absorption trough from 4.5~$\\mu$m to 4.9~$\\mu$m and a narrow peak at 4.67~$\\mu$m, both of which are attributed to the 1-0 vibration-rotation band of CO \\citep{nol97,opp98}.  These features reveal a CO abundance that is over 1000 times larger than expected under conditions of CO$\\leftrightarrow$CH$_4$ thermochemical equilibrium, indicating that CO is rapidly transported outward from warmer, CO-rich layers of the atmosphere \\citep{feg96,gri99,sau00,sau03}.  The $M$-band fluxes of two other T dwarfs, SDSS~J125453.90--012247.4 and 2MASS~J05591914--1404488, are reportedly well below the levels expected for CO$\\leftrightarrow$CH$_4$ equilibrium, which suggests that vertical mixing of CO within the atmospheres of T dwarfs is common \\citep[L02;][]{sau03}.  In this paper, we present new 3.4--$4.1~\\mu$m and 4.6--$4.8~\\mu$m photometry of ultracool dwarfs obtained with the United Kingdom Infrared Telescope (UKIRT) using the Mauna Kea Observatory (MKO) $L'$ and $M'$ filters.  These data increase by factors of 0.6 and 1.6, respectively, the numbers of ultracool dwarfs for which MKO $L'$ and $M'$ measurements have been reported.  We examine the near-infrared colors and magnitudes of these dwarfs as functions of spectral type.  Using recently published trigonometric parallaxes, we show color--magnitude diagrams in the MKO $K$, $L'$, and $M'$ bandpasses, and we determine the bolometric luminosities and effective temperatures of ultracool dwarfs.  We compare these results with the predictions of recent atmospheric models that consider the effects of cloud sedimentation on the broadband spectra of these objects.  Finally, we consider the effects of nonequilibrium CO$\\leftrightarrow$CH$_4$ chemistry on the $M'$ luminosities of ultracool dwarfs and on direct searches for even cooler objects at wavelengths around $5~\\mu$m. ",
        "conclusions": "Our compilation of new and previously reported MKO $L'$ and $M'$ photometry has permitted us to characterize ultracool dwarfs comprehensively at wavelengths longward of the commonly used $J$, $H$, and $K$ bands.  We find that $K$--$L'$ increases monotonically with decreasing $T_{\\rm eff}$, but the nearly constant $T_{\\rm eff} \\approx 1450$~K of spectral types L7--T4 limits the utility of $K$--$L'$ as an indicator of spectral type.  Likewise, $L'$--$M'$ is nearly constant between types L6 and T3, indicating that the dramatic changes in the 1--$2.5~\\mu$m spectra of L--T transition dwarfs are not duplicated in their $L'$- and $M'$-band spectra.  This dichotomous behavior supports recent theories that the rapid migration, disruption, and/or thinning of condensate clouds at $T_{\\rm eff} \\lesssim 1400$~K occur at altitudes that are coincident with the regions of $z$- through $K$-band emission but are well below the $L'$-band and $M'$-band ``photospheres.''  The $L'$ and $M'$ luminosities of the early-T dwarfs do not exhibit the pronounced humps or inflections noted by others in the $I$ through $K$ bands, but insufficient data exist for types L6--T5 to assert that $M_{L'}$ and $M_{M'}$ are strictly monotonic within this range of types.  We used our $L'$ photometry, flux-calibrated $JHK$ spectra, and recently published trigonometric parallaxes to compute $L_{\\rm bol}$, BC$_K$, and $T_{\\rm eff}$ for ultracool dwarfs.  We find that BC$_K$ is a well behaved function of spectral type with a dispersion of $\\sim 0.1$~mag for types M6--T5.  BC$_K$ is significantly more scattered among the later T dwarfs, which may indicate the sensitivity of H$_2$ CIA in the $K$-band to varying surface gravity for $T_{\\rm eff} \\lesssim 1400$~K.  BC$_K$ is neither a monotonic nor single-valued function of $J$--$K$ because of the color reversal induced by the onset of CH$_4$ aborption at 2.2--$2.4~\\mu$m at spectral type L8.  BC$_K$ is a single-valued function of $K$--$L'$ except at $K$--$L' \\approx 1.6$, which corresponds to the L--T transition.  $T_{\\rm eff}$ declines steeply and monotonically for types M6--L7 and T4--T9, but is nearly constant at $\\sim 1450$~K for types L7--T4 with assumed ages of $\\sim 3$~Gyr.  Our photometry and bolometric calculations indicate that Kelu-1 (L3) and SDSS~J0423--0414 (T0) are probable binary systems.  We compute log$(L_{\\rm bol}/L_{\\odot}) = -5.73 \\pm 0.05$ and $T_{\\rm eff} = 600$--750~K for 2MASS~J0415--0935 (T9), making it the least luminous and coolest brown dwarf presently known.  We have compared the measured absolute magnitudes of L3--T9 dwarfs with those predicted by the precipitating-cloud models of \\citeauthor{mar02}\\ for varying surface gravities, $g$, and sedimentation efficiencies, $f_{\\rm sed}$. The models spanning $4.5 \\leq {\\rm log}~g \\leq 5.5$ and $f_{\\rm sed} = 3$, 5, and ``$\\infty$'' (no clouds) reproduce well the $M_K$ and $M_{L'}$ of all the dwarfs in our sample. The models indicate that the L3--T4.5 dwarfs generally have higher gravities (log~$g = 5.0$--5.5) than the T6--T9 dwarfs (log~$g = 4.5$--5.0).  The lower-gravity models underestimate $M_{M'}$ for the late-T dwarfs by 0.5--1~mag.  This overestimation of the $M'$ luminosity for $T_{\\rm eff} \\lesssim 1000$~K is attributed to absorption at 4.5--$4.9~\\mu$m by CO, which is not expected under the condition of thermochemical equilibrium assumed in the models.  The impact of nonequilibrium chemistry on the broadband near-infrared fluxes of cool brown dwarfs has only recently been appreciated.  Consequently, the effective-temperature limits of space-based $5~\\mu$m searches for infra-T dwarfs and EGPs, such as those planned with the recently-launched {\\it Spitzer Space Telescope}, will be somewhat higher than originally expected."
    },
    "0402/astro-ph0402196_arXiv.txt": {
        "abstract": " ",
        "introduction": " ",
        "conclusions": "The heating of the GC clouds is a long-standing problem. The discrepancy of the dust and gas temperatures and the apparently lack of ionized gas associated to the molecular clouds are usually invoked to rule out radiative heating mechanisms. However, the picture arising from the ISO observations is a complex medium where the molecular clouds are irradiated from the exterior by hot radiation arising from relatively hot and distant sources. The surfaces of the clouds are ionized and there is a warm (150 K) neutral gas component in the interface between the ionized and the molecular gas. The 30-40 K dust component should be associated with the 150 K gas component and both heated in PDRs. It is noteworthy that the discrepancy of gas and dust temperatures only rules out gas heating by collisions with hot dust but it does not rule out {\\it any} radiative heating mechanisms as photoelectric effect on the dust grains in a PDR. Indeed a gas temperature of 150 K and a dust temperature of $\\sim35$ K as measured in the GC clouds it is exactly what it is expected in a 10$^3$ cm$^{-3}$ and G$_0$=10$^3$ PDR (Hollenbach et al. 1991).  On the other hand, the extended distribution and high abundances of fragile molecules when iluminated by ultra-violet radiation (see for instance Martin-Pintado et al. 2001) and the large line-widths are the best probe of the presence of turbulence or low velocity shocks in the GC clouds. Shocks, turbulence and PDRs probably coexist in the GC. Unfortunately, it is still not clear  how much gas is heated in PDRs and how much is heated by other mechanisms. In any case, the new ISO data show that the importance of radiation in the heating of the GC clouds should be revisited and that a significant fraction of the warm gas in the GC clouds arise in standard PDRs.      \\vspace{5mm} {\\bf Acknowledgments:} The authors acknowledge the collaboration of  A. Fuente, P. de Vicente, T.L.Wilson and S. Huttemeister in some of the work presented in this contribution. N.J.R-F acknowledges fruitful discussions with J.R. Goicoechea and  M. Morris during the conference and support by a Marie Curie Fellowship of the European Community program``Improving Human Research Potential and the Socio-economic Knowledge base'' under contract number HPMF-CT-2002-01677. J. M-P acknowledges {\\it Ministerio de Ciencia y Tecnologia} grants ESP2002-01627 and AYA2002-10113-E."
    },
    "0402/astro-ph0402643_arXiv.txt": {
        "abstract": " ",
        "introduction": "The ultraviolet/optical continuum radiation of quasars, which is often called the Big Blue Bump, is thought to be from an accretion flow around a supermassive black hole. The radiative output of quasars is dominated by this component, and thus it is a crucial component to understand the accretion flow.  This radiation is often assumed to be an optically thick and thermal emission from an accretion disk.  However, continuum edge features, which are key predictions of many models for the atmospheric emission from accretion disks, have not been confirmed satisfactorily.\\cite{KB99,Bl01} \\ Identifying these features would prove the thermal nature of the emission in the first place, and thus it is of fundamental importance.  There have been intensive and unsuccessful searches for broadened Lyman edge discontinuities, starting with Ref.~\\citen{An89}.  Later some slope changes have been found around 1000\\AA\\ in some individual cases and composite spectra\\cite{Kr99,Zh97,Te02}.  This could be related to the opacity change at the Lyman edge which is severely smeared by relativistic effects.  However, in detail, even state-of-the-art atmosphere models are not in a good agreement with the observed feature.\\cite{Bl01} \\ In addition, a foreground absorption possibility certainly complicates the interpretation of the spectral features in this wavelength region.  The Balmer edge wavelengths can be more advantageous in these two respects: (1) the spectral smearing would be less effective in the context of accretion disk models, since this longer wavelength feature originates from much farther out in the gravitational potential well; (2) the Balmer edge wavelengths are less vulnerable to a foreground absorption possibility since the Balmer features are not resonant ones unlike the Lyman features. However, the emissions from outside the nucleus, i.e. the broad-line region (BLR) and outer regions, heavily contaminates this Balmer edge wavelength range. These are high-order Balmer emission lines and Balmer continuum in emission as well as broad FeII emission lines, collectively called the small blue bump.  There is a way to overcome this obstacle, however. Many normal quasars show a small polarization ($P \\sim 1$\\%), which can be used at least in some cases to eliminate any unwanted contamination from outside the nucleus, as we describe below. ",
        "conclusions": "The continuum-confined polarization seen in these objects is quite in contrast to Seyfert 2 galaxies where the continuum and the broad emission lines are {\\it both} polarized.  Based on this, the polarization is thought to occur due to scattering {\\it outside} the BLR.\\cite{AM85} \\ However, in our cases, the emission from the BLR is not polarized. Therefore, the polarization is considered to occur {\\it interior} to the BLR. Thus, the polarized flux is likely to show the intrinsic spectral shape of the emission from interior to the BLR, scraping off all the other contaminating emission from outside.  The cause of the polarization is not still clear, but probably it involves some form of electron scattering. Dust scattering is not likely, since the region which is causing the polarization is thought to be interior to the BLR and thus within the dust sublimation radius.  The scattering region could be the atmosphere of the putative accretion disk around the central black hole (though see more discussions in Paper~I). Alternatively, it could be a diffuse optically-thin region surrounding the emitter of the Big Blue Bump.  In either of these simplest cases, the Balmer edge feature seen in absorption in the polarized flux is intrinsic to the Big Blue Bump emission. In this case, the feature indicates that the emission is indeed thermal, and the emitter is optically thick.  However, some alternative explanations for the observed polarized flux are not ruled out, as discussed in Paper I. Briefly, the feature might have been imprinted in some foreground region or in the scattering/polarizing region itself. However, a simple foreground absorption is unlikely due to the general lack of a huge, corresponding Lyman edge absorption in the total flux.  If the Balmer edge feature were to be imprinted in-situ when the polarized flux is formed in the scattering region outside the Big Blue Bump emitter, then the distinction from the case of the atmospheric feature might be just semantic, since the scattering region should be optically thick, and should have a large velocity dispersion to produce a large broadening in the edge feature.  In the re-observation of Ton 202 in May 2003, which is one year apart from the observation shown in the top panels of Fig.\\ref{fig-5obj}, we did not detect such a dramatic feature. The details are presented in Paper II, but this is apparently due to polarization variability.  This is not unexpected, since, as we discussed above, the polarization is thought to originate from a quite compact region, at least comparable to or smaller than the size scale of the BLR.   In conclusion , the edge feature is most simply interpreted as an intrinsic feature of the Big Blue Bump emission, showing its thermal and optically thick nature. However, we cannot rule out some alternative interpretations, and we need more data and modeling to confirm that our interpretation is correct."
    },
    "0402/astro-ph0402556_arXiv.txt": {
        "abstract": "{There is observational evidence that nuclear winds and X-rays can  heavily influence the physical conditions and chemical abundances of  molecular gas in the circumnuclear disks (CND) of Active Galactic Nuclei  (AGN). In this paper we probe the chemical status of molecular gas in the CND  of \\object{NGC~1068}, a prototypical  Seyfert 2 galaxy. Precedent claims that the chemistry of molecular gas in the nucleus of  \\object{NGC~1068} is {\\sl abnormal} by galactic standards were based on the high  HCN/CO luminosity ratio measured in the CND. Results from new observations obtained in this survey have { served to derive}  abundances of molecular species such as SiO, CN, HCO$^+$,  HOC$^+$, H$^{13}$CO$^+$ and HCO. These  estimates are complemented by a re-evaluation of molecular abundances for  HCN, CS and CO, based on previously published single-dish and interferometer  observations of \\object{NGC~1068}. We report on the first detection of   SiO emission in the CND of \\object{NGC~1068}. The estimated large  abundance of SiO in the CND, X(SiO)$\\sim$(5-10)$\\times$10$^{-9}$, cannot be  attributed to shocks related to star formation, as there is little evidence  of a recent starburst in the nucleus of \\object{NGC~1068}. Alternatively, we  propose that silicon chemistry is driven by intense X-ray processing of  molecular gas. We also report on the first extragalactic detection of the  reactive ion HOC$^+$. Most remarkably, the estimated HCO$^+$/HOC$^+$  abundance ratio in the nucleus of \\object{NGC~1068}, $\\sim$30--80, is the smallest ever measured in molecular gas. The abundances derived for all  molecules that have been the subject of this survey are compared with the  predictions of models invoking either oxygen-depletion or X-ray chemistry in  molecular gas. Our conclusions favour an overall scenario where the CND of  \\object{NGC~1068} has become a giant X-ray Dominated Region (XDR). ",
        "introduction": "Active Galactic Nuclei (AGN) are able to inject vast amounts of energy into their { host galaxies}, carried by strong radiation fields and rapidly moving jets. It is predictable that AGN should have a disruptive influence on the gas reservoir near their central engines. There is multi-wavelength observational evidence that the general properties of neutral interstellar matter in AGN differ from { those} of quiescent star-forming disks and starburst galaxies (Genzel et al. \\cite{genz98}; Laurent et al. \\cite{laur00}). In particular, molecular gas close to the central engines of active galaxies can be exposed to a strong X-ray irradiation. While the accretion disks of AGN are strong UV emitters, the bulk of the UV flux can be attenuated by neutral gas column densities of only N(H)$\\sim$10$^{21}$cm$^{-2}$. Hard X-ray photons (2--10~keV) can penetrate neutral gas column densities out to N(H)$\\sim$10$^{23}$-10$^{24}$cm$^{-2}$, however. Therefore, X-ray dominated regions (XDR) could become the dominant sources of emission for molecular gas in the harsh environment of circumnuclear disks (CND) of AGN, as originally argued by Maloney et al. (\\cite{mall96}).  First observational evidence that the physical and chemical properties of molecular gas in the CND of AGN depart from `normality' came from the single-dish and interferometer observations of HCN and CO emission in \\object{NGC~1068} (Tacconi et al. \\cite{tacc94}; Sternberg et al. \\cite{ster94}). This prototypical Seyfert 2 galaxy hosts a circumnuclear starburst ring of $\\sim$2.5-3\\,kpc--diameter (see Fig.~\\ref{PosFig}); the ring delimits a 2.3~kpc stellar bar detected by Scoville et al. (\\cite{scov88}) in the NIR. The strong emission detected in the 1--0 and 2--1 CO lines coming from the starburst ring corroborates that massive star formation is fed by a significant gas reservoir (Planesas et al. \\cite{plan89}, \\cite{plan91}; Helfer \\& Blitz \\cite{helf95}; Schinnerer et al.  \\cite{schi00}). Significant CO emission arises also from a 200\\,pc CND of M(H$_2$)$\\sim$5$\\times$10$^{7}$M$_{\\sun}$ { (inferred using a N(H$_2$)/I(CO) conversion factor of 2.2$\\times$10$^{20}$~cm$^{-2}$~(K~km~s$^{-1}$)$^{-1}$, from Solomon \\& Barrett \\cite{solo91})}.  The CND, partly resolved into two  knots, surrounds the position of the active nucleus identified as the compact radio-source S1 in the map of Gallimore et al. (\\cite{gall96a}). Most remarkably, the CND is prominent in HCN emission (Tacconi et al. \\cite{tacc94}). According to the analysis of Sternberg et al. (\\cite{ster94}), the high HCN/CO intensity ratio measured by Tacconi et al. (\\cite{tacc94}) ($\\sim$1--10) leads to an abnormally high HCN/CO abundance ratio in the nucleus of \\object{NGC~1068}: N(HCN)/N(CO)$\\sim$a few~10$^{-3}$-10$^{-2}$, i.e., the highest ratio ever found in the centre of any galaxy.  Different explanations have been advanced to quantify the possible link between the anomalous HCN chemistry and the presence of an active nucleus in \\object{NGC~1068}. The selective depletion of gas-phase oxygen in the dense molecular clouds would explain the high HCN-to-CO abundance ratio (Sternberg et al. \\cite{ster94}; Shalabiea \\&  Greenberg \\cite{shal96}). The same oxygen depletion scheme predicts a lower-than-normal abundance of all oxygen-bearing species. Alternatively, an increased X-ray ionization of molecular clouds near the AGN could enhance the abundance of HCN (Lepp \\& Dalgarno \\cite{lepp96}). Furthermore, X-rays could evaporate small ($\\sim$10~\\AA) silicate grains, increasing the fraction in gas phase of all refractory elements and subsequently enhancing the abundance of some molecules (e.g., SiO) in X-ray irradiated molecular gas (Voit \\cite{voit91}; Mart\\'{\\i}n-Pintado et al. \\cite{mart00}). While the aforementioned scenarios { succeed} to reproduce the measured enhancement of HCN relative to CO in \\object{NGC~1068}, their predictions about the abundances of other molecular species differ significantly. The lack of tight observational constraints for these models, prompted by the first mm-observations made in \\object{NGC~1068}, has hampered thus far the choice of an optimum scenario, however.   In this paper we discuss the results of a molecular survey made in \\object{NGC~1068} with the IRAM 30m mm-telescope. \\object{NGC~1068} is the optimum target to { quantify} the feed-back of  activity and star formation on the chemistry of molecular gas. Furthermore,  the spatial resolution of the 30m telescope is well suited to discern between the emission coming from the star forming ring and that coming from the CND.  We discuss the results obtained from new mm-observations of 6 molecular species. The list includes:  SiO(v=0, J=2--1 and J=3--2), HCO(J=3/2--1/2, F=2--1),  H$^{13}$CO$^{+}$(J=1--0), HCO$^{+}$(J=1--0), HOC$^{+}$(J=1--0) and  CN(N=2--1). For { comparison purposes}, we include in our analysis the  results from previous single-dish and interferometer observations of CO  (J=1--0 and 2--1 from Schinnerer et al. \\cite{schi00}; J=4--3 from Tacconi et  al. \\cite{tacc94}), HCN (J=1--0 and 4--3) (Tacconi et al. \\cite{tacc94}) and  CS (J=2--1) (Tacconi et al. \\cite{tacc97}). This data { base} has served for estimating the abundances of eight molecular species in the CND of  \\object{NGC~1068} using LVG model calculations. The inferred abundances are  compared with the predictions of models invoking either oxygen-depletion or  X-ray chemistry in molecular gas.  We present in Section~\\ref{sec2} the 30m observations made for this survey as  well as the data compiled from previous works on \\object{NGC~1068}.  Section~\\ref{sec3} presents the results obtained from our SiO study and their implications for the CND chemistry. Section~\\ref{sec4} is devoted to discuss  the chemistry of the HOC$^{+}$/HCO$^{+}$ active ions. The molecular gas  inventory of the CND is globally presented and discussed in Section~\\ref{sec5}.  We discuss in Section~\\ref{sec6} the interpretation of these results in the  framework of different chemistry models and summarize the main conclusions of  this work in Section~\\ref{sec7}. ",
        "conclusions": ""
    },
    "0402/astro-ph0402283_arXiv.txt": {
        "abstract": "This guide presents, in a uniform manner, all the information collected by the IUE satellite on normal galaxies.  It contains information on 274 galaxies and it supersedes the previous IUE guide to normal galaxies (Longo \\& Capaccioli 1992). The data shown here are restricted to galaxies defined as ``normal'' by the observer, and entered as such in the IUE data file header. The information is also restricted to low-dispersion spectra obtained through the large apertures of IUE. For the first time, we provide spectral information from {\\bf well-defined} and identifiable locations in the target galaxies. These are mostly located close to the photocenter of each object, although there are few exceptions. Each representative spectrum of a galaxy consists of  a short-wave (SW) and a long-wave (LW) IUE low-dispersion spectrum (where available) combined into a single spectrum covering the wavelength range 1150\\AA\\, to 3350\\AA\\,. We selected the two spectra to be combined so as to be, preferably, the deepest exposures available in the INES archive. Each representative spectrum is accompanied by two images of the galaxy, on which the locations of the SW entrance aperture and the LW entrance aperture are marked. ",
        "introduction": "The International Ultraviolet Explorer (IUE) Satellite project was a joint effort between NASA, ESA and the PPARC. The IUE spacecraft and instruments were operated in a Guest Observer mode to allow ultraviolet spectrophotometry at two resolutions in the wavelength range from  1150\\AA\\, to 3200\\AA\\ : low resolution $\\frac{\\lambda}{\\Delta\\lambda}$=300 ($\\sim$1,000 km/sec.) and a high resolution mode $\\frac{\\lambda}{\\Delta\\lambda}$=10,000 ($\\sim$19 km/sec.). The IUE spacecraft, its scientific instruments as well as the data acquisition and reduction procedures, have been described in ``Exploring the Universe with the IUE Satellite\", Part I, Part VI and Part VII (Astrophysics and Space Sciences Library volume 129, Y. Kondo, Editor-in-Chief, Kluwer Acad. Publ. Co.) and references therein. A more recent overview of the IUE project is given in the conference proceedings of the last IUE conference ``Ultraviolet Astrophysics beyond the IUE Final Archive''  (ESA SP-413, 1998, Eds. W.Wamsteker and R. Gonzalez Riestra) and in ``IUE Spacecraft Operations Final Report'' (ESA SP-1215, 1997, A. Perez Calpena \\& J.Pepoy).  Additional information on the IUE project can also be found at URL: http://www.vilspa.esa.es/iue/iue.html.  From the very beginning of the project (launched on 26 January 1978), it was expected that the archival value of the data obtained with IUE would be very high. This expectation has been borne out fully after 18.6 years of orbital operations (on 30 September 1996 the science operations with the IUE spacecraft were stopped). The average IUE Archive data retrieval rate, during the operational phase of the project, has been some 61,000 spectra per year. This compares to a new data access rate of 5,500 spectra per year. Considering that the demand for observing time continued to exceed the available time by a factor of three until the end of science operations, it is clear that the IUE Archive remains an important source of data. The IUE ULDA/USSP (Uniform Low Dispersion Archive/ULDA Support Software Package) was developed by ESA in the mid-eighties (Wamsteker et al., 1989, Astron. \\& Astrophys. Suppl. Ser., {\\bf 79}, pg. 1-10) as the first astronomical archive with direct access for users on a world-wide basis. Over the ten years that the ULDA has been supporting IUE data retrieval, it has driven more than 50\\% of all IUE Archive usage. The quantity of data in the IUE Archive is sufficiently large that it is not necessarily simple to address the data efficiently in the context of an astrophysical problem, even  though access to the data is extremely easy. Therefore, the series of INES Access Guides is intended to facilitate the use of the IUE Archive for scientists with a specific astrophysical problem in mind.  The extremely good reception of the distributed archive model by the scientific community has led to the decision to develop the Final Archive server with a similar philosophy, in the form of INES (IUE Newly Extracted Spectra). The INES system is a complete system design, with \\begin{itemize} \\item improvements to the data \\item a structure to facilitate the direct application to scientific analysis, and \\item an integrated and distributed data retrieval system. \\end{itemize}  Detailed information on the INES System and its data content has been published in a series of papers ``The INES System''. The specific references are: Rodriguez-Pascual \\etal 1999, Astron. \\& Astrophys. Suppl. Ser., {\\bf 139}, pg.183-198; Cassatella \\etal 2000, Astron. \\& Astrophys. Suppl. Ser., {\\bf 141}, pg.331-342; Gonzalez-Riestra \\etal 2000 Astron. \\& Astrophys. Suppl. Ser, {\\bf 141}, pg.343-356; Wamsteker \\etal 2000, Astroph. \\& Space Sci., in press. The INES Users Guide, also collecting these publications, has been published by ESA Publications Division in the INES Newsletter (March 2000).  At  the time of this writing 19 National Hosts have functional installations  of the INES system  Version 2.0 (see below), giving to the end user, direct access to all 104,000 IUE spectra, with the Principal Center (LAEFF) at URL  http://ines.vilspa.esa.es/ines/ and its Mirror Site for North America at the Canadian Astronomy Data Centre (CADC) http://ines.hia.nrc.ca in Victoria. The ESA involvement in the IUE Project activities will come to an end in 2000, and, from that moment onward, the INES system will become a part of the astronomical heritage of the IUE project. More National Hosts are forseen to come on-line in the future. The series of INES Access Guides is forseen to continue and consists of a number of subject-oriented books, for which a specialist in the field has been invited to take the scientific responsibility. INES  Access Guide No. 2 treats the data for Normal Galaxies and has been compiled by  Liliana Formiggini and Noah Brosch of the Wise Observatory and the School of Physics and Astronomy, Tel Aviv University. This volume supersedes the ULDA Guide No. 3. Further volumes of INES Access Guides  will be published whenever the necessary data compilation has been completed by the authors. The list of previously published ULDA \\& INES Access Guides is given below, as well as the INES Guides currently in preparation. For details of the access to INES through the National Hosts we refer to the information supplied at http://ines.vilspa.esa.es/ines/, or recommend contacting the INES Helpdesk at LAEFF at VILSPA, Madrid, Spain ineshelp@iuearc.vilspa.esa.es. Other queries about the data or any specific questions about data content in relation to the INES system should also be directed there or to the National Host Institutes.  \\hspace{10cm}    Willem Wamsteker   \\subsection{Previously issued IUE-ULDA Access Guides:} \\begin{description} \\item No. 1  ESA SP-1114    C. la Dous {\\bf Dwarf Novae and Nova-like Stars}. \\item No. 2  ESA SP-1134    M. Festou {\\bf Comets}. \\item No. 3  ESA SP-1146    G. Longo, M. Capaccioli {\\bf Normal Galaxies}. \\item No. 4  ESA SP-1153 (Vol. A \\& B)  T.J.-L. Courvoisier, S. Paltani {\\bf Active Galactic Nuclei}. \\item No. 5  ESA SP-1181 (Vol. I \\& II) C. la Dous, A. Gimenez {\\bf Chromospherically Active Binary Stars}. \\item No. 6  ESA SP-1189    E. Cappellaro, M. Turatto, J. Fernley {\\bf Supernovae} \\item No. 7  ESA SP-1203     M. Franchini, C. Morossi, M.L. Malagnini {\\bf K Stars} \\item No. 8  ESA SP-1205    A.I. Gomez de Castro, M. Franqueira {\\bf T Tauri Stars} \\end{description}  \\subsection{INES Guides:} \\begin{description} \\item No. 1  ESA SP-1237    A.I. Gomez de Castro, A. Robles {\\bf Herbig-Haro Objects} \\item No. 2  ESA SP-1239 L. Formiggini,  N. Brosch {\\bf Normal Galaxies} supersedes ULDA Guide No.3 \\end{description}  \\subsection{INES Guides in preparation:} \\begin{description} \\item No. 3    A. Niedzielski  {\\bf Wolf-Rayet Stars} \\item No. 4    M. Festou    {\\bf Comets} Supersedes ULDA Guide \\# 2 \\item No. 5   A. Cassatella \\& R. Gonzalez-Riestra  {\\bf Novae} \\end{description}  \\subsection{INES National Hosts (as of 18 February 2000):} \\begin{tabular}{c|c} Country         &               Institution \\\\ \\hline Argentina     &    Facultad de Ciencias  Astronomicas y Geofisicas, \\\\ &               Buenos Aires \\\\ Austria      &      Kuffner Observatory, Vienna \\\\ Belgium      &      Royal Observatory of  Belgium, Ukkel \\\\ Brazil       &      Instituto Astronomico e  Geofisico, Sao Paulo \\\\ Canada     &        Canadian Astronomy Data Center, Victoria \\\\ &           (Principal   Centre Mirror Site) \\\\ China       &       Center for Astrophysics, \\\\ &             Hefei (local access only) \\\\ ESO          &    European Southern  Observatory, Garching \\\\ &            (local access only) \\\\ Israel       &      Wise Observatory, Tel Aviv University, Tel Aviv \\\\ France      &       Centre de Donees Astronomiques, Strasbourg \\\\ Italy       &       Astronomical Observatory,     Trieste \\\\ Japan        &      National Astronomical  Observatory, Tokyo \\\\ Korea       &       Chungbuk National University, Cheongju \\\\ The Netherlands  &  Astronomical Institute,  Utrecht \\\\ Poland       &      Center for Astronomy, Torun \\\\ Russia      &    Institute of Astronomy of  the Russian Academy of \\\\ &              Sciences, Moscow \\\\ Spain      &        LAEFF, Madrid (Principal Centre) \\\\ Sweden      &      Astronomical Observatory Uppsala \\\\ &       (serving the Nordic  Countries) \\\\ Taiwan       &      Institute of Physics and Astronomy, Chung-Li \\\\ United Kingdom  &   Rutherford Appleton Laboratory, Didcot \\\\ United States   &   Space Telescope Science   Institute, Baltimore \\end{tabular}   \\newpage ",
        "conclusions": "We present here a compilation of all galaxies classified as ``normal'', {\\it i.e.} not containing an AGN, which were observed by the IUE satellite. This is an updated and enlarged version of the previous IUE ULDA Guide to Normal Galaxies, produced by Longo \\& Capaccioli, and supersedes it. We took advantage of the final processing of IUE observations within the INES archive to extract all the observations of galaxies classified as normal and we  present their UV spectra after matching the fluxes between the SW and LW segments that were observed toward the same location in the galaxy.  We believe that although most IUE observations of normal galaxies were only of their central regions, as defined by the optical light distributions, this compilation will be of some use in understanding the UV light produced by the different stellar populations in objects lacking a central, compact, massive object."
    },
    "0402/astro-ph0402410_arXiv.txt": {
        "abstract": "We perform N-body simulations of the dynamical evolution of a dwarf galaxy falling into the Milky Way  galaxy in order to understand the formation scenario of the peculiar globular cluster $\\omega$ Centauri. We use  self-consistent models of the bulge and the disc of the Milky Way, as well as of the dwarf galaxy, and explore a range of dwarf models with different density distributions. Namely, we use King  (1966) and Hernquist (1990) density profiles to model the density distribution in the dwarf. The central region of our King model has a density profile approximately $\\propto r^{-2}$, while that of the Hernquist model is $\\propto r^{-1}$. The difference in the dwarf's density distributions leads to distinct evolutionary scenarios. The King model dwarf loses its mass exponentially as a function of apocentric distance, with the mass-loss rate depending on the initial mass and size of the dwarf. Regardless of the initial mass and size, the King model dwarf remains more massive than $10^8$ \\msun after a few  Gyr of evolution. The Hernquist model dwarf experiences an accelerated mass loss, and the mass of the remnant falls below $10^8$ \\msun within a few Gyr. By exploring an appropriate set of parameters, we find a Hernquist model that can attain the mass and orbital characteristics of $\\omega$ Cen after a few Gyr. ",
        "introduction": "\\label{sec:Introduction}  The accretion and subsequent tidal disruption of dwarf galaxies are important agents in the formation and evolution of the Milky Way (Searle \\& Zinn 1978). Such accretion events would presumably lead to formation of tidal streams in the halo of the Milky Way associated with the disrupted dwarf galaxies. The discovery of the tidal stream related to the Sagittarius dwarf galaxy (Ibata et al. 2001) indeed supports the idea that halos of galaxies can have a number of streams maintaining their coherence during a few Gyrs.  There is another class of halo objects which might be associated with past accretion events. During the accretion of nucleated dwarf elliptical galaxies, the compact stellar nuclei may survive the disruptive event and continue orbiting until the present time. Freeman \\& Bland-Hawthorn (2002) even speculate that the globular clusters in the Milky Way could be the stripped relics of ancient protogalactic stellar systems. To identify such objects would be of great importance in understanding the nature of such building blocks of galaxies. In this paper we focus on one object which may be the relic of a past accretion event - the globular cluster $\\omega$ Centauri. $\\omega$ Centauri is one of the most peculiar globular clusters in the Milky Way: with a mass of $5\\times10^6$\\msun, it is the most massive globular cluster. It has an unusual flatness with an ellipticity of about 0.12 associated with fast internal rotation  \\citep{free2001}. It orbits in a retrograde direction relative to the disc with an apocentric radius of about 6 kpc, a pericentric radius of 1.2 kpc, and a maximum height above the plane of 1 kpc \\citep{dine1999}.  What really distinguishes $\\omega$ Cen from other Galactic globular clusters, is its chemical abundance pattern. It is the only known globular cluster that has well-defined signatures of self-enrichment. Thus, according to Smith et al. 2000 and Vanture et al. 2002, the abundance pattern observed within $\\omega$ Cen can be understood as a combination of Type II supernovae and, more importantly, AGB-star enrichment during a period of about 3 Gyr. Gnedin et al. (2002) have demonstrated that $\\omega$ Cen on its present orbit and with its present mass can not retain heavy elements dispersed in the AGB phase of stellar evolution. The passages through the disk would sweep out the intracluster gas and the cluster, once established on this orbit, would not be able to produce stars with enhanced s-process elements. Thus, the cluster should have evolved chemically before settling on its present orbit, or, in other words, it was accreted by the Galaxy. Another possibility to explain the chemical peculiarities of $\\omega$ Cen is the merger of two or more chemically distinct systems. Vanture et al. (2002) conclude however, that in this scenario, it would be difficult to explain an increase in the of ratio of heavy s-process elements to iron with metallicity.  A plausible explanation is that   $\\omega$ Cen is the nucleus of a dwarf galaxy captured and disrupted by the the Milky Way (e.g., Freeman 1993). In this scenario, a dwarf sinks to the center of the Galaxy due to dynamical friction,  simultaneously losing mass in the Galactic tidal field. All of the dwarf's mass is eventually stripped but the nucleus, with a mass of $5\\times10^6$ \\msun, is left on the current orbit of $\\omega$ Cen. The current angular momentum of $\\omega$ Cen is thus inherited from the dwarf galaxy, on a retrograde orbit. The complex chemical composition and extended star formation history of $\\omega$ Cen can be explained in this picture by the chemical evolution of the nucleus in a deep potential well of a dwarf, such that the nucleus would be able to retain stellar ejecta during subsequent activity of supernovae and AGB stars.  Even though a capture scenario is seen as a plausible explanation for the peculiar properties of $\\omega$ Cen, it is unclear whether it is possible to realize an orbit and mass evolution of the dwarf that can reproduce the current parameters of $\\omega$ Cen. Zhao (2002) studied the orbital decay of a dwarf that may lead to the current position of $\\omega$ Cen by launching a dwarf on an orbit that started 50 kpc away from the Galactic center. Using a semi-analytical model, he concluded that a progenitor of $\\omega$ Cen can not decay to its present orbit. He found that strong tidal shocks quickly reduce the mass of the dwarf so that dynamical friction becomes too weak to drag the remnant to the inner regions of the Galaxy. Zhao (2002) concluded that the only possibility to explain $\\omega$ Cen - phenomenon is that its progenitor was born only $15$~kpc away from the Galactic center.  Zhao's analysis, however,  was simplified in a number of ways. He assumed that the mass outside the tidal radius is instantaneously stripped,  and the contribution from the bulge and disc components to the dynamical drag was not taken into account. In this study, we perform a series of numerical simulations, using self-consistent bulge and disc models for the Milky Way, and self-consistent dwarf models that have different density profiles. As a result, we find a dwarf-capture scenario that successfully produces an $\\omega$ Cen-like object from a normal dwarf galaxy. This result is also reported in more detail in a separate paper \\citep{korc2003}. In this paper, we explore the model parameters and focus on the physical processes of the dynamical evolution of such infalling dwarfs making this study beneficial in understanding the properties of merging parent-satellite galaxy systems.  The numerical realization of this scenario is a complicated problem because of the large range of masses and sizes of the interacting components, and due to the drastic change of the mass of the sinking dwarf. We explain our numerical methods in Sect.~\\ref{sec:Numerical_Methods}, where we describe the Galaxy and the dwarf models. The results of the simulations are presented in Sec.~\\ref{sec:Results}. Some analytical considerations for our results are given in Sect. 4. Finally we formulate our conclusions in Sect.~\\ref{sec:Conclusions}. ",
        "conclusions": "\\label{sec:Conclusions}  We have studied the dynamical evolution of a falling dwarf galaxy, which in turn is influenced by the strong gravitational shocks from the bulge and disc of the host galaxy. We focused our study on the mass-loss history, and on the orbital migration of the dwarf. The dwarf galaxy sinks toward the center of the host galaxy due to dynamical friction, which depends on the dwarf's mass. Simultaneously, the dwarf is losing mass due to strong bulge and disc gravitational shocks. As a result, the mass-loss history and orbital decay are closely connected and depend in turn on the internal density distribution of the dwarf. We have examined two different dwarf models, namely a highly-concentrated King model and  the Hernquist model. The Hernquist model has a $r^{-1}$ density cusp while the King model has a density profile nearly proportional to $r^{-2}$. This difference leads to qualitatively different evolutionary courses for the two models. Regardless of the mass and the size, all the King models stop orbital sinking in a few Gyr, once the remnant mass has reached $\\sim 10^8$ \\msun. The central regions of the King model dwarf remain gravitationally bound, and the remnant continues to orbit for a long time without considerable mass loss. Therefore, if the initial density distribution of the dwarf galaxy is close to isothermal, its remnant can not have a mass much smaller than $10^8$ \\msun. On the other hand, the Hernquist model dwarfs keep losing mass even after dynamical friction becomes ineffective. As a consequence, the remnant can be stripped down to a small object. If there is a compact gravitationally  bound nucleus in the center of the dwarf with a mass of $\\la 10^7$ \\msun and will the size of a globular cluster, it will survive for few Gyr on the present orbit of $\\omega$ Centauri (e.g., Fall \\& Zhang 2001). Our results support the capture scenario for the origin of the globular cluster $\\omega$ Centauri."
    },
    "0402/astro-ph0402626_arXiv.txt": {
        "abstract": "{ We estimate the group shape of the recently compiled Pecolation-Inferred Galaxy Group (2PIGG) catalogue. Using a set of integral equations we invert the projected distribution of group axial ratios and recover the corresponding intrinsic distribution under the assumption that groups are pure spheroids. In agreement with the analysis of the  UZC-SSRS2 group sample (Plionis, Basilakos \\& Tovmassian 2004) we find that groups are extremely elongated prolate-like systems. We also find an interesting trend between shape and group richness with poorer groups being significantly flatter than richer ones. ",
        "introduction": "Groups of galaxies are the lowest level cosmic structures, after galaxies themselves, in the hierarchy  that leads to the largest  virialized structures, the clusters of galaxies. It appears that most galaxies are found in groups and they are therefore extremely important in our attempts to understand the cosmic structure formation processes. Since virialization will tend to sphericalize initial anisotropic distributions of matter, the shape of different cosmic structures is a possible indication of their evolutionary stage.  Theoretical and observational works have shown that the cosmic structures (clusters and superclusters) are dominated by prolate like shapes (cf. Carter \\& Metcalfe 1980; Plionis, Barrow \\& Frenk 1991; Cooray 1999; Basilakos, Plionis \\& Maddox 2000; Zeldovich, Einasto \\& Shandarin 1982; de Lapparent, Geller \\& Huchra 1991; Plionis, Jing \\& Valdarnini 1992; Jaaniste et al. 1998; Sathyaprakash et al. 1998; Valdarnini, Ghizzardi \\& Bonometto 1999; Basilakos, Plionis \\& Rowan-Robinson 2001). In the case of Groups of galaxies a recent study of the UGC-SSRS2 group catalogue (Plionis, Basilakos \\& Tovmassian 2004) have shown that they are extremely flat systems. Furthermore, their intrinsic shape appears to be that of a prolate-like spheroid (see also Oleak et al. 1995 and references therein).  In this letter we use the recently constructed 2PIGG group catalogue (Eke et al. 2004) which is based on the Two-Degree Field Galaxy Redshift Survey (2dFGRS) and contains $7020$ groups with four or more galaxy members to estimate their projected and intrinsic shape as a function of their galaxy membership. ",
        "conclusions": "We have measured the projected axial ratio distribution of the 2PIGG groups of galaxies within a roughly volume-limited region ($cz\\le 30000$ km/sec). Assuming that groups constitute a homogeneous spheroidal population, we numerically invert the projected axial ratio distribution to obtain the corresponding intrinsic one. The only acceptable model is that of prolate spheroids which is in excellent agreement with the analysis of the UGC-SSRS2 sample of poor groups (Plionis et al. 2004).  There is an obvious group richness-flatness relation, seen in both projected and intrinsic axial ratio distributions, with poorer groups being also the flatter. For example, the peak of the projected axial ratio distribution shifts from $q\\sim$ 0.22 for groups with only 4 members to $\\sim$0.55 for those with more than 19 members. The increase of the group sphericity with richness, which extends also to clusters, could be explained as an indication of a higher degree of virialization, which is expected to be more rapid in more massive systems."
    },
    "0402/astro-ph0402189_arXiv.txt": {
        "abstract": "{We report follow-up observations of a new white dwarf/red dwarf binary HS\\,2237+8154, identified as a blue variable star from the Hamburg Quasar Survey.  Ellipsoidal modulation observed in the $R$ band as well as the $\\Ha$ radial velocity variations measured from time-resolved spectroscopy determine the orbital period to be $\\Porb=178.10\\pm0.08$\\,min. The optical spectrum of HS\\,2237+8154 is well described by a combination of a $\\Teff=11\\,500\\pm1500$\\,K white dwarf (assuming $\\log g=8$) and a dM3.5$\\pm$0.5 secondary star. The distance implied from the flux scaling factors of both stellar components is $d=105\\pm25$\\,pc. Combining the constraints obtained from the radial velocity of the secondary and from the ellipsoidal modulation, we derive a binary inclination of $i\\simeq50^{\\circ}-70^{\\circ}$ and stellar masses of $\\Mwd=0.47-0.67$\\,\\Msun\\ and $\\Msec=0.2-0.4$\\,\\Msun.  All observations imply that the secondary star must be nearly Roche-lobe filling. Consequently, HS\\,2237+8154 may be either a pre-cataclysmic variable close to the start of mass transfer, or~--~considering its orbital period~--~a cataclysmic variable that terminated mass transfer and entered the period gap, or a hibernating nova. ",
        "introduction": "The currently known population of cataclysmic variables (CVs) comprises more than 1000 systems, with period measurements available for $\\simeq500$ of them \\citep{downesetal01-1, kubeetal03-1, ritter+kolb03-1}. Compared to these numbers, only very few CV progenitors have been identified so far. \\citet{schreiber+gaensicke03-1} have recently analysed 30 well-observed pre-CVs and found that the sample of known pre-CVs is strongly biased towards systems containing hot (young) white dwarfs and late type (low-mass) companion stars. Testing and improving our understanding of the evolution of pre-CVs through the common envelope phase and of the angular momentum loss mechisms that subsequently brings them into a semidetached configuration will enormously benefit from a larger and unbiased sample of such stars.  We have identified a detached white dwarf/red dwarf binary (Fig.\\,\\ref{f-fc}) in the course of our studies of the stellar content of the Hamburg Quasar Survey \\citep[e.g.][]{gaensickeetal00-2, nogamietal00-1, szkodyetal01-1, gaensickeetal02-2, araujo-betancoretal03-2}. Our follow-up observations show that HS\\,2237+8154 may belong to the large population of ``cold \\& old'' pre-CVs predicted by \\citet{schreiber+gaensicke03-1}, however, it may as well be a CV in a peculiar phase of its evolution. ",
        "conclusions": "We have identified a detached white/red dwarf binary with an orbital period of $\\Porb=178.10\\pm0.08$\\,min. Our spectroscopic and photometric observations constrain the binary inclination to be $i\\simeq50^{\\circ}-70^{\\circ}$ and the stellar masses to be $\\Mwd=0.47-0.67$\\,\\Msun\\ and $\\Msec=0.2-0.4$\\,\\Msun. Modelling the optical spectrum of HS\\,2237+8154 we find $\\Teff=11\\,500\\pm1500$\\,K for the white dwarf and a spectral type dM3.5$\\pm$0.5 for the secondary. The distance of the system implied by the flux scaling factors of both stellar components is $d=105\\pm25$\\,pc.  The observations strongly suggest that the secondary star in HS\\,2237+8154 is nearly filling its Roche lobe, and we conclude that the system is either a pre-CV just about to evolve into a semi-detached configuration, a CV that evolved into a detached configuration and entered the period gap, or a hibernating nova. All three possible scenarios are extremely interesting in the context of CV evolution, and we encourage future work to refine the binary parameter of HS\\,2237+8154."
    },
    "0402/astro-ph0402140_arXiv.txt": {
        "abstract": "The cosmological star formation rate in the combined Chandra Deep Fields North and South is derived from our X-Ray Luminosity Function for Galaxies in these Deep Fields. Mild evolution is seen up to redshift order unity with SFR $\\sim (1 + z)^{2.7}$. This is the first directly observed {\\it normal star-forming galaxy} X-ray luminosity function (XLF) at cosmologically interesting redshifts (z$>$0).  This provides the most direct measure yet of the X-ray derived cosmic star-formation history of the Universe.   We make use of Bayesian statistical methods to classify the galaxies and the two types of AGN, finding the most useful discriminators to be the X-ray luminosity, X-ray hardness ratio, and X-ray to optical flux ratio. There is some residual AGN contamination in the sample at the bright end of the luminosity function. Incompleteness slightly flattens the XLF at the faint end of the luminosity function. \\\\  The XLF has a lognormal distribution and agrees well with the radio and infrared luminosity functions. However, the XLF does not agree with the Schechter luminosity function for the H$\\alpha$LF indicating that, as discussed in the text, additional and different physical processes may be involved in the establishment of the lognormal form of the XLF. \\\\  The agreement of our star formation history points with the other star formation determinations in different wavebands (IR, Radio, H$\\alpha$) gives an interesting constraint on the IMF.  The X-ray emission in the Chandra band is most likely due to binary stars although X-ray emission from non-stellar sources (e.g., intermediate-mass black holes and/or low-luminosity AGN) remain a possibility.  Under the assumption it is binary stars the overall consistency and correlations between single star effects and binary star effects indicate that not only is the one parameter IMF(M) constant but also the bivariate IMF($M_1, M_2$) must be constant at least at the high mass end.  Another way to put this, quite simply, is that X-ray observations may be measuring directly the {\\bf \\it binary star formation history} of the Universe. \\\\  X-ray studies will continue to be useful for probing the star formation history of the universe by avoiding problems of obscuration. Star formation may therefore be measured in more detail by deep surveys with future x-ray missions. ",
        "introduction": "There have been many recent studies of star formation in galaxies and of the star formation history of the universe derived from data in the radio, IR, and optical \\citep{lilly1996, madau1998, rowanrobinson1997, haarsma2000, Cole01,Baldry02,teplitz2003}. In the range from the present epoch to redshifts of order unity, recent critical compilations and discussions of \\citet{sullivan2001}, \\citet{hopkins2003}, \\citet{sullivan2004} and \\citet{hogg2004} show that the results from the multi-waveband data have a dispersion of 1-2 orders of magnitude in the comoving cosmic star formation density . As noted by these authors, there are important physical corrections that need to be made to go from the observations in a particular band to the cosmic star formation rate which include physical understanding of the dust extinction, the Initial Mass function (IMF), and stellar population models. Reasonable interpretations of the current observations in this redshift range have been presented that argue, on the one hand, that there is a steep decline in the star formation rate to the present epoch \\cite{hogg2004} or, on the other, that the cosmic star formation density has a shallow decline in the same redshift range \\cite{wilson2002}. Therefore it is important to utilize all wavebands to study this phenomenon from different aspects and with different selection effects. The X-ray band has now just opened up to detailed studies of star formation in galaxies at cosmological distances.  Hitherto, even deep X-ray surveys studied only the cosmological populations of evolving active galaxies and quasars. X-ray studies of individual nearby galaxies were performed and the underlying hot gas and stellar x-ray source components analyzed \\citep{fabbiano1989}. However, the deep 1-2 Megasecond surveys in the Chandra Deep Field South and North, respectively, now show a major cosmological population (in the range from present day to redshifts of order unity) of X-ray emitting normal star forming galaxies at faint flux levels. Stacking analysis allows one to extend the x-ray properties studied to even larger redshift ranges \\citep{hornschemeier2002}.  Some notable results on star-forming galaxies have already been derived in deep X-ray survey work.  In the $Chandra$ Deep Field-North (hereafter CDF-N) it has been found that the faint 15$\\mu$m population, composed primarily of luminous infrared starburst galaxies (e.g. \\citet{chary2001}) are associated with X-ray-detected emission-line galaxies (Alexander et al. 2002).  {\\it Chandra} and {\\it XMM-Newton} stacking analyses of relatively quiescent (non-starburst) galaxies have constrained the evolution of X-ray emission with respect to optical emission to at most a factor of 2--3 (e.g., Hornschemeier et al. 2002; Georgakakis et al. 2003).  Studies of the quiescent population of galaxies has also provided some initial constraints on the evolution of star-formation in the Universe ( SFR $\\sim (1+z)^k$ for $z <<1$ where k $<3$; Georgakakis et al. 2003).  We now discuss some relevant details of the X-ray observations. We derive X-ray Luminosity Functions and the corresponding cosmological star formation history using X-ray data from the 2~Ms and 1~Ms exposures of the CDF-N and {\\it Chandra} Deep Field South (hereafter CDF-S).  This analysis requires the extreme depth of the CDF surveys as non-active galaxies arise in appreciable numbers at extremely faint X-ray fluxes; the fraction of X-ray sources that are normal and star-forming galaxies rises sharply below 0.5--2~keV fluxes of $\\approx1\\times10^{-15}$~erg~cm$^{-2}$~s$^{-1}$ \\citep[e.g., Figure~6 of][; see also the fluctuation analysis results of Miyaji \\& Griffiths 2002]{hornschemeier2003}.  The luminosity functions for galaxies in the radio, optical and infrared have been studied extensively (c.f. \\citet{tresse2002} for a recent discussion). Prior to this work, the X-ray luminosity function for normal galaxies was estimated from the optical galaxy luminosity function by \\citet{georgantopoulos1999}.  Schmidt, Boller \\& Voges (1996) reported on a galaxy XLF derived in the local ($v < 500$~km~s$^{-1}$) Universe as well, however their sample does include some fairly active galaxies (e.g., Cen A) and does not appear to be an entirely clean normal/star-forming galaxy luminosity function.  Our work seeks to construct a relatively clean {\\it normal star-forming } galaxy luminosity function, and possesses the great advantage of doing this at cosmologically interesting distances (our two luminosity functions have median redshifts $z\\approx0.3$ and $z\\approx0.8$).  We may thus study the {\\it evolution} of X-ray emission from galaxies with respect to cosmologically varying quantities such as the global star-formation density of the Universe.  Our galaxy XLFs have their own set of selection effects that we discuss later but they can add to the overall picture of the physics of galaxy evolution. For example, it is obvious that corrections for obscuration will not be as important here as in the analysis of H$\\alpha$ Luminosity functions discussed in detail below.  We expect correlations between the various classic cosmological indicators of star formation (H$\\alpha$, IR, radio) and our X-ray studies. In general, the fluxes we measure in the respective wavebands associated with the different methods are all ultimately connected with the evolution and death and transfiguration of massive stars. For example, the radiative luminosity of OB stars, the mechanical energy input from supernovae and the accretion power of high mass X-ray binaries (HMXBs) are all ultimately derived from massive stars. It is therefore not surprising that tight correlations have been found empirically between the X-ray flux, the radio flux and the IR continuum.  The most recent careful analysis in the local Universe is from \\citet{ranalli2003} which is discussed in detail below. Interestingly, the correlations appear to hold for at least the radio and X-ray bands at high redshift \\cite{bauer2002,grimm2003}. We use these latest empirically derived correlations betwen radio, IR, and X-ray star formation indicators as the central relations that allow us to infer star formation rates from the X-ray data on galaxies in the Chandra Deep Fields North and South. In addition, it is useful to also consider theoretical studies on the cosmic X-ray evolution of galaxies which have been presented by \\citet{cavaliere2000}, \\citet{ptak2001} and \\citet{ghosh2001}.  Observations of detailed X-ray stellar populations in individual nearby galaxies have been carried out by \\citet{zezas2002},\\citet{kilgard2002} and \\citet{colbert2003}.  Using archival $Chandra$ data on 32 galaxies in the nearby Universe, Colbert et al. (2003) have established that the X-ray emission from accreting binaries in galaxies is correlated with the current level of star formation in their host galaxies.  This relation appears to be both macroscopic in that the entire X-ray point source luminosity of galaxies scales with star-formation rate and microscopic in that the slope of the X-ray binary luminosity function {\\it within} galaxies correlates with star formation.  In our studies, much care has gone into accurately selecting the population of normal star forming galaxies. Contamination by AGNs is a potentially serious problem. We have a multi-parameter probability space for the selection of the normal star forming galaxies and we have therefore used Bayesian methods to classify the galaxies and the two types of AGN. The most useful discriminators are the X-ray luminosity and X-ray hardness ratio as well as the optical to X-ray ratios. AGN contamination is most serious at the bright end of the luminosity function. Our analysis is a step towards a fully modern statistical Monte Carlo Markov-Chain study but we need to have a better idea of the overall structure of the probability space before we undertake this (c.f. \\citet{hobson2002} and \\citet{hobson2003}). Our analysis assumes, as a first step, that the priors are Gaussian and we show that this is a reasonable and useful assumption.  These studies of the cosmological evolution of normal star forming galaxies in the x-ray bands are the first of many detailed studies that can be undertaken with future missions such as XEUS. Here we have made a start on the basic luminosity functions and cosmic star formation histories.  The organization of this Paper is as follows: in section 2 we describe the data acquisition, the selection of the galaxy sample and the processing of the data. Section 3 presents the derivation of the galaxy XLF. In Section 4, we carefully compare the galaxy XLF derived here with the IR LF using the recent empirically derived analysis from \\citet{ranalli2003} (hereafter RCS). A similar procedure is undertaken for comparison with the H$\\alpha$LF.  We also give the derived cosmic star formation rates from the XLF.  In section 5 we discuss the implications of our results and the potential consequences for future missions. In section 6 we give our four principal conclusions. Throughout this paper, if not explicitly stated otherwise, we use $H_0$ = 70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m = 0.3$ and $\\Omega_{\\Lambda}=0.7$. ",
        "conclusions": "There are five main conclusions of this work.  1. We have established the X-ray luminosity functions of galaxies in the {\\it Chandra} deep Fields North and South. These XLFs have a lognormal distribution. The XLFs are completely consistent with both the infra-red luminosity functions (IRLFs) and the GHz radio luminosity functions(RLFs), which both have lognormal distributions.  2. The evolution of the XLFs is consistent with a pure luminosity evolution(PLE) of $\\sim (1+z)^{2.7}$.The more robust integral of the XLFs can be transformed to an x-ray derived cosmic star formation history. This cosmic star formation rate (SFR) is consistent with the (mild) evolution in SFR $\\propto (1 + z)^{2.7}$. In general, the overall SFR estimates in the range of redshift from unity to the present epoch show a wide scatter although there is a general trend of mild evolution consistent with the above. To reduce the scatter and better understand the evolution of the SFR in this redshift range one needs in the future:(1) more sensitive wide area studies such as the GALEX mission (2) extensive multi-band studies; the XLFs and X-ray derived SFRs now bring in a new waveband here and (3) deeper physical understanding of the star formation indicators.   3. The H$\\alpha$LFs have the form of the Schechter luminosity function which fits the galaxy luminosity functions in J and K. The H $\\alpha$ distribution is quite different from the the lognormal distributions for the XLF, IRLF and RLF discussed above. A detailed physical understanding is not yet available.   4. The X-ray derived star formation history gives an interesting constraint on the IMF. The X-ray emission in the Chandra band is most likely due to binary stars. The overall consistency and correlations between single star effects and binary star effects in different wavebands indicate that the bivariate IMF($M_1, M_2$) must be universally constant, at least at the high mass end, since the  X-ray observations may be measuring directly the {\\bf \\it binary star formation history} of the Universe."
    },
    "0402/astro-ph0402376_arXiv.txt": {
        "abstract": "We have re-examined the relation between the mass of the central black holes in nearby galaxies, $M_{bh}$, and the stellar mass of the surrounding spheroid or bulge, $M_{bulge}$. For a total of 30 galaxies bulge masses were derived through Jeans equation modeling or adopted from dynamical models in the literature. In stellar mass-to-light ratios the spheroids and bulges span a range of a factor of eight. The bulge masses were related to well-determined black hole masses taken from the literature. With these improved values for $M_{bh}$, compared to \\cite{mag98}, and our redetermination of $M_{bulge}$, we find the $M_{bh}-M_{bulge}$ relation becomes very tight. We find $M_{bh} \\sim M_{bulge}^{1.12\\pm0.06}$ with an observed scatter of $\\lesssim$ 0.30 dex, a fraction of which can be attributed to measurement errors. The scatter in this relation is therefore comparable to the scatter in the relations of $M_{bh}$ with $\\sigma$ and the stellar concentration. These results confirm and refine the work of \\cite{mar03}. For $M_{bulge}\\sim 5\\times 10^{10} M_{\\odot}$ the median black hole mass is $0.14\\% \\pm 0.04 \\%$ of the bulge mass. ",
        "introduction": "The good correlations between the mass of the central black hole and the physical properties of the surrounding stellar bulge provides evidence that black holes play a key role in the evolution of galaxies. So far, the tightest relation is that between the black hole mass $M_{bh}$ and the stellar velocity dispersion $\\sigma$ of the bulge stars \\citep{fer00, geb00}. Apart from that, other properties correlate with the mass of the black hole at the center of galaxies. \\cite{gra02} showed that $M_{bh}$ correlates tightly with the concentration of the host bulge as quantified by the Sersic index $n$. \\cite{mag98} explored the relation between $M_{bh}$ and bulge mass, $M_{bulge}$, finding $M_{bh}\\sim 0.005 M_{bulge}$, but with very large scatter. It is timely to reconsider this relation, since black hole mass measurements, now mostly based on HST data, have become much more reliable. In fact it is now clear that the black hole masses modeled by \\cite{mag98} were overestimated by as much as a factor of ten, since the black hole's sphere of gravitational influence was not well resolved in their data. This necessarily implies that the black hole mass fraction is lower than originally estimated \\citep{mer01}. \\\\ Recently, \\cite{mar03} showed that in the near infrared the correlation of the bulge luminosity and the black hole mass becomes much tighter than in the optical. They also demonstrated a tight relation between $R_e \\sigma^2_e$ and the black hole masses, where $R_e \\sigma^2_e$ represents a simple virial bulge mass estimate.\\\\ In this Letter we explore further the connection between the mass of the central black hole and the dynamical mass of the galaxy's bulge or spheroid in more detail, by combining direct $M_{bh}$ estimates, deemed reliable from other work, with $M_{bulge}$ determinations based on Jeans equation modeling, as opposed to virial estimates. ",
        "conclusions": "In Figure \\ref{fig2}, we plot $M_{bh}$ against the dynamical bulge mass $M_{bulge}$ for the 30 galaxies in the sample. In contrast to the relation presented by \\cite{mag98}, there is a tight relation without strong outliers.  A bisector linear regression fit \\citep{akr96} to the data leads to the relation $$\\log(M_{bh}/M_{\\odot}) = (8.20\\pm0.10)\\,+\\, (1.12\\pm0.06)\\,\\log(M_{bulge}/10^{11}M_{\\odot}),$$\\\\ where we included constant fractional errors of 0.18 dex on the bulge mass, the published uncertainties for the black hole masses (see Table \\ref{tab1}), and an intrinsic scatter of 0.3 dex. For this fit we have excluded the one possible outlier, NGC 4342; its inclusion leaves the slope of the relation unchanged. We calculated the mean values and their 1$\\sigma$ uncertainties using the bootstrap method \\citep{efr93}. Bootstrapping is preferable, since we do not have rigorous error bars for $M_{bulge}$, or for all $M_{bh}$ estimates. Whether we adopt intrinsic scatter or $\\delta$M$_{bulge}$ of 0.2 dex or 0.3 dex for the fit changes the slope by $\\lesssim$1\\% and the intercept by less than 0.3\\%. A least squares fit using the FITEXY routine \\citep{press92} changes the slope by $\\sim$15\\%, yielding 1.32$\\pm0.17$. All values agree within the stated uncertainties. \\\\ The upper limit on the intrinsic dispersion in the $M_{bh}-M_{bulge}$ relation, namely the observed dispersion assuming no measurement errors, is $\\sim 0.30$ dex. A significant portion of this scatter can plausibly be attributed to the observational errors in black hole masses. The implied median black hole mass fraction at bulge masses of $\\sim 5\\times 10^{10}M_{\\odot}$ is $M_{bh}$/M$_{bulge}=(1.4\\pm0.4)\\cdot10^{-3}$. This fraction is in agreement with the estimates from \\cite{mer01} and \\cite{mar03}. \\\\ At face value, the slope of the $M_{bh}-M_{bulge}$ relation exceeds unity with 1.5$\\sigma$ significance, both for the Akritas \\& Bershady estimator and the FITEXY routine, as also used by \\cite{mar03}. However, the data are still in agreement with a $M_{bh}-M_{bulge}$ proportionality at the $< 2 \\sigma$ level and we do not want to emphasize the non-linearity.  The mass-to-light ratios $\\Upsilon$ we found through the dynamical models are spread over a wide range from 0.15 to 8.0 $M_{\\odot}/L_{\\odot}$. Excluding the smallest value, which comes from NGC1068 (a galaxy with starburst activity), we still find a range for $\\Upsilon$ of a factor of eight. \\\\ Revisiting the Magorrian relation with more reliable black hole masses leads to a relation with a strongly reduced scatter. Our relation shows at most one outlier (NGC4342, among 30 objects) and we did not apply any selection criteria apart from the reliability of the black hole masses. \\\\  Our results confirm and expand the findings of \\cite{mar03}, who relate black hole masses to infrared luminosities and also to virial bulge mass estimates ($M\\sim \\sigma^2 r_e$). Their $M_{bh}-M_{bulge}$ relation for all galaxies is statistically in agreement with the relation we find. They find a slightly higher observed scatter, potentially because their virial estimate is less precise than the Jeans equation estimates, used here. \\\\ Through determining and compiling M$_{bulge}$ measurements for the objects with robust M$_{bh}$ estimates, we could demonstrate that the scatter in the black hole mass to bulge mass relation is nearly as small as that in the $M_{bh}-\\sigma$ \\citep{geb00,fer00} ($\\sim$0.3 dex) and $M_{bh}$-concentration \\citep{gra02} ($\\sim$0.31 dex) relation. Therefore, the relation between the black hole mass and the velocity dispersion is not unique and it seems as if the large scatter in the original Magorrian relation is due to erroneous estimation of the black hole masses. \\\\ Still, in the local universe $M-\\sigma$ is of invaluable practical use, since velocity dispersions are easy to measure. However, towards higher redshift ($z\\gtrsim 2$) the relation between black hole mass and stellar bulge mass gains importance. It is then exceedingly difficult to measure the velocity dispersion, but the bulge mass can be estimated via the measured luminosity and an upper limit of the stellar mass-to-light ratio, derived from the maximal age of the stellar population at that redshift."
    },
    "0402/astro-ph0402006_arXiv.txt": {
        "abstract": " ",
        "introduction": "The understanding of the generation and maintenance of a large scale magnetic field in astrophysical objects is a problem of exceptional importance and difficulty. It is widely accepted that the magnetic field is generated by the turbulent flow of the electrically-conducting fluid. Inhomogeneous velocity fluctuations stretch magnetic lines and amplify the magnetic field. These small scale fluctuations of turbulent flow are primarily responsible for the generation of magnetic fields. The problem is that it is difficult to resolve them. Thus their influence on the resolved large-scale magnetic field has to be modelled. A traditional closure scheme is based on the $\\alpha -$effect according to which small-scale fluctuations can be described by an average term involving the curl of the mean magnetic field, $\\mathbf{B,}$ written as $\\nabla \\times (\\alpha \\mathbf{% B)}.$ The asymptotic analysis of the induction equation exploiting the assumption of two separated scales for turbulent flow leads to the effective macroscopic equation for the large scale magnetic field $\\mathbf{B}(t,% \\mathbf{x})$ [1-5] \\begin{equation} \\frac{\\partial \\mathbf{B}}{\\partial t}=\\nabla \\times (\\alpha \\mathbf{B}% )-\\nabla \\times (\\beta \\nabla \\times \\mathbf{B)}+\\nabla \\times (\\mathbf{u}% \\times \\mathbf{B}),  \\label{dynamo} \\end{equation} where $\\ \\mathbf{u}\\ $ is the mean velocity field of the turbulent flow, $\\ \\alpha \\ $ is the coefficient of the $\\alpha $-effect and $\\beta $ is the turbulent magnetic diffusivity. Traditionally the phenomenon of generation of magnetic fields was analyzed by considering perturbations of the trivial state $\\mathbf{B}=0$ and looking for exponential solutions to the deterministic PDE (\\ref{dynamo}) with appropriate boundary conditions. While this standard stability analysis successfully predicts the dynamo action for the supercritical case, there are situations for which this eigenvalue analysis fails to predict the subcritical onset of instability [6-10]. It was pointed out in \\cite{FI1} that the closure scheme involving only deterministic $\\alpha \\beta -$parameterization is not completely satisfactory since unresolved fluctuations may produce random terms on the right hand side of the dynamo equation (\\ref{dynamo}). Moreover, it follows from astronomical observations that large scale magnetic fields exhibit a rich random variability both in space and time that cannot be described by the deterministic equation (\\ref {dynamo}).  The importance of noise effects in the dynamo problem has been recognized previously and several attempts have been made to account for the effects of spatial and temporal fluctuations in small scale magnetic and velocity fields on the generation process. Kraichnan considered fluctuations in the $% \\alpha $-parameter and found a negative contribution \\ to turbulent diffusivity from helicity fluctuations \\cite{Kr}. Hoyng with colleagues in [12-14] studied in detail the effect of random fluctuations in the $\\alpha $% -parameter by considering the system of stochastic linear equations for eigenmodes corresponding to the dynamo equations. They found the excitation of those modes such that their magnetic energy is proportional to $\\gamma ^{-1},$ where $\\gamma $ is the damping rate. Stochastic dynamics of magnetic field generation have been also analyzed by Farrell and Ioannou in \\cite{FI1}% , where they examine a mechanism by which small-scale fluctuations excite the large scale magnetic field. They modelled these fluctuations by an additive noise term in the mean field equation and identified the crucial role of non-normality on the dynamo process. Numerical simulations of the magnetoconvection equations were performed in \\cite{Proctor} with analysis of effects of noise and non-normal transient growth. Inhomogeneous turbulent helicity fluctuations were considered in \\cite{Sil}. One should mention the dynamo model that exhibits aperiodic switching between regular behavior and chaotic oscillation \\cite{PST2}. Stochastic dynamo theory, using the term of an incoherent dynamo, was proposed by Vishniac and Brandenburg in \\cite{VB}. They showed how random fluctuations in the helicity can generate a large-scale magnetic field for the $\\alpha \\Omega -$% dynamo. We note that this model is closely related to the present work. However, they did not consider the transient amplification due to the non-normality of the dynamo equation. It turns out that non-normal dynamical systems exhibit \\ an extraordinary sensitivity to random perturbations which leads to great amplification of the second moments of the stochastic dynamical systems (see \\cite{Farrel1,Grossmann,Fedotov}). Our recent work has demonstrated the possibility for stochastic magnetic energy amplification in the subcritical situation where the dynamo number is less than critical \\cite {F1}. These observations motivate further studies of noise and nonlinearity effects which we consider below in the context of a simplified stochastic no-% $z$ model (see \\cite{Moss}). An issue we address in this paper is how the random fluctuations may be appropriately incorporated into the classical $% \\alpha \\Omega -$dynamo model. Our analysis especially focuses upon the multiplicative noise due to random fluctuations in the $\\alpha $-parameter and the non-normality of the dynamo equation operator. Recall that an operator is said to be non-normal if it does not commute with its adjoint in the corresponding scalar product. The determination of the effect of the noise, along with non-normality, on the amplification of magnetic energy during the early kinematic stage is the primary contribution of this paper.  In section II, we discuss the deterministic $\\alpha \\Omega -$dynamo model,\\ its equilibrium points and transient growth effects. In section III, we consider a linear stochastic model for the subcritical case. We derive equations for the second moments and find their stationary values. We demonstrate the important differences between the non-normal system and the normal one under the influence of additive noise. We explore the influence of multiplicative noise due to fluctuations in the $\\alpha $-parameter on the amplification of magnetic energy in the subcritical case during the kinematic stage. We derive the criteria under which the second moments grow exponentially with time (kinematic regime). Finally, in section IV, we perform numerical experiments showing that there exists a series of noise-induced phase transitions in the $\\alpha \\Omega -$dynamo model. ",
        "conclusions": "We have studied the stochastic amplification of large scale magnetic fields in a differentially rotating system in a subcritical regime and discussed the possible implications of this for the magnetic field in galaxies. The main purpose was to address the stochastic generation that cannot be explained by traditional linear eigenvalue analysis. We have chosen the simplified stochastic $\\alpha \\Omega -$dynamo model for galaxies in a thin-disk approximation and thereby concentrated on the influence of\\ additive and multiplicative noises along with non-normality on the amplification of the magnetic field in the subcritical case (when the dynamo number is less than critical). In the linear case, we have derived the equations for second moments describing the magnetic energy and demonstrated the important differences between the non-normal system and the normal one under the influence of additive noise. For the multiplicative noise, the criteria for the stochastic instability during the early kinematic stage was established in terms of the critical value of noise intensity due to $\\alpha $-fluctuations. In the non-linear case, we have performed numerical simulations of non-linear stochastic differential equations for the $\\alpha \\Omega -$dynamo and found a series of noise induced phase transitions: qualitative changes in the behavior of the trajectories due to the increase in the noise intensity parameter.  It should be noted that the equations (\\ref{basic1}),(\\ref{basic}) as applied to galaxies are a theoretical simplification, but the do provide a useful framework for understanding the effect of random fluctuations. Our finding for the stochastic parametric instability can be straightforwardly applied to partial differential equations (\\ref{main}) in which case one has to consider the coupled system for eigenmodes \\cite{Hoyng}. In this respect the two equations (\\ref{basic1}),(\\ref{basic}) can be regarded as a dynamical system for first eigenmodes (order parameters), where the influence of other degrees of freedom is parameterized by additive and multiplicative noises. We believe that our theory can also be extended and applied (after some modifications) to the solar dynamo in which a toroidal field is generated by the action of a shear flow (typical non-normal effect).  \\textbf{Acknowledgment}. In this project we benefited from the financial support of this work by the Royal Society-Russia-UK Joint Project Grant. SF, who was supported by Center of Turbulence Research, Stanford, is grateful to Parviz Moin and Heinz Pitsch for a hospitality and fruitful discussions.  \\bigskip"
    },
    "0402/gr-qc0402088_arXiv.txt": {
        "abstract": "In a previous work \\cite{mbeleka}, we have modelled the rotation curves (RC) of spiral galaxies by including in the equation of motion dynamical terms from an external real self-interacting scalar field, $\\phi$, minimally coupled to gravity and which respects the equivalence principle in the absence of electromagnetic fields.  This model appears to have three free parameters :  the turnover radius, $r_{0}$, the maximum rotational velocity, $v_{max} = v(r_{0})$, plus a strictly positive integer, $n$.  Here, the coupling of the $\\phi$-field to other kinds of matter is emphasized at the expense of its self-interaction.  This reformulation presents the very advantageous possibility that the same potential may be used now for all galaxies.  New correlations are established. ",
        "introduction": "The equation of motion of a neutral test body reads in the presence of the $\\phi$-field \\begin{equation} \\label{eq motion} \\frac{du^{\\mu}}{ds} \\,+ \\,{\\Gamma}^{\\mu}_{\\alpha\\beta} \\,u^{\\alpha} \\,u^{\\beta} = \\frac{d\\phi}{ds} \\,u^{\\mu} \\,- \\,{\\partial}^{\\mu} \\phi, \\end{equation} where the geodesic equation is recovered in the case of a non variable $\\phi$-field. In the weak fields and low velocity limit, assuming spherical symmetry and circular orbits, the rotational velocity, $v$, at radius $r$ is then given by \\begin{equation} \\label{angular momentum vs phi} \\ln{(rv)} = \\ln{J} + \\phi, \\end{equation} where $J$ is a constant which would represent the angular momentum per unit mass if the $\\phi$-field were not present within the galaxy\\footnote{Here, we just deal with the tangential equation since we have assumed circular orbits.  The radial equation would yield the radial velocity but, it is found negligible with respect to the tangential velocity, out of the bulge.}.  Assuming that the excitation of the $\\phi$-field is very small compared to its vacuum expectation value, ${\\phi}_{vev}$, the equation of the $\\phi$-field reads in the first approximation \\begin{equation} \\label{eq phi} {\\partial}_{\\mu}\\,{\\partial}^{\\mu} \\phi = - \\,g \\,\\chi \\,{\\phi}_{vev} \\,T, \\end{equation} where $T$ denotes the trace of the energy-momentum tensor of the source of the $\\phi$-field, $\\chi = 8\\pi G/c^{4}$ is the Einstein gravitational constant and $g$ a universal dimensionless coupling constant.  Hence, one gets for a static spherical matter distribution \\begin{equation} \\label{eq phi static spherical} \\frac{d^{2}\\phi}{dr^{2}} \\,\\,+ \\,\\,\\frac{2}{r} \\,\\frac{d\\phi}{dr} = g \\,\\frac{8\\pi G}{c^{2}} \\,{\\phi}_{vev} \\,\\rho, \\end{equation} where $\\rho = {\\rho}_{bulge} \\,+ {\\rho}_{disk} \\,+ \\,{\\rho}_{halo}$ denotes the mass density of the matter fields other than the $\\phi$-field itself.  Assuming in addition a sufficiently thin disk (stellar plus gaseous disks) and a (quasi-isothermal) spherical dark halo with mass density such that \\begin{equation} \\label{halo density} {\\rho}_{halo} \\propto 1/r^{2 \\,+ \\,1/n}, \\end{equation} the static spherical solution of $\\phi$ is found proportional to $r^{-1/n}$ (up to the vacuum expectation value) within the galaxy out of the bulge.  Hence, the rotational velocity reads \\begin{equation} \\label{velocity vs radius} v = v_{max} \\,G_{n}(r/r_{0}), \\end{equation} where the functions $G_{n}$ are defined as follows \\begin{equation} \\label{universal RC} G_{n}(x) = \\frac{1}{x} \\,\\exp{[ \\,n(1 - x^{-1/n}) \\,]}.  \\end{equation} For our purpose, $n \\geq 2$.  Indeed, this is one of the two necessary conditions for the $\\phi$-field mimics a dark matter mass profile ($\\propto r^{1-1/n}$). ",
        "conclusions": "It is possible that long range scalar fields external to gravity play a significant role not only at the cosmological level but also at the scale of galaxies or even the solar system \\cite{mbelekb}."
    },
    "0402/astro-ph0402230_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:intro}  Since its introduction roughly 20 years ago \\cite{Berger84}, adaptive mesh refinement (AMR) has emerged as an important class of numerical techniques for improving the accuracy and dynamic range of grid-based calculations for fluid dynamics problems. Such problems, especially compressible flow, develop steep gradients (shock waves and contact discontinuities) which, in the absence of mesh refinement, become sources of error for the global solution (e.g., \\cite{Woodward84}). Through appropriate local mesh refinement, AMR can be thought of as a numerical technique for optimizing the quality of a numerical solution for a given computational cost (e.g., \\cite{Berger89}).  The last ten years have seen the application of AMR methods to problems in space physics \\cite{deZeeuw98,Groth99,Groth00}, astrophysics \\cite{Falle93, Falle95, Klein94, Klein98, Klein03, Truelove97, Truelove98, Khokhlov98, Kritsuk01, Falle02, Calder02}, and cosmology \\cite{Bryan97, Bryan00, Kravtsov97, Klypin99, ABN00, ABN02, Loken02, Tassis03, Nagai03}. Here, a variety of physical processes may operate singly or together in astrophysical fluids to expand the range of important length- and time-scales. Such processes include gravity and gravitational instability, reaction kinetics, magnetic reconnection, radiation transfer, ionization fronts, etc. AMR has also been applied to the solution of gravitational N-body problems \\cite{Couchman91, Kravtsov97}, and to hybrid particle-fluid simulations in cosmology \\cite{Bryan97, Bryan99}. In such applications, the physics is intrinsically multi-scale, and AMR can be thought of as a numerical technique for extending the dynamic range of resolved physics, regardless of the computational cost. It is these sorts of applications which are reviewed here, and where AMR can make a scientific and not just an economic impact.  AMR is also having a positive impact on the methodology of computational physics itself. I am referring to the validation of computational codes through resolution studies. With AMR, it is now practical on today's supercomputers to perform resolution studies over a sufficient range of scales to obtain convergent results on properties of interest which may be compared with laboratory experiments (e.g., \\cite{Calder02}).  Table 1 shows the impressive diversity of topics AMR has been applied to. Particularly interesting is the range of physical processes AMR has been adapted to model. This is listed in the second column of Table 1. As of today, AMR has been successfully applied to ideal gas dynamics (Newtonian and special relativistic), reactive gas dynamics, MHD (ideal and resistive), self-gravitating gas dynamics and MHD, N-body dynamics, and hybrid fluid/N-body systems. As we heard at this conference, methods are under development for radiative transfer (Howell), radiation hydrodynamics (Weaver), and solid mechanics (Falle). It is clear AMR is a method of wide applicability, and one that is growing in its impact. This growth is fueled in part by the public availability of AMR codes and frameworks (see Section 3).  \\begin{table} \\centering \\begin{tabular}{|l|l|l|} \\hline \\bf{Topic}   & \\bf{Physics} & {\\bf Select References} \\\\ \\hline \\hline Code validation & HD, reactive HD & \\cite{Calder02} \\\\ \\hline Solar and space physics & MHD & \\cite{Groth99} \\\\ \\hline Supernovae and nucleosynthesis & reactive HD & \\cite{Kifonidis00, Timmes00, Gamezo03} \\\\ \\hline Interstellar medium & HD, MHD & \\cite{Klein94, Kritsuk01, Cid96, Balsara01} \\\\ \\hline Star formation & grav HD, grav RHD & \\cite{Truelove98, Klein03} \\\\ \\hline Astrophysical jets & HD, rel HD & \\cite{Falle95, Hughes97, KF97} \\\\ \\hline N-body dynamics & particles, grav & \\cite{Kravtsov97, Klypin99} \\\\ \\hline Hydrodynamic cosmology & hybrid & \\cite{Bryan97, Nagai03} \\\\ \\hline \\hline  \\end{tabular} \\newline \\caption{Classes of AMR applications.} \\label{applications} \\end{table} \\medskip  In this paper I survey the use and impact of adaptive mesh refinement simulations in numerical astrophysics and cosmology. Two basic techniques are in use to extend the dynamic range of Eulerian grid simulations in multi-dimensions: cell refinement (CR), and patch refinement (PR), otherwise known as block-structured adaptive mesh refinement (SAMR). Details of these two approaches are given elsewhere in this volume. In this review, no attempt is made to assess the relative merits of these two approaches. Rather, the discussion focuses on how AMR is being used and how AMR is making a scientific impact in a diverse set of fields from space physics to the cosmology of the early universe. At the end, I provide a partial list of software resources for those interested in learning more about AMR simulations. [{\\bf NB.} The following is meant to be representative rather than complete. I apologize to authors in advance if their work is not mentioned.] ",
        "conclusions": ""
    },
    "0402/astro-ph0402224_arXiv.txt": {
        "abstract": "Clusters of galaxies can be seen as giant astrophysical laboratories enclosing matter in a large enough volume, so that the matter composition can be taken as representing the composition of our Universe. X-ray observations allow a very precise investigation of the physical properties of the intracluster plasma allowing us to probe probe the cluster structure, determine its total mass, and measure the baryon fraction in clusters and in the Universe as a whole. We can determine the abundance of heavy elements from O to Ni which originate from supernova explosions and draw from this important conclusions on the history of star formation in the cluster galaxy population. From the entropy structure of the intracluster medium we obtain constraints on the energy release during early star bursts. With the observational capabilities of the X-ray observatories XMM-Newton and Chandra this field of research is rapidly evolving. In particular, first detailed observations of the intracluster medium of the Virgo cluster around M87 have provided new insights. The present contribution gives an account of the current implications of the intracluster medium observations, but more importantly illustrates the prospects of this research for the coming years. ",
        "introduction": "Clusters of galaxies, the largest well defined objects, are the largest well characterized astrophysical laboratories at our disposal, except for the Universe as a whole. The clusters masses range from less than $10^{14}$ to several $10^{15}$ M$_{\\odot}$. This mass range is extended down to about several $10^{12}$ M$_{\\odot}$ by galaxy groups which can be considered as small scale versions of clusters. Hundreds to Thousands of cluster galaxies are making up merely a few percent of the total cluster mass while more mass is in the gaseous intracluster medium (ICM). The dominant fraction is Dark Matter. Being formed through gravitational collapse where gravity acts on all forms of matter equally, galaxy clusters contain to a good approximation a composition of matter well representative of the Universe as a whole (e.g. White et al. 1993).  In this contribution I will focus on the information we can gain on galaxy evolution from the study of the chemical composition and the thermodynamic structure of the ICM. One of the astronomically most interesting epochs in the history of our Universe was that of the most intense star and galaxy formation at redshift between 2 and 4 (e.g. Madau et al. 1998). The traces of what happened then can be found in the ICM today: i) we observe that the ICM is enriched by metals which must have been synthezised by supernovae in the cluster galaxies mostly in these early star formation epochs; ii) we further observe in the entropy structure of groups and clusters of galaxies the effect of an early heating of the intergalactic medium by energy release in the star formation epochs.  The ICM is a hot plasma with temperatures of several ten Million degrees which has its thermal radiation maximum in the soft X-ray regime. It is a fortunate coincidence that it is the same wavelength region in which we can use the current technology of imaging X-ray telescopes. Therefore we know most about the ICM through X-ray astronomy. We have furthermore two new satellite X-ray observatories, ESA's XMM-Newton mission and NASA's Chandra mission, with greatly improved spectral and imaging capabilities, which are providing a overwhelming new insight into the physics of the ICM. In this contribution I will use a varying scaling of physical parameters the Hubble constant (depending on the sources) quoted which will be labled by the scaling parameter $h$, e.g. $h_{70} = H_0 / 70$ km s$^{-1}$ Mpc$^{-1}$.   \\begin{figure}[ht] \\plotone{Boehringerfig1.ps} \\caption{Gas mass fraction in units of $h_{50}^{-1.5}$ for 106 of the brightest clusters of galaxies detected by ROSAT in X-rays (from Reiprich 2001). } \\end{figure} ",
        "conclusions": ""
    },
    "0402/astro-ph0402538_arXiv.txt": {
        "abstract": "We present a modified TREESPH code to model galaxies in 3d. The model includes a multi-phase description of the interstellar medium which combines two numerical techniques. A diffuse warm/hot gas phase is modelled by SPH while a sticky particle scheme is used to represent a cloudy medium. Interaction processes, such as star formation and feedback, cooling and mixing by condensation and evaporation, are taken into account. Here we apply our model to the evolution of a Milky Way type galaxy. After an initial stage, a quasi-equilibrium state is reached. It is characterised by a star formation rate of $\\sim$1\\,${\\rm M}_{\\odot}\\,{\\rm yr}^{-1}$. Condensation and evaporation rates are in balance at 0.1-1\\,${\\rm M}_{\\odot}\\,{\\rm yr}^{-1}$. ",
        "introduction": "In most 3d models of galaxies the interstellar medium (ISM) is described as a single phase: either as a diffuse gas e.g.\\ using \\revised{smoothed particle hydrodynamics (SPH)} (\\revised{Lucy, 1977; Gingold \\& Monaghan, 1977}) or as a clumpy phase e.g.\\ by sticky particles as in Theis \\& Hensler (1993). \\revised{Single-phase models have been successfully applied in the context of cosmological simulations (Steinmetz \\& M\\\"uller, 1994, 1995; Navarro, Frenk \\& White, 1997) as well as in simulations of isolated galaxies (Friedli \\& Benz, 1995; Raiteri, Villata \\& Navarro, 1996; Berczik, 1999). A problem in these simulation can be the so-called overcooling (White \\& Frenk, 1991; Cole et al., 1994) leading to an overproduction of dwarf galaxies.} \\revised{More recent models use a multi-phase ISM. This allows a more realistic description of star formation (SF) and feedback processes which could be a solution to this problem.}  \\revised{A multi-phase ISM can be} accomplished by modifying the SPH-algorithm (\\revised{Hultman \\& Pharasyn, 1999;} Ritchie \\& Thomas, 2001; Springel \\& Hernquist, 2003). We follow a different approach and combine both treatments of a one-phase model in a particle based code: the hot diffuse gas phase is described by a SPH formalism, whereas the cold molecular clouds are represented by sticky particles. In this model, the two gas phases are dynamical independent. The coupling between the gaseous phases is achieved by condensation/evaporation (C/E) and by a drag force due to ram pressure. Energy is dissipated by cloud-cloud-collisions or by radiative cooling. Furthermore, stars are formed in clouds and the stars return mass and energy to the gas by feedback processes (supernovae (SNe), planetary nebulae (PNe)). A similar concept is followed by Semelin \\& Combes (2002) and Berczik et al. (2003).  \\begin{figure*}[!t] \\begin{center} \\begin{tabular}{ccc} \\psfig{width=8cm,angle=0,file=harfstf1a.eps} && \\psfig{width=8cm,angle=90,file=harfstf1b.ps} \\\\ \\end{tabular} \\caption{ The left diagram shows the SF scheme: ($t_0$) the cloud is inactive (no SF); ($t_1$) an embedded star cluster has been formed with a local SF efficiency; ($t_2$) the cloud is fragmented by SNe energy input. \\revised{In the right plot the local SF efficiency depending on the mass of the star forming cloud and the local pressure is shown. It is based on a model proposed by Elmegreen \\& Efremov (1997).}} \\label{harfst_fig1} \\end{center} \\end{figure*}  In SPH codes the star formation (SF) is usually based on the Schmidt law, i.e.\\ the SF rate depends on gas density to some power and a characteristic time scale. For sticky particles the SF can be coupled to cloud-cloud collisions, or a single cloud forms stars with a constant SF efficiency. Here, a SF scheme using a different approach is presented: the process of SF is treated individually for each cloud. The SF efficiency depends on local properties of the ISM and the star forming clouds, thereby enabling self-regulation of SF by feedback.  The code is described in more detail in Sec.~\\ref{harfst_sec_code}. In Sec.~\\ref{harfst_sec_res} a model of a Milky Way type galaxy is presented. A short summary and future prospects are given in Sec.~\\ref{harfst_sec_sum}. ",
        "conclusions": "\\label{harfst_sec_sum}  We presented a particle code to model the evolution of galaxies with a multi-phase description of the ISM including a new approach to SF. A model of a Milky Way type galaxy shows a reasonable behavior in terms of SF rate and stability of the disk.  In the near future more simulations are needed to determine how the results depend on initial conditions, \\revised{e.g. the initial set-up for the hot gas or initial particle numbers} and how other parameters in the model, e.g. the energy feedback by SNe, affect the evolution of the model galaxy. This should result in a reference model, that could be used in simulations of interacting galaxies, e.g.\\ to shed more light on trigger mechanisms of star bursts."
    },
    "0402/astro-ph0402012_arXiv.txt": {
        "abstract": "{ Recently, twin-peak QPOs have been observed in a 3:2 ratio for three Galactic black-hole microquasars with frequencies that have been shown to scale as 1/$M$, as expected for general relativisitic motion near a black hole.  It may be possible to extend this result to distinguish between the following two disparate models that have been proposed for the puzzling ultraluminous X-ray sources (ULXs): (1) an intermediate-mass black hole ($ M \\sim 10^3 M_{\\odot}$) radiating very near the Eddington limit and (2) a conventional black hole ($ M \\sim 10 M_{\\odot}$) accreting at a highly super-Eddington rate with its emission beamed along the rotation axis.  We suggest that it may be possible to distinguish between these models by detecting the counterpart of a Galactic twin-peak QPO in a ULX: the expected frequency for the intermediate-mass black hole model is only about 1 Hz, whereas, for the conventional black hole model the expected frequency would be the $\\sim 100$ Hz value observed for the Galactic microquasars.  } ",
        "introduction": " ",
        "conclusions": "The detection of a twin-peak 3:2 QPO in a ULX would immediately determine its mass (see Fig.~\\ref{figure}) and thereby solve the puzzle: is the correct model for a ULX based on an intermediate-mass black hole ($ M \\sim 10^3 M_{\\odot}$) or a conventional black hole ($ M \\sim 10 M_{\\odot}$) embedded in a Polish doughnut?  Finally, we note that Mirabel's quasar-microquasar analogy suggests that one should also expect twin-peak QPOs with a 3:2 ratio in the microhertz range for quasars (Fig.~\\ref{figure}; $ M/M_{\\odot} \\sim 10^7 $ to $ 10^9$). In this connection, it is interesting to note that a 17 minute periodicity has recently been reported from the compact radio source Sgr A$^*$, ``a 3.6-million-solar-mass black hole'' at the Galactic Centre (Genzel et al. 2003). This periodicity seems to correspond exactly to the 1 mHz frequency expected on the scaling discussed here (Fig.~1). However, the result is inconclusive because only a single oscillation frequency was observed. If the 17 minute flare period does indeed correspond to the upper (or lower) of the twin-peak QPOs in microquasars, it would be interesting to see whether a 26 minute (or 12 minute) quasi-periodicity may also be present in the source.  After this paper was completed, Aschenbach et al. (astro-ph/0401589) reported X-ray QPOs from Sgr A*. The claimed periods include 1150 s (19 minutes) and 700s (12 minutes)."
    },
    "0402/astro-ph0402648_arXiv.txt": {
        "abstract": "The Sloan Digital Sky Survey has discovered a $z=2.4917$ radio-loud active galactic nucleus (AGN) with a luminous, variable, low-polarization UV continuum, \\hi\\ two-photon emission, and a moderately broad \\lya\\ line (FWHM$\\simeq$1430\\,\\kms) but without obvious metal-line emission. SDSS J113658.36+024220.1 does have associated metal-line absorption in three distinct, narrow systems spanning a velocity range of 2710 \\kms. Despite certain spectral similarities, \\sdssj\\ is not a Lyman-break galaxy. Instead, the \\lya\\ and two-photon emission can be attributed to an extended, low-metallicity narrow-line region. The unpolarized continuum argues that we see \\sdssj\\ very close to the axis of any ionization cone present.  We can conceive of two plausible explanations for why we see a strong UV continuum but no broad-line emission in this `face-on radio galaxy' model for \\sdssj: the continuum could be relativistically beamed synchrotron emission which swamps the broad-line emission; or, more likely, \\sdssj\\ could be similar to \\pg, a quasar in which for some unknown reason the high-ionization emission lines are very broad, very weak, and highly blueshifted. ",
        "introduction": "\\label{INTRO}  One of the goals of the Sloan Digital Sky Survey \\markcite{yor00}(SDSS; {York} {et~al.} 2000) is to obtain spectra for $\\sim$10$^5$ quasars, in addition to the $\\sim$10$^6$ galaxies which comprise the bulk of the spectroscopic targets \\markcite{str02,sdss89}({Strauss} {et~al.} 2002; {Blanton} {et~al.} 2003).  From astrometrically calibrated drift-scanned imaging data \\markcite{gun98,sdss153}({Gunn} {et~al.} 1998; {Pier} {et~al.} 2003) on the SDSS $ugriz$ AB asinh magnitude system \\markcite{fuk96,sdss26,sdss82,sdss85,sdss105}({Fukugita} {et~al.} 1996; {Lupton}, {Gunn}, \\& {Szalay} 1999; {Hogg} {et~al.} 2001; {Stoughton} {et~al.} 2002; {Smith} {et~al.} 2002), quasar candidates are selected primarily using color criteria designed to target objects whose broad-band colors differ from those of normal stars and galaxies \\markcite{sdssqtarget}({Richards} {et~al.} 2002a). Some of these quasars are bound to possess unusual properties.  Here we describe one such unusual quasar, SDSS J113658.36+024220.1 (hereafter SDSS J1136+0242), which exhibits strong \\lya\\ emission but no detected metal-line emission. We discuss its unusual properties (\\S \\ref{SPEC}), consider possible explanations for them (\\S \\ref{WASSUP}), and conclude with suggestions for further investigation (\\S \\ref{CON}). Where needed, we assume \\ho=70 \\kmsm, $\\Omega_M$=0.7 and $\\Omega_{\\Lambda}$=0.3. ",
        "conclusions": "\\label{CON}  The newly discovered, radio-loud AGN \\sdssj\\ is unusual in possessing a strong continuum without accompanying strong metal-line emission. Of $\\sim$2400 quasars in the SDSS Data Release 1 quasar catalog \\markcite{dr1q}({Schneider} {et~al.} 2003) with $z\\geq2.12$, sufficient to place \\lya\\ within the spectral coverage of the SDSS spectrographs, it is the only object with strong, relatively narrow \\lya\\ emission and no accompanying metal-line emission (excepting broad absorption line quasars where such emission is absorbed along with part or all of the continuum).  The spectrum of \\sdssj\\ can generally be understood as that of a very optically luminous radio galaxy, in terms of \\lya\\ emission flux, velocity width, associated narrow \\lya\\ and \\civ\\ absorption, \\civ/\\lya\\ ratio (from a low-metallicity NLR), and even two-photon continuum strength (\\S\\ref{2Q} and \\S\\ref{NLAGN}).  However, \\sdssj\\ has a very strong UV continuum not seen in radio galaxies (\\S\\ref{CONT}).  This continuum is unpolarized, suggesting a face-on orientation for \\sdssj. The continuum could be relativistically beamed synchrotron emission which swamps the intrinsic broad-line emission, but the SED of \\sdssj\\ is not a good match to known synchrotron-dominated objects. A more likely possibility is that \\sdssj\\ is a \\pg\\ analog: a low-polarization radio quasar which for some unknown reason has very weak, very broad, highly blueshifted high-ionization emission lines. A better optical spectrum is needed to determine if such features are present in \\sdssj. This scenario predicts polarization perpendicular to the axis of the radio jet, which is tentatively observed to be the case in both \\sdssj\\ and \\pg\\ but requires confirmation with better polarization data.  In the unlikely event the \\lya\\ is from a broad-line region of typical density, the dip in the continuum otherwise attributed to \\hi\\ two-photon emission is unexplained. Also, the BLR would have to consist predominantly of low-ionization or low-metallicity gas (Appendix \\ref{BLAGN}). If the former, it is unclear what about \\sdssj\\ causes its lack of the high-ionization gas which dominates the broad-line regions of most AGNs. If the latter, metallicities $\\lesssim$1\\% of solar are required.  Further investigation of \\sdssj\\ would benefit from better optical spectra (preferably spatially resolved to search for and measure any spatial extent of \\lya), better optical polarimetry (or at least another epoch of the same quality), deep, high-resolution radio imaging and polarimetry to determine the jet axis, near-IR spectroscopy to study the rest-frame optical spectrum, and multiwavelength photometry to determine if its spectral energy distribution can be reconciled with those of blazars."
    },
    "0402/astro-ph0402362_arXiv.txt": {
        "abstract": "We present new results of our wide-field redshift survey of galaxies in a 182 square degree region of the Shapley Supercluster (SSC) based on observations with the FLAIR-II spectrograph on the UK Schmidt Telescope (UKST). In this paper we present new measurements to give a total sample of redshifts for 710 bright ($R\\leq 16.6 $) galaxies, of which 464 are members of the SSC ($8000<v<18000\\kms $). Our data reveal that the main plane of the SSC ($v\\approx 14500\\kms $) extends further than previously realised, filling the whole extent of our survey region of 10 degrees by 20 degrees on the sky (35 Mpc by 70 Mpc, $H_0=75\\kmsMpc$). There is also a significant structure associated with the slightly nearer Abell 3571 cluster complex ($v\\approx12000\\kms$) with a caustic structure evident out to a radius of 6 Mpc.  These galaxies seem to link two previously identified sheets of galaxies and establish a connection with a third one at $\\overline V$= 15000\\kms near $R.A.= 13h$. They also tend to fill the gap of galaxies between the foreground Hydra-Centaurus region and the more distant SSC. We calculate galaxy overdensities of $5.0\\pm0.1$ over the 182 square degree region surveyed and $3.3\\pm0.1$ in a 159 square degree region excluding rich clusters. Over the large region of our survey the inter-cluster galaxies make up 46 per cent of all galaxies in the SSC region and may contribute a similar amount of mass to the cluster galaxies. ",
        "introduction": "In the past few decades, large galaxy redshift surveys have revealed structures on ever-increasing scales. The largest structures found are superclusters, collections of thousands of galaxies with sizes as large as 100 Mpc. The mere existence of these structures places important constraints on theories of the formation of galaxies and clusters. The Shapley supercluster, the subject of this paper, is one of the most massive concentrations of galaxies in the local universe (Scaramella et al.\\ 1989; Raychaudhury 1989), so it is also of particular interest to consider its effect on the dynamics of the Local Group.  The Shapley supercluster (SSC) is a remarkably rich concentration of galaxies centred around R.A.$=13^h25^m$ Dec. $= -30^{\\circ}$ which has been investigated by numerous authors since its discovery in 1930 (see Quintana et al. 1995). It consists of many clusters and groups of galaxies in the redshift range $0.04<z<0.055$. The SSC lies in the general direction of the dipole anisotropy of the Cosmic Microwave Background (CMB), (Smoot et al. 1992), and is located 130\\hMpc\\ beyond the Hydra-Centaurus supercluster (itself $\\simeq 50$\\hMpc\\ away from us).  Quintana et al.\\ (1995) estimated that for $\\Omega_{o}= 0.3$ and $H_{o}$= 75\\kmsMpc the gravitational pull of the supercluster may account for up to 25\\% of the peculiar velocity of the Local Group required to explain the CMB dipole anisotropy in which case the mass of the supercluster would be dominated by inter-cluster dark matter. A major study of the SSC was made by Bardelli et al.\\ (2000, 2001 and references therein) who also studied inter-cluster galaxies in the core region of the supercluster. They derived a total mean overdensity of the SSC of $N/\\overline N \\sim 11.3$ on a scale of $10.1h^{-1}$ Mpc and found that the central part of the supercluster contributes about 26\\kms\\ to the peculiar velocity of the Local Group.   The early studies of the Shapley supercluster were limited (primarily by observational constraints) to the various rich Abell galaxy clusters in the region, but this might give a very biased view of the overall supercluster as they represent only the most concentrated peaks in the luminous matter distribution. We have commenced an investigation into the larger scale distributions of galaxies throughout the entire SSC region and close environs using data from wide-field multi-fibre spectrographs such as the FLAIR-II system on the U.K. Schmidt Telescope (UKST) of the Anglo-Australian Observatory (Parker \\& Watson, 1995, Parker 1997). With such multiplexed facilities we are able to measure many more galaxy redshifts over large angular extents and obtain a more complete picture of the composition and disposition of galaxies in the entire supercluster.  A preliminary analysis based on 301 new, bright ($R\\leq 16$) galaxy redshifts obtained with FLAIR-II in a 77~deg$^{2}$ region of the supercluster has already been published by our group (Drinkwater et al.\\ 1999, hereafter D99). The measured galaxies were uniformly distributed over the selected inter-cluster regions, and most ($230 \\equiv 75\\%$) were found to be members of the supercluster. The D99 sample traced out two previously unknown sheets of galaxies which appeared to link various Abell clusters in the supercluster. We also found that in a 44 deg$^{2}$ sub-area of the supercluster {\\em excluding} the Abell clusters, these sheets alone represent an overdensity of a factor of $2.0\\pm0.2$ compared to a uniform galaxy distribution.  Within the initial survey area of D99 the new galaxies contribute an additional 50\\% to the known contents of the Shapley supercluster in that region, implying a corresponding increase in contribution to the motion of the Local Group. A much larger galaxy sample over a similar area of sky was presented by Quintana et al.\\ (2000). They measured nearly 3000 galaxy redshifts in an area of about $12\\deg \\times 6\\deg$. They estimated the upper bound on the mass of the central region and found the overdensity to be substantial, but still only able to contribute to a small fraction ($3\\Omega_m^{-0.4}$\\%) of the observed motion of the Local Group (Reisenegger et al.\\ 2000).  In this paper we present radial velocities for an additional 409 bright ($R\\leq16.5$) galaxies spread over an extended region to the East, West and South of the main SSC concentration. We analyse a combined magnitude-limited sample from this paper and D99 of 710 galaxies with measured redshifts, of which 464 (65 per cent) are members of the supercluster based on their redshift range ($8,000 < cz < 18,000\\kms$). They seem to link the two previously identified galaxy sheets found in D99 and also establish a connection with a third feature at $\\overline V$= 15000\\kms near R.A.=~13h. The new sample also tends to fill a previous gap in the galaxy distribution between the foreground Hydra-Centaurus region and the SSC. Future study of the SSC will cover even larger regions using data from the FLASH survey of Kaldare et al.\\ (2003) (also with the UKST/FLAIR-II system) and the 6dF survey of the southern sky currently in progress (Watson et al.\\ 2000).  The observed galaxy sample and observations are described in Section~2, the results are presented in Section~3 and a brief discussion of the significance of these new measurements is given in Section~4.  A full and detailed analysis and interpretation of the new redshifts from D99, this work and a compilation of all literature redshifts in the SSC is the subject of a separate paper (Proust et al., in preparation). Unless otherwise noted in this paper we use a Hubble constant of $H_0=75\\kmsMpc$ giving a distance of 200 Mpc to the main supercluster ($cz=15,000\\kms$) so 1 degree subtends a distance of 3.5 Mpc. ",
        "conclusions": "Our new observations of galaxies towards the Shapley supercluster have, by surveying a large contiguous region including areas away from known clusters, revealed substantial new, large, inter-connected structures. The original SSC region has now been shown to form part of a much larger structure extending uniformly in at least two sheets over the whole surveyed region principally to the South of the SSC core. We have detected an additional 426 members of the SSC in our combined (D99 and this paper) survey area, increasing very significantly the number of known SSC galaxies. Our new data suggest that the SSC is at least 50\\% more massive than previously thought with a significant part of the extra mass in the closer sub-region. The SSC must therefore has a more important effect on the Local Group that originally envisaged. We defer a detailed calculation of the size of this effect and a fuller analysis and interpretation of these new data, to a subsequent paper (Proust et al, in preparation) where some additional data will also be included.  We now plan to investigate regions well away from the SSC center between R.A. 12h30 to 14h30 and from $-23^{\\circ}$ to $-42^{\\circ}$ in order to: i) define the complete topology of the SSC, ii) analyze the individual Abell clusters contained in the SSC and iii) determine more precisely its gravitational effect on the Local Group. We have already obtained the necessary extended catalogue of galaxies for this larger region from additional red ESO/SRC SSC plates scanned by the MAMA machine. The new 6dF replacement for FLAIR-II at the UKST with its automated robotic fibre positioner and 150~fibres in a $6^{o}$ field will be an ideal facility to continue our study of the Shapley Supercluster."
    },
    "0402/astro-ph0402154_arXiv.txt": {
        "abstract": "{XMM-{\\it Newton} observations are presented for the $z=0.83$ cluster of galaxies  MS\\,1054$-$0321, the highest redshift cluster in the {\\it Einstein} Extended Medium Sensitivity Survey (EMSS). The  temperature inferred by the XMM-{\\it Newton} data, T$=$7.2$^{+0.7}_{-0.6}$ keV,  is much lower than the temperature previously reported from ASCA data, T$=$12.3$^{+3.1}_{-2.2}$ keV (Donahue et al. 1998), and a little lower than the Chandra temperature, T$=$10.4$^{+1.7}_{-1.5}$ keV, determined by Jeltema et  al. 2001. The discrepancy between the newly derived temperature and the previously  derived  temperatures  is  discussed in detail. If one allows the column density to be a free parameter, then the best fit temperature becomes T$=$8.6$^{+1.2}_{-1.1}$ keV,  and the best fit column density becomes N$_{H}=1.33^{+0.15}_{-0.14}\\times10^{20}$ atoms cm$^{-2}$. The iron line  is well  detected in the  XMM-{\\it Newton} spectrum with a value for  the abundance of Z$=$0.33$^{+0.19}_{-0.18}$ Z$_{\\sun}$, ~in very good agreement with previous determinations. The derived XMM X-ray luminosity for the overall cluster in the  2--10 keV energy band is $L_{X}=(3.81\\pm0.19) \\times10^{44}$ h$^{-2}$ erg s$^{-1}$ while the  bolometric luminosity is $L_{BOL}=(8.05\\pm0.40) \\times10^{44}$ h$^{-2}$ erg s$^{-1}$. The XMM-{\\it Newton} data confirm the substructure in the cluster X-ray morphology already seen by {\\emph  ROSAT} and in much more detail by Chandra. We find that only two of the three clumps detected in the weak lensing mass reconstruction image are visible in X-rays, as already noted by \\cite{jel01}. The central weak lensing  clump is coincident  with  the  main  cluster  component and has a temperature T$=$8.1$^{+1.3}_{-1.2}$ keV. The western weak lensing clump coincides with the western X-ray component which is much cooler with a temperature T$=$5.6$^{+0.8}_{-0.6}$ keV. The optically  measured velocity  dispersion, obtained from  145 cluster redshifts, is consistent  with the velocity dispersion  expected from the $\\sigma_{V}-T_{X}$ relationship once the XMM-{\\it Newton} temperature is used.  \\ms fits well in the  $\\sigma \\propto T^{1/2}$ correlation available in the literature and derived from information collected for all clusters  at redshfits of z$\\leq$1.27 known today and with a measured X-ray temperature. The cluster temperature seems to be commensurate with the  predictions from its X-ray luminosity from the L$_{X}-T_{X}$ relation of  local  clusters. Given the newly determined temperature, \\ms is no longer  amongst  the hottest clusters known. ",
        "introduction": "\\ms is the most luminous and distant cluster of galaxies in the EMSS (Gioia et al. 1990; Stocke et al. 1991). For some time it was the only very massive, high  redshift cluster of galaxies (Gioia \\& Luppino 1994), hence it was extensively studied and observed at  many  different  wavelengths. \\cite{don98} performed an X-ray analysis using both ASCA and ROSAT  and determined  a high  luminosity (L$_{BOL}=1.1\\times10^{45}h^{-2}_{100} ~erg~s^{-1}$) and  high gas temperature (T$=$12.3$^{+3.1}_{-2.2}$ keV). Using this temperature Donahue et al. estimated a virial mass of $\\sim7.4\\times10^{14} h^{-1}_{100} M_{\\sun}$  ~within a radius r$_{500}\\footnote{within a region whose density is 500 times the critical density.} = 1.5 h^{-1}_{100}$ Mpc. The X-ray emission of \\ms is resolved by the ROSAT/HRI in two, possibly  three clumps, plus an extended component indicating that the cluster is not relaxed and that some  caution must be adopted when determining  parameters from the X-ray data. \\cite{neu00} used the same ASCA temperature and an isothermal  $\\beta$ model fit to the ROSAT/HRI data to obtain a total mass of $1.6\\times10^{15} ~h^{-1}_{50} M_{\\sun}$ within a virial radius r$_{V} = 1.65 ~h^{-1}_{50}$ Mpc. The significant substructure observed in the HRI data  by the above investigators was confirmed by Jeltema et al. 2001 (J01) who used the superb resolution of the Chandra instruments. J01 give a new determination of the temperature, T$=$10.4$^{+1.7}_{-1.5}$ keV, lower than the ASCA temperature and measured, consequently, a lower virial mass of 6.2$\\times10^{14} ~h^{-1}_{100} ~M_{\\sun}$ within r$_{200}\\footnote{within a region whose density is 200 times the critical density.} = 0.76 ~h^{-1}_{100}$ Mpc. A similar temperature for \\ms was found by \\cite{joy01} using Sunyaev Zeldovich Effect data. They derive a temperature  T$=$10.4$^{+5}_{-2}$ keV and a total mass within 94\\arcsec ~(corresponding to 389 kpc in the cosmology adopted in this paper) equal to M($\\leq94'')=(4.6\\pm0.8)\\times10^{14} ~h^{-1}_{100} M_{\\sun}$. \\cite{vik02} used the Chandra instruments,  and derived an even lower value for the temperature, T$=$7.8$\\pm$0.6 keV, in agreement with the determination  by \\cite{toz03}  who analyzed the same Chandra data and obtained T$=$8.0$\\pm$0.5 keV. In the optical band an extensive study by  \\cite{tran99} and \\cite{vdok00}, based on 24 and 81 spectroscopic cluster member redshifts, respectively, finds agreement between the observed velocity dispersion ($\\sigma_V=1170\\pm$150 km s$^{-1}$; Tran et al. 1999) and the high  X-ray measured  ASCA temperature (T$=$12.3$^{+3.1}_{-2.2}$ keV; Donahue et al. 1998),  albeit within the large errors. From Hubble Space Telescope (HST) imaging, supplemented by ground telescope data, van Dokkum et al. (2000) detected  a large overdensity of red cluster galaxies with irregular and elongated distribution. A very high fraction (17 per cent) of the cluster galaxies are classified as ``mergers/peculiar'' on the basis of double nuclei, tidal tails and distorted morphologies (see van Dokkum et al. 1999; van Dokkum et al. 2000).  Hoekstra, Franx \\& Kuijken (2000)  performed a weak lensing study using a two-color  mosaic of deep WFPC2 images and determined  a mass of (1.2$\\pm$0.2$)\\times10^{15}$ $h^{-1}_{50} M_{\\sun}$ within an aperture radius 1~$h^{-1}_{50}$ Mpc, in good agreement with the previous estimators (Luppino \\& Kaiser 1997; Clowe et al. 2000). Assuming an isothermal mass distribution \\cite{hoek00}  found a corresponding velocity dispersion of 1311$^{+83}_{-89}$ km s$^{-1}$, again consistent with the X-ray and optical estimates. The weak lensing signal clearly shows three distinct clumps in the mass reconstruction distribution in agreement with the light distribution, that all appear to have similar masses.  An extremely deep 5-GHz survey of the cluster region  down to 32 $\\mu$Jy (6$\\sigma$; Best et al. 2000) detects a high number of radio sources (34 vs 25 expected from blank-field radio source counts). Surprisingly, no radio emission  is detected towards the merger galaxies found in the HST images by \\cite{vdok00} consistent with the hypothesis that low-luminosity radio sources may be triggered by initial weak interactions rather than direct mergers (Best et al. 2000). \\cite{joh03} detected an excess of point sources  in the archival deep Chandra image of the field of MS\\,1054$-$0321. Combined with identifications of cluster AGN  from the radio work of \\cite{best02}, \\ms seems to have a significantly enhanced AGN acitivity with respect to local galaxy clusters.  From all these data a consistent picture emerges of \\ms as a young, massive, highly luminous cluster with significant substructure both in the  X-ray and optical data, confirmed by the weak lensing mass reconstruction distribution. \\ms was also observed by XMM-{\\it Newton}.  The new data analysis and interpretation  are the subject of this paper. A new redshift determination from additional optical spectroscopic data is used in combination with the XMM-{\\it Newton} data to  provide a more accurate picture of  the cluster. Comparisons are  also made with previous analyses. Throughout the paper we use  a Hubbble constant of $H_{0} = 100 h$ km s$^{-1}$ Mpc$^{-1}$,  q$_{0}=0.5$ (for comparison with previous work on this cluster) and a cosmological constant of $\\Lambda=0$, except where noted. One arcmin corresponds to 248.6 h$^{-1}$ kpc at the cluster redshift.  \\begin{figure} \\centering \\caption{X-ray image of the MOS2 detector showing the extended emission of MS\\,1054$-$0321. This image is presented to highlight the four different regions analyzed in X-rays. Figure 1 can be seen in the pdf file of the paper at http://www.ira.cnr.it/~gioia/PUB/publications.html} \\label{Fig1} \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=10cm, angle=-90]{fig2_spectrum.ps} \\caption{Binned X-ray spectrum and residuals for MS\\,1054$-$0321. The spectrum was binned to have a minimum of 40 (50) counts per bin for each MOS (pn). The solid line represents the the best fit MEKAL model.} \\label{Fig2} \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=8cm, angle=-90]{fig3_chi_cont.ps} \\caption{68.3\\%, 90\\%, and 99\\% confidence $\\chi^{2}$ contours ($\\Delta\\chi=$2.30, 4.61, and 9.21) for the cluster iron abundance and temperature (point source to the south excluded).} \\label{Fig3} \\end{figure} ",
        "conclusions": "The main result from the new XMM-{\\it Newton} data is the lower measured value for the temperature of \\ms with respect to previous determinations. The XMM-{\\it Newton}  temperature, T$=$7.2$^{+0.7}_{-0.6}$ keV, is definitely lower than the ASCA temperature T$=$12.3$^{+3.1}_{-2.2}$ keV (Donahue et al. 1998), and also somewhat lower, even taking into account the uncertainties, than the  Chandra  temperature, T$=$10.4$^{+1.7}_{-1.5}$ keV, reported by J01. It is however  fully consistent with the Chandra temperature, T$=$7.8$\\pm$0.6 keV, published by \\cite{vik02} (see their Table 2), and  with the Chandra temperature, T$=$8.0$\\pm$0.5 keV, published by \\cite{toz03} (see their Table 3). The iron abundance inferred from the XMM-{\\it Newton} observations is found to be in very good agreement with previous  determinations by  both J01 (Z$=$0.38$\\pm$0.15 Z$_{\\sun}$) and by \\cite{toz03}  (Z$=$0.35$\\pm$0.07 Z$_{\\sun}$) after their abundances are rescaled to the \\cite{gs98}  values while Vikhlinin et al. (2002) fix the abundance to  Z=0.3 Z$_{\\sun}$. \\ms is in many aspects very similar  to other distant X-ray selected   clusters like RXJ\\,1717+67 at  $z=0.81$  (Gioia et al. 1999), or RXJ\\,0152.7$-$1357 at $z=0.83$  (Della Ceca et al. 2000; Ebeling et al. 2000) which show  elongated or double morphologies  in optical, weak lensing and X-rays and which have moderate temperature values.  While it is possible that ASCA included other sources in its larger field of view, the lower X-ray  temperature found for this cluster by this analysis, seems to be more  accurate given the agreement of three out of four measurements (this paper; Vikhlinin et al., 2002; Tozzi et al., 2003). In order to verify the correctness of our XMM-{\\it Newton} results we have analyzed  the Chandra archival data of \\ms ~and used all the information in the J01 paper to duplicate exactly their data reduction. There is a number of corrections to be made to the data. One of them, the ACISABS correction\\footnote{http://www.astro.psu.edu/users/chartas/xcontdir/xcont.html}, was not available at the time of the J01 analysis and was implemented later in the suggested steps for Chandra data reduction. This correction takes into account the quantum efficiency degradation at low energies, below 1 keV. The application of the ACISABS correction shifts the temperature towards lower values. There are several flares in the Chandra data, mostly visible only in the ACIS S3 back-illuminated chip, which contribute to raise the temperature. In the following analysis we use: a 1.5\\arcmin ~extraction radius at the same cluster center as J01, the 0.8--7 keV energy range  for the  spectral fit,  subtraction of  point sources including the one to the south at the same position as J01, local background estimate, use of the Anders \\& Grevesse abundance values, counts binned in intervals so as to have at least 20 counts per bin, fit with a Raymond-Smith (Raymond \\& Smith 1977) thermal plasma  model, N$_{H}$ fixed to  the Galactic value, temperature and abundance left free to vary. The two different parameters used are a lower exposure time of 67,478s (to cut out the time intervals affected by the flares) and application of the  ACISABS correction. We obtain a value for the Chandra temperature T$=$7.4$^{+1.4}_{-0.9}$ keV,  and an abundance Z$=$0.20$\\pm$0.14 Z$_{\\sun}$, fully consistent with our XMM-{\\it Newton} spectral result. If instead we release  the two corrections (ACISABS, lower exposure time) then we obtain the same values as  published by J01,  namely T$=$10.1$^{+1.7}_{-1.5}$ keV and  Z$=$0.22$\\pm$0.16 Z$_{\\sun}$. Adopting all the corrections and the same parameters as in \\cite{vik02} we obtain exactly their same value  for the temperature T$=$7.8$\\pm$1 keV. Thus we are confident that our XMM-{\\it Newton} data analysis is correct since the two different datasets, XMM-{\\it Newton} and Chandra, yield exactly the same result.  We check now how \\ms fits in the  L$_{X}-T_{X}$ relation of galaxy clusters. Understanding the evolution of the  L$_{X}-T_{X}$ relation is important not only to understand the  physics behind the formation of galaxy clusters, but also because it  provides a link between observations of clusters and derivation  of  cosmological parameters. The L$_{X}-T_{X}$ relation has been well studied at low redshift. There is no consensus yet on the evolution of the L$_{X}-T_{X}$ relation with redshift (see among others: Borgani et al. 2001; Holden et al. 2002; Novicki et al. 2002; Vikhlinin et al. 2002) probably due to the lack of large  samples of  galaxy clusters at cosmologically significant redshift. By comparing the bolometric luminosity of \\ms with the best fit relationship log(L$_{BOL,X})=(2.88\\pm0.15)$log(T/6keV)$+$(45.06$\\pm$0.03) by Arnaud \\& Evrard (1999), or using similar observationally determined relationships, one would expect for \\ms a temperature T$=$8.2$\\pm$0.15 keV which is well in agreement within the errors with the measured XMM-{\\it Newton} temperature.  We check next the $\\sigma_{V}-T_{X}$ relationship. We can estimate the velocity dispersion for  \\ms from the X-ray data and compare the result with the optical data. The velocity dispersion implied by the XMM-{\\it Newton} temperature   using the relation T$_{X}/\\mu m_{p} \\sigma^{2}$  with $\\beta=$1 (that is assuming energy equipartition between the galaxies and the gas) is   $\\sigma_{V}=1047^{+50}_{-44}$ km s$^{-1}$, consistent with the observed  optical velocity  dispersion within the errors, implying that the velocity dispersion reflects the temperature of the gas. Girardi et al. (1996) have derived a best fit relation between the velocity dispersion and the X-ray temperature equal to Log($\\sigma$)$=$(2.53$\\pm$0.04)+(0.61$\\pm$0.05)log(T), with more than 30\\% reduced scatter with respect to previous  work by taking into account distortions in the velocity fields, asphericity of the cluster or presence of substructures. Using the temperature T$=$7.2 keV of \\ms in the above relation, the resulting velocity dispersion value, $\\sigma_{V}=1129^{+66}_{-57}$ km s$^{-1}$,  is in very good agreement with the measured optical value,  $\\sigma_V=1153\\pm$80 km s$^{-1}$. The relation between temperature and velocity dispersion has been published by several authors. Last in order of time, this relation can be found in Lubin, Mulchaey \\& Postman (2003) where the authors have collected the information known today for all clusters at redshfits of z$\\leq$1.27 with a measured X-ray temperature. There is a big scatter in the $\\sigma_{V}-$T$_{X}$ relation (see their figure 4) but the average relation is consistent with $\\sigma \\propto T^{1/2}$ up to these  redshifts. We find that \\ms ~follows closely this correlation.  We can also estimate the mass of \\ms using the X-ray temperature value. With the assumptions that the mean density in the virialized region is $\\sim$200 times the critical density at the redshift of the cluster and that the cluster is  isothermal (Evrard, Metzler \\& Navarro 1996; Donahue et al. 1998), we can use the  scaling law method as illustrated in  Arnaud \\& Evrard (1999).  From the simulations of Evrard, Metzler \\& Navarro (1996) for the mass-temperature relation one can estimate the virial mass within a radius r$_{200}(T)=$1.85(T/10keV)$^{1/2}(1+z)^{-3/2}$ h$^{-1}$ Mpc by using the equation  M$_{vir}\\approx(1.45\\times10^{15} h^{-1} M_{\\sun})(1+z)^{-3/2}(kT_{X}/10 keV)^{1.5}$. From the XMM-{\\it Newton} temperature the virial mass is approximately M$_{vir}=(3.58^{+0.53}_{-0.44})\\times10^{14}$ M$_{\\sun}$ ~within r$_{200}=0.63 ~h^{-1}$ Mpc, where the uncertainties on the mass reflect the uncertainties on the temperature. This mass is  lower than the one derived by J01 $(M\\leq 0.76 ~h^{-1}_{100}$ Mpc = 6.2$\\times10^{14} ~h^{-1}_{100} ~M_{\\sun}$) because they use a higher Chandra temperature which, given the relation between r$_{200}$ and T, converts in a larger r$_{200}$ (their 0.76 $h^{-1}_{100}$ Mpc vs our 0.63 $h^{-1}_{100}$ Mpc). The virial mass derived from the XMM-{\\it Newton} temperature is  definitely lower than the mass derived from weak lensing by Hoekstra, Franx \\& Kuijken (2000) of 1.2$\\pm$0.2$ \\times10^{15}$ $h^{-1}_{50} M_{\\sun}$ who used a different aperture radius of 1~$h^{-1}_{50}$ Mpc. It is not possible with the present XMM-{\\it Newton} data to assume an aperture radius of 1~$h^{-1}_{50}$ Mpc (1.75\\arcmin ~using the Hoekstra et al. cosmology) since we are running into the gaps between the chips for the pn detector.  \\begin{figure} \\centering \\caption{21,600s I-band image of MS\\,1054$-$0321 taken at the 2.2m of the University of Hawai'i. The smoothed overlayed X-ray contours are in units of 3, 5, 7, 10, 15, 20, 30, 35, 40, 43, 45 $\\sigma$ over the background. The image measures $2048 \\times 2048$  pixels covering a field of $7'.5\\times7'.5$ ($1.86h\\times1.86h$ Mpc at $z =$ 0.83). Figure 7 can be seen in the pdf file of the paper at http://www.ira.cnr.it/~gioia/PUB/publications.html} \\label{Fig7} \\end{figure} Given the XMM-{\\it Newton} results, \\ms appears to be a more ``normal'' cluster with high, but not extreme  X-ray luminosity or temperature values. Its  temperature seems to be commensurate with the predictions from its X-ray luminosity from the L$_{X}-T_{X}$ relation of local clusters published in the literature.  The optically measured velocity dispersion is consistent with the velocity dispersion expected from the $\\sigma_{V}-T_{X}$  relationship.  Given the clear evidence for substructure in the cluster morphology, indicative of a still unrelaxed cluster, it is surprising that the global  properties of \\ms fit so well in any  local cluster correlation."
    },
    "0402/astro-ph0402632_arXiv.txt": {
        "abstract": "Ideal hydrodynamic models of the intracluster medium (ICM) in the core regions of galaxy clusters fail to explain both the observed temperature structure of this gas, and the observed morphology of radio-galaxy/ICM interactions.  It has recently been suggested that, even in the presence of reasonable magnetic fields, thermal conduction in the ICM may be crucial for reproducing the temperature floor seen in many systems.  If this is indeed correct, it raises the possibility that other transport processes may be important. With this motivation, we present a numerical investigation of the buoyant evolution of AGN-blown cavities in ICM that has a non-negligible shear viscosity.  We use the ZEUS-MP code to follow the 3-dimensional evolution of an initially static, hot bubble in a $\\beta$-model ICM atmosphere with varying degrees of shear viscosity.  With no explicit viscosity, it is found that the combined action of Rayleigh-Taylor and Kelvin-Helmholtz instabilities rapidly shred the ICM cavity and one does not reproduce the intact and detached ``ghost cavities'' observed in systems such as Perseus-A.  On the other hand, even a modest level of shear viscosity (corresponding to approximately 25\\% of the Spitzer value) can be important in quenching the fluid instabilities and maintaining the integrity of the bubble.  In particular, we show that the morphology of the NW ghost cavity found in Perseus-A can be reproduced, as can the flow pattern inferred from the morphology of H$\\alpha$ filaments.  Finally, we discuss the possible relevance of ICM viscosity to the fact that many of the active ICM cavities (i.e., those currently associated with active radio-lobes) are not bounded by strong shocks, the so-called ``shock problem''. ",
        "introduction": "Recent years have seen a growing realization that the cores of rich galaxy clusters are complex and dynamic environments.  In particular, it is becoming clear that the radio-loud active galactic nuclei (AGN) often hosted by the cD galaxies in rich clusters can have a major influence on the hydrodynamics and thermodynamics of the core regions of the intracluster medium (ICM).  As we discuss below, the rich datasets coming from the {\\it Chandra X-ray Observatory} and {\\it XMM-Newton} now demand theoretical models that go beyond the simple picture of jet-blown ``bubbles'' rising in a static ICM described by ideal hydrodynamics.  Bulk ICM motions (including turbulence), magnetohydrodynamics (MHD), thermal conductivity, and viscosity may all be relevant to data that are currently being taken.  Even before the launches of {\\it Chandra} and {\\it XMM-Newton}, {\\it Einstein} and {\\it ROSAT} studies revealed prominent radio-galaxy/cluster interactions in three clusters; Perseus-A (B\\\"ohringer et al. 1993; Heinz, Reynolds \\& Begelman 1998), Virgo-A (Feigelson et al. 1987; B\\\"oringer et al.  1995) and Cygnus-A (Carilli, Perley \\& Harris 1994; Harris, Carilli \\& Perley 1994). {\\it ROSAT} showed Perseus-A and Cygnus-A to possess intracluster medium (ICM) cavities coincident with the prominent radio-lobes in these two sources, suggesting supersonic inflation of a bubble in the ICM by the jetted AGN (Clarke, Harris \\& Carilli 1997).  Virgo-A, on the other hand, showed a cooler (and soft X-ray brighter) region of ICM associated with the outer eastern ``ear'' seen in low-frequency radio maps of Virgo-A (Owen, Eilek \\& Kassim 2000).  It was first suggested by B\\\"ohringer et al.  (1995) that this phenomenon might be caused by lower entropy gas from the cluster center being dragged upwards in the ICM atmosphere by a buoyantly rising radio-lobe.  More recent observations by {\\it Chandra} and {\\it XMM-Newton} show that radio-galaxy induced ICM substructure is surprisingly ubiquitous and complex.  The basic results described above, i.e., the existence of ICM cavities associated with active radio-lobes and the presence of cool material that appears to lie in the wake of old, buoyantly rising radio-lobes, have been confirmed in numerous systems (e.g., Hydra-A; McNamara et al. 2000, Abell~2052; Blanton et al. 2001, Virgo-A; Young, Wilson \\& Mundell 2002, Perseus-A; Fabian et al. 2000, 2003a). However, these new data have raised several mysteries and are increasingly at odds with simple models for radio-galaxy/ICM interactions.  Firstly, in any ideal hydrodynamic model, the cavities must be inflated supersonically or else they would be destroyed by Rayleigh-Taylor (RT) instabilities faster than they are inflated. Curiously, the strong and hot ICM shocks that one expects to find around the active cavities are notably absent; instead, many ICM cavities are surrounded by ICM shells that are {\\it cooler} than the ambient ICM.  We shall refer to this as the ``shock problem''. Secondly, ICM cavities that are not associated with any obvious radio-lobe (``ghost cavities'') have been discovered.  Examples are found in the Perseus cluster (Fabian et al. 2003a), Abell~2597 (McNamara et al. 2001), and Abell~4059 (Heinz et al. 2002, Choi et al. 2004).  In some cases, ghost cavities are coincident with regions of low-frequency (74\\,MHz) radio emission supporting the hypothesis that they correspond to old radio-lobes from previous and now extinct AGN outbursts (Fabian et al. 2002a).  Interestingly, as we will explicitly demonstrate in this paper, ideal hydrodynamic models fail to reproduce the observed morphology of at least some ghost cavities.  Finally, most clusters have been found to possess a ``temperature-floor'' in the sense that the radiative cooling of the ICM appears not to proceed below temperatures of $kT\\sim 1-2\\keV$ (Tamura et al. 2001; Peterson et al.  2001).  Again, no clear explanation for this fact is provided by an ideal hydrodynamic model.  Br\\\"uggen \\& Kaiser (2002) suggest that stirring of the ICM core by a central radio galaxy may be responsible for the temperature floor; however, such a scenario only postpones rather than prevents radiative cooling, and is hard to reconcile with the strong metallicity gradients observed in the cores of some clusters (David et al. 2001, Matsushita et al. 2002; Sanders \\& Fabian 2002).  This has led several authors (Narayan \\& Medvedev 2001; Fabian et al.  2002b; Voigt et al. 2002; Kim \\& Narayan 2003) to resurrect the notion that thermal conduction may be important in determining the thermodynamics of ICM cores.  Faced with the multiple failures of simple hydrodynamic models for cluster cores, we must carefully examine the other physical processes that might be relevant.  Both the dynamical and thermodynamical effects of magnetic fields are largely unexplored and are the subject of on-going large-scale simulations.  However, the qualitative success of models including thermal conduction in solving the temperature floor problem begs a study of other transport processes and, in particular, the effects of shear viscosity on the dynamics of a radio-galaxy/ICM interaction.  \\begin{figure} \\centerline{ \\psfig{figure=fig1.ps,width=0.42\\textwidth} } \\caption{X-ray image of the core regions of the Perseus cluster from the {\\it Chandra X-ray Observatory} (blue; from Fabian et al. 2003a) overlaid with the H$\\alpha$ image from the WIYN telescope (Fabian et al. 2003b; Conselice et al. 2001).  The flattened X-ray cavity is clearly visible in the central regions of this image. Furthermore, the H$\\alpha$ emitting filaments display well-defined arcs, suggestive of a vortex-like flow pattern in the region behind the buoyant cavity.  This image is 4.33\\,arcmin (96\\,kpc) on a side and is oriented such that North is upwards.} \\label{fig:perseus} \\end{figure}  Perseus-A and the core of the Perseus cluster continues to be one of the best studied examples of a rich cluster core and a complex radio-galaxy/cluster interaction.  The most detailed X-ray investigation to date, based on a deep (200\\,ks) observation by {\\it Chandra}/ACIS-S, has been presented by Fabian et al. (2003a).  These authors report the discovery of wave-like disturbances in the ICM on spatial scales of $\\sim 50\\kpc$, approximately twice the spatial scales of the most obvious ghost cavity to the north-west of Perseus-A.  They discuss a scenario in which viscous dissipation of these disturbances may act as a significant heat source for the ICM core.  It is shown that, provided viscosity operates reasonably close to its ideal unmagnetized value, it is possible for viscous dissipation of radio-galaxy induced disturbances to balance radiative cooling of the ICM.  This suggestion has been supported by recent simulation work by Ruszkowski, Br\\\"uggen \\& Begelman (2003). Circumstantial evidence for the presence of significant ICM viscosity is also provided by an examination of the morphology of H$\\alpha$ filaments.  Several of the filaments appear to trace well-defined arcs in the region below the ghost cavity (Fig.~\\ref{fig:perseus}; also see Fabian 2003b).  This argues against the presence of strong turbulence in the ICM core, possibly resulting from the action of viscosity.  If we make the stronger presumption that the H$\\alpha$ filaments follow streamlines in the ICM, the morphology of the filaments suggests the existence of a vortex ring within the ICM just below the NW ghost cavity.  With this background and motivation, this paper presents hydrodynamic simulations of the buoyant evolution of an AGN-blown cavity in a viscous ICM.  In order to bring clarity to the discussion of such a complex system, this paper deals with the focused question of how ICM viscosity effects the observed morphology and associated flow patterns of old (ghost) cavities.  Detailed investigations of the effects of viscosity on the growth of active cavities and the thermodynamic state of the ICM will be addresses in future work.  Section~2 reviews the importance of viscosity in typical clusters, and touches upon the robustness of viscosity in the presence of magnetic fields.  Section~3 describes the basic set-up of our simulations, as well as our results on the morphology and flow patterns.  Section~4 discusses some of the limitations of this work, and possible implications of ICM viscosity. Finally, conclusions are presented in Section~5. ",
        "conclusions": "The true role of kinematic viscosity in ICM/radio-galaxy interactions remains unclear, primarily due to the great uncertainty associated with magnetic suppression of Spitzer-type transport processes in the plasma.  However, in stark contrast to the inviscid models, our models of bubble evolution in modestly viscous ICM can reproduce both the morphology and inferred flow patterns seen around the best studied of the ghost ICM cavities, the NW cavity in Perseus-A.  The principal effect of viscosity is to stabilize the bubble against both RT and KH instabilities, thereby allowing it to remain intact as it floats upwards in the ICM atmosphere.  The flattening of such bubbles as they rise is a natural explanation for the elongated morphology of the NW cavity in Perseus-A.  The ``smoke-ring'' like flow pattern below the NW cavity inferred from the morphology of the H$\\alpha$ filaments can then arise due to minor fragmentation of the rising bubble (leading to a trailing torus) or genuine vortex shedding.  If the ICM is indeed viscous at the level treated by our models, the subsequent slowing of the evolution timescales for ICM/radio-galaxy interactions may be the solution to the ``shock problem'', i.e. the lack of strong shocks bounding the ICM cavities that are associated with active radio-lobes.  In this scenario, radio galaxies such as Per-A would be rather less powerful than previously thought (Heinz, Reynolds \\& Begelman 1998), inflating their associated ICM bubbles subsonically.  If the evolution is sufficiently slow, radiative cooling might allow the ICM shell surrounding the cavity to cool, again in agreement with Chandra observations of Perseus-A.  Ultimately, disentangling the effects of transport processes (viscosity and thermal conduction) and magnetic fields will require further simulation work coupled with detailed observations.  On the theoretical side, it is important to initiate a new generation of simulations which incorporate full MHD {\\it as well as} transport processes.  Qualitatively new and rich dynamics becomes possible upon the inclusion of transport processes within MHD.  For example, Balbus (2000) discovered the ``magnetothermal instability'' which occurs in any MHD atmosphere with a temperature profile that decreases with height provided there is thermal conduction along the magnetic field lines. In addition, the viscous transport coefficients will be anisotropic in the presence of a magnetic field, and this may lead to yet further complexity in the dynamics.  Such instabilities could well be crucial for understanding the dynamics of radio-galaxy/ICM interactions.  Observationally, deep (megasecond) {\\it Chandra} observations of the nearest radio-galaxy/cluster interactions might allow detailed imaging spectroscopy of fluid instabilities associated with the radio-galaxy/ICM interaction.  Furthermore, upon the launch of {\\it Astro-E2} in 2005, we will be able to directly probe turbulent velocities within the ICM of nearby (bright) clusters through broadening of their X-ray emission lines.  With the caveat that full MHD viscous simulations have yet to be done, the presence or absence of turbulence in the ICM, especially in regions well separated from potential driving sources (such as AGN or merging subclusters) could be an important indicator of the presence of viscosity."
    },
    "0402/astro-ph0402318_arXiv.txt": {
        "abstract": "We present three {\\it Far Ultraviolet Spectroscopic Explorer (FUSE)\\/} observations of the Narrow-Line Seyfert~1 galaxy NGC\\,4051. The most prominent features in the far-ultraviolet (FUV) spectrum are the \\ion{O}{6} emission and absorption lines and the \\ion{H}{1} Lyman series absorption lines which are detected up to the Lyman edge. We also identify weak emission from \\ion{N}{3}, \\ion{C}{3}, and \\ion{He}{2}. The \\ion{C}{3} line shows absorption while none is detected in the \\ion{N}{3} and \\ion{He}{2} lines.  In \\ion{H}{1} and \\ion{C}{3} we detect two main absorption systems at outflow velocities of $-50\\pm30$ and $-240\\pm40$ \\kms , as well as a possible third one at $\\sim -450$ \\kms . These systems are consistent in velocity with the 10 absorption systems found previously in \\ion{C}{4}, \\ion{N}{5}, and \\ion{Si}{4}, though the individual systems are blended together in the FUV spectrum. We estimate column densities of the two main absorption systems and find that the \\ion{H}{1} column density is lower for systems with larger outflow velocity. We detect no flux or spectral variations of NGC\\,4051 at FUV wavelengths during three epochs spanning one year. This is consistent with the optical light curve which shows no variations between the three epochs. It is also consistent with the X-ray light curve which shows consistent flux levels at the three epochs of the {\\it FUSE\\/} observations, although the X-ray light curve shows strong variations on much shorter timescales. ",
        "introduction": "The Narrow-Line Seyfert~1 (NLS1) galaxy NGC\\,4051 is one of the best-studied active galactic nuclei (AGN) across the electromagnetic spectrum due to its brightness and proximity ($V\\approx 13.5$; $z=0.002295\\pm 0.000043$ from optical emission lines; de Vaucouleurs et~al. 1991). It has been part of many studies since it was first mentioned by Hubble (1926) and classified by Seyfert (1943). Collinge et~al. (2001) recently presented the first high-resolution X-ray and ultraviolet (UV) spectra of NGC\\,4051. The {\\it Chandra\\/} High-Energy Transmission Grating Spectrometer (HETGS) spectrum reveals absorption and emission lines from hydrogen-like and helium-like ions of O, Ne, Mg, and Si. Two distinct blueshifted absorption systems were detected: a high-velocity system at $-2340\\pm130$ \\kms\\ and a low-velocity system at $-600\\pm130$ \\kms. In a {\\it Chandra\\/} Low-Energy Transmission Grating Spectrometer (LETGS) observation taken $\\sim 2$~yr later, van der Meer et~al. (2003) claim to detect a new X-ray absorption system at $\\sim -4500$ \\kms\\ while the $-2340$ \\kms\\ system is barely visible. An even higher velocity outflowing absorption system at $\\sim -6500$ \\kms\\ or $\\sim -16500$ \\kms\\ (depending on the identification of an absorption line at 7.1 keV in an {\\it XMM-Newton} EPIC observation) was suggested recently by Pounds et al. (2003). The UV spectrum taken with the Space Telescope Imaging Spectrograph (STIS) shows strong absorption mainly in the lines of \\ion{C}{4}, \\ion{N}{5}, and \\ion{Si}{4}. Ten intrinsic UV absorption systems are seen with velocities between $-650$ and 30 \\kms\\ and FWHMs ranging from $\\sim 30$ to $\\sim 160$ \\kms .  While the low-velocity X-ray absorption system is consistent in velocity with many of the UV absorption systems, the high-velocity X-ray absorption seems to have no UV counterpart.  NGC\\,4051 shows rapid and large-amplitude X-ray flux variability by up to an order of magnitude on timescales of hours (e.g., Marshall et al. 1983; Uttley et~al. 2000; Collinge et~al. 2001; McHardy et al. 2003; and references therein). Uttley et~al. (2000) found a strong correlation between the variability in the extreme UV (124--188~eV) and X-ray (2--10~keV) bands, and they suggested this to indicate that both bands are sampling the same power-law continuum. Shemmer et~al. (2003) find the variable optical continuum flux to correlate with the variable X-ray flux. The optical continuum leads the X-ray by $2.4\\pm1.0$ days. They also found that the variation amplitude in the optical is much smaller than in the X-ray, as seen in other NLS1s.  In this paper we bridge between the X-ray and near UV high-resolution spectra by presenting the first {\\it Far Ultraviolet Spectroscopic Explorer (FUSE)\\/} spectra of NGC\\,4051.  The {\\it FUSE\\/} spectra cover the 900--1180\\,\\AA\\ band with a resolution of $\\sim 0.06$\\,\\AA\\ (corresponding to $\\sim 20$~km~s$^{-1}$) . Our main proposed goals were (1) to study the FUV kinematic counterparts to the known X-ray and UV absorption lines, (2) to study any FUV absorption-line variability, and (3) to relate any FUV variability to that seen at X-ray and other wavelengths. The observations and data reduction are detailed in \\S~2. In \\S\\S~3 and 4 we present the temporal and spectral analyses, and in \\S~5 we discuss our results.  Finally, \\S6 presents a summary of our main results. ",
        "conclusions": "\\subsection{Far Ultraviolet Emission and Absorption} \\label{fuvem}  In these first FUV spectra of NGC\\,4051, we detect emission from \\ion{O}{6}, \\ion{N}{3}, \\ion{He}{2}, and \\ion{C}{3}. The \\ion{O}{6} emission lines are fitted with two Gaussians with FWHMs of $\\sim 700$ and $\\sim 4000$ \\kms ; the overall line profiles are significantly broader than for the optical lines.  This is in agreement with the majority of other NLS1s which show broad wings (several thousands of \\kms ) on their high-ionization UV lines (e.g., Rodriguez-Pascual, Mas-Hesse, \\& Santos-Lleo 1997). We also find the Gaussian centroids to be blueshifted by $-200$\\,\\kms\\ relative to the optical lines. This is again in agreement with the findings for other AGN (e.g., Richards et al. 2002 and references therein).  We detect absorption on the blue wings of the \\ion{O}{6} and \\ion{C}{3} emission lines, as well as absorption in the \\ion{H}{1} Lyman series. In \\ion{H}{1} we identify two main absorption systems and possibly a third at outflow velocities of $-50\\pm30$, $-240\\pm40$, and $\\sim -450$ \\kms . These three absorption systems are also present in \\ion{C}{3} (Figure~\\ref{o6velspec}c). The three FUV systems correspond in outflow velocity to UV systems 7--8, 5, and 2--3 which were found by Collinge et~al. (2001) in \\ion{N}{5}, \\ion{Si}{4}, and  \\ion{C}{4}. Thus, it is plausible that each of the three systems detected in the FUV is a blend of several narrower systems which cannot be resolved by {\\it FUSE}.\\footnote{We note that UV system 5 (and perhaps other systems) of Collinge et~al. (2001) is saturated and broad; hence it might well be a blend of several systems which are intrinsically blended (rather than just blended due to instrumental resolution).}  An interesting relation is present between the strengths of the absorption systems and their outflow velocities. The absorption system with the largest outflow velocity ($\\sim -450$ \\kms ) seems to have the weakest absorption, and it is present only in the lines up to about Ly$\\zeta$, i.e., it has a relatively low column density. The absorption system with the smallest outflow velocity ($\\sim -50$ \\kms ) has the strongest absorption with relatively high column density ($2.2^{+1.3}_{-1.1}\\times10^{17}$ \\cmii ), and all the lines that can be detected within the limit of the instrumental resolution (up to Ly$\\xi$) are saturated. The intermediate-velocity absorption system (at $\\sim -250$ \\kms ) has an intermediate absorption depth, detected only up to Ly$\\lambda$, and it has a column density of $1.0^{+0.6}_{-0.5}\\times10^{16}$ \\cmii .  This relation might be explained if the systems with higher velocities are closer to the central radiation source and hence are in a more intense radiation field which yields less \\ion{H}{1} in these systems.  This might indicate that the outflow is decelerating so that its velocity decreases as its distance from the ionizing source increases (a decelerating outflow was also found recently in NGC\\,3783; Gabel et al. 2003).  In the lowest velocity outflow system, which is highly saturated, we detect on the trough's wing from 100 to $-$100 \\kms\\ the effect of a velocity dependent covering factor (see Figure~\\ref{o6velspec}a), i.e., the absorption profile is produced by a different covering factor for each outflow velocity and not by the internal Voigt profiles of the lines (e.g., Arav, Korista, \\& de Kool 2002).  We can use the multi-epoch observations to put some constraints on the line-of-sight acceleration of the absorbers.\\footnote{This might not be the physical acceleration of the absorbing material, which could be crossing our line of sight to the nucleus in a ``standing'' pattern.} First we use the {\\it FUSE} observations which are separated by 354.4 days in the rest frame. An upper limit on the absorption profile difference between the two epochs is 30 \\kms\\ (this is the trough location measurement uncertainty; here we used the \\ion{O}{6}, \\ion{C}{3}, and Ly$\\gamma$ to Ly$\\theta$ absorption lines). Thus, the upper limit on the acceleration of the absorber is $9.8\\times 10^{-7}$ \\acc . Next we used the time between the STIS observation and the last {\\it FUSE} observation which is 1086.9 days in the rest frame. From the observed consistency of 100 \\kms\\ between the two (e.g., Figure~\\ref{o6velspec}b), we derive an upper limit on the absorber's acceleration of $1.1\\times 10^{-6}$ \\acc . The upper limit we find is at the lower range of the UV-outflow deceleration discovered recently in NGC\\,3783 of $\\sim (1$--$2.5)\\times 10^{-6}$ \\acc\\ (Gabel et al. 2003).  The \\ion{O}{6} absorption we detect from NGC\\,4051 is an example of the ``smooth'' absorption morphology described in Kriss (2002). This morphology includes objects where the \\ion{O}{6} absorption is so broad and blended  that individual \\ion{O}{6} components cannot be identified.  Another NLS1 with a ``smooth'' \\ion{O}{6} absorption morphology is Ark\\,564 (Romano et~al.  2002). Other NLS1s have ``single'' morphology (like Ton\\,S180 and I\\,Zw\\,1) or a ``blend'' morphology (like Mrk\\,478; see Kriss 2002 for details). Thus, NLS1s appear to be similar in their range of \\ion{O}{6} absorption-trough morphologies to broad-line Seyfert 1s.  The deep absorption troughs imprinted on the \\ion{O}{6} and \\ion{C}{3} broad emission lines indicate these absorbers reside at a distance from the AGN center larger than the corresponding broad line region (BLR) size. In NGC\\,4051 the BLR size has been measured, using optical Balmer lines, to be $3\\pm 1.5$ light days ($[7.8\\pm3.9]\\times 10^{15}$ cm; Shemmer et~al 2003, see also Peterson et al. 2000). The BLR size for the UV lines (higher ionization lines) is known to be smaller than that of the optical lines (lower ionization lines) by a factor of $\\sim$ 2--3 (e.g., Peterson \\& Wandel 2000).  This is attributed to the fact that the ionizing flux density is higher closer to the central source. Taking these considerations into account, we estimate a lower limit for the distance of the FUV absorber from the central continuum source to be $\\sim 2\\times 10^{15}$ cm.  The amount of nuclear FUV/UV light scattered around the absorbing material and into our line of sight is practically zero since the absorption troughs of the low-ionization ions (\\ion{H}{1}, \\ion{C}{3}, \\ion{C}{4}, and \\ion{N}{5}) reach saturation at zero flux. We find the higher ionization ion \\ion{O}{6} to have a covering factor of $\\sim$~0.9--0.95. This might indicate that the higher ionization absorbers have a smaller covering factor than the low-ionization absorbers.  \\subsection{Ultraviolet --- X-ray Connection}  Several studies have suggested a link between the UV and X-ray absorbers in AGN (e.g., Mathur, Elvis, \\& Wilkes 1995; Crenshaw et~al. 1999). The data presented in this paper enable us to examine this connection in NGC\\,4051 .  In Figure~\\ref{complines} we compare the \\ion{N}{5}~$\\lambda$1238 absorption from the STIS spectrum, the \\ion{O}{6}~$\\lambda$1038 absorption from the {\\it FUSE\\/} spectrum, and combined line profiles of He-like ions and H-like ions from the {\\it Chandra\\/}\\,HETGS data (Figure 3 of Collinge et~al. 2001). The resolution and S/N of the HETGS spectrum do not allow precise measurement of the X-ray absorption systems. The velocity shifts and FWHMs of the FUV and UV absorption systems are consistent with the low-velocity ($-$600 \\kms ) X-ray absorber which is seen in both the He-like and H-like lines.\\footnote{ We note that the profile of the low-velocity absorber has its main trough between $\\sim~ -500$ and $-900$ \\kms\\ in X-ray H-like lines, whereas in X-ray He-like lines it spans a wider range from $\\sim -100$ to $-900$ \\kms\\ (see Figures~\\ref{complines}c and \\ref{complines}d). The more highly ionized absorber may be more concentrated at higher outflow velocities, or alternatively emission lines may be filling in the low-velocity part of the H-like absorption trough.} Both the He-like and the H-like profiles seem to extend to higher blueshift velocities than the FUV and UV absorbers; however, given the HETGS resolution and the poor S/N of the X-ray spectrum, it is not possible to determine the significance of this effect. We do not identify in the FUV and UV absorption lines the high-velocity absorption system (at $-$2340 \\kms ) observed in the X-ray spectrum by Collinge et~al. (2001). In particular, we do not detect this system in \\ion{O}{6}. The high-velocity absorption system in the X-ray spectrum was detected only in the H-like ions of Si, Mg, Ne and O; it is only marginally seen in the He-like ions of these elements (Figure~\\ref{complines}c). Thus, it is not implausible that for the lower ionization ion \\ion{O}{6} we do not detect this absorption system. Further observations, probably with {\\it Chandra\\/}, are required to investigate the putative $-2340$~km~s$^{-1}$ absorption system further.  \\begin{figure} \\centerline{\\includegraphics[width=8.5cm]{Kaspi.fig6.eps}} \\caption{Line velocity profiles. (a) The \\ion{N}{5}~$\\lambda$1238 absorption from the STIS spectrum. (b) The \\ion{O}{6}~$\\lambda$1038 absorption from the {\\it FUSE\\/} spectrum. (c) and (d) are the combined line profiles of He-like and H-like ions, respectively, from the {\\it Chandra\\/} HETGS data binned to 200 \\kms\\ (Figure 3 of Collinge et al. 2001). The dashed vertical lines show 0, $-600$, and $-2340$ \\kms\\ velocities. The wide absorption trough around $-2000$ \\kms\\ in panel (b) is from \\ion{O}{6}~$\\lambda$1032. Airglow geocoronal emission lines and interstellar Galactic absorption lines are marked with~$\\oplus$. \\label{complines} } \\end{figure}  McHardy et al. (1995) and Nicastro et~al. (1999) modeled {\\it ROSAT\\/} observations of NGC\\,4051 and claimed significant ionized X-ray absorption variability. Our observations, in contrast, fail to find any significant FUV absorption variability (though the limited temporal sampling might have caused a coincidence of observing on times when conditions were similar).  This initially appears surprising if the UV and X-ray absorbers are connected. One possible explanation is that the low-resolution {\\it ROSAT\\/} data were confused by complex X-ray spectral variability; changes in the power-law continuum and/or soft excess could have created a false impression of ionized absorption variability. Alternatively, this could be attributed to FUV line saturation. Since most of the FUV absorption lines we detect appear saturated, changes in the absorber might lead to little change in the absorption lines. A possible exception is the \\ion{C}{3} absorption which does not seem to be heavily saturated; the lack of apparent \\ion{C}{3} absorption variability supports the idea that the absorber does not change substantially on time scales of months.  \\subsection{The Spectral Energy Distribution of NGC\\,4051}  In Figure~\\ref{sed} we show a spectral energy distribution (SED) of NGC\\,4051 from the infrared to the X-ray. All data were corrected for the Galactic absorption (see caption for details), but no intrinsic reddening correction was applied.  In the radio region (which is not shown in the figure) the object is radio quiet with $\\nu F_\\nu \\approx 10^{-16}$--$10^{-15}$ \\ergcms . The data we present are not contemporaneous, but overall most of the data represent an average over time of the object's flux (see the figure caption for details). The SED resembles the one presented in Komossa \\& Fink (1997), although here the SED is more detailed and has larger wavelength coverage.  \\begin{figure*} \\centerline{\\includegraphics[width=17.5cm]{Kaspi.fig7.eps}} \\caption{The SED of NGC\\,4051. Optical data are the average spectrum from Shemmer et al. (2003). {\\it IUE} data are mean archive spectra taken over $\\sim 10$ yr. The {\\it FUSE} data are the average spectrum from this work. {\\it EUVE} data are from Uttley et al. (2000), and the error bar represents the variability range. STIS and HETGS data are from Collinge et al. (2001); the HETGS data represent an average of the high and low states. The LETGS data are reported by van der Meer et al. (2003) and were taken while NGC\\,4051 was in a lower X-ray flux state.  The spectra were heavily binned for clarity.  The infrared data were taken from the NASA Extragalactic Database (NED) and were obtained by {\\it IRAS} in the far-IR and ground-based observatories in the near-IR. The data were corrected for the Galactic absorption, even though it is low. We used $E(B-V)=0.013$ mag and the extinction curve of Cardelli, Clayton, \\& Mathis (1989) for the UV to IR spectra, and $N_{\\rm H}=1.3\\times 10^{20}$ \\cmii\\ (Elvis, Wilkes, \\& Lockman 1989) with the WABS model in XSPEC (Morrison \\& McCammon 1983) for the X-ray spectra. No intrinsic reddening correction was applied. \\label{sed} } \\vglue -0.5cm \\end{figure*}  The mean spectrum of NGC\\,4051 from the FUV to the optical is represented by data from four instruments and is averaged over timescales of years. The data from all four instruments are in good agreement (notable is the agreement of the STIS spectrum which is at a consistent flux level with the averaged {\\it IUE} spectrum observed one-to-two decades before). The optical continuum of NGC\\,4051 varies little (e.g., Shemmer et~al. 2003), the UV continuum is also consistent with little variation, and the current work shows that NGC\\,4051 did not vary significantly in the FUV. Nevertheless, there are notable variations in the optical line fluxes (e.g., Shemmer et~al. 2003). Since the optical-to-FUV continuum does not vary much, a plausible conclusion is that the line variations are driven by the rapid and large soft X-ray and extreme UV variations and are not driven by the near or far UV as suggested before for this object.  The observed SED of NGC\\,4051 can be fitted with a power law ($f_\\nu \\propto \\nu^{\\alpha}$)  in limited spectral ranges.  We find $\\alpha$ to be $-1.46\\pm 0.01$ in the optical range (5500--7500~\\AA ) and $-2.03\\pm 0.02$ in the UV range (1160--3000~\\AA ). The FUV cannot be fitted with a power law as it is contaminated by emission and absorption, though Figure~\\ref{sed} suggests the UV slope continues into the FUV. In the X-ray range we find from the HETGS spectrum in the range 2--6 keV that $\\alpha = -0.84\\pm0.04$. In Table~\\ref{luminosity} we give the luminosity of NGC\\,4051 in several bands assuming $H_{0}=70$ km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_{M}=0.3$, and $\\Omega_{\\Lambda}=0.7$.  \\begin{deluxetable}{ccc} \\tablecolumns{3} \\tablewidth{0pt} \\tablecaption{Luminosity of NGC 4051 in Several Bands \\label{luminosity}} \\tablehead{ \\colhead{Wavelength [\\AA ]} & \\colhead{Data used} & \\colhead{Luminosity [$10^{40}$~erg\\,s$^{-1}$]}  } \\startdata 4000--7500   & Optical  & 61  \\\\ 1300--3300   & {\\it IUE}      & 37  \\\\ 1200--1700   & STIS     & 11  \\\\ \\phn{950--1150}& {\\it FUSE} & 3.8 \\\\ \\phn{6.2--24.8} & HETGS & 16 \\\\ 1.9--6.2 & HETGS & 13 \\enddata \\tablecomments{Luminosity is computed using a cosmology of $H_0=70$ km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_{M}=0.3$, and $\\Omega_{\\Lambda}=0.7$, which results in a luminosity distance of $9.8\\times10^6$ pc to NGC\\,4051. \\vglue -0.5cm } \\end{deluxetable}  Overall the SED shows a gradual rise in the emitted energy from the hard X-rays to the far-IR. Most of the observed energy from NGC\\,4051 is seen in the far-IR, while the FUV emission is about an order of magnitude weaker. The UV and FUV spectra do not show indications for a big blue bump. Such an observed SED could be explained by dust along the line of sight which reddens the optical-to-FUV spectrum and re-emits energy in the far-IR. This would be similar to the observed SED of Ark\\,564 in which significant reddening was found (Crenshaw et~al. 2002; Romano et~al. 2003). However, data from {\\it EUVE\\/} and {\\it ROSAT\\/} strongly constrain the amount of cold gas along the line of sight that could be associated with the putative dust (e.g., McHardy et~al. 1995; Uttley et~al. 2000). One solution to this difficulty could be to invoke an ionized absorber with embedded dust (e.g., Brandt, Fabian, \\& Pounds 1996; Kraemer et~al. 2000), although even in this case there is little evidence at present for X-ray absorption edges from the dust grains themselves (e.g., Komossa \\& Fink 1997).  Also, the relatively high flux in the far-IR could have a significant contribution from the host galaxy (e.g., Ward et al. 1987; Elvis et al. 1994).  Further evidence against a dust-reddening interpretation comes from consideration of the \\ion{He}{2}~$\\lambda 1640/\\lambda 4686$ ratio. We estimate the observed ratio to be a surprisingly high $19.2\\pm 6.3$ from a STIS measurement of the \\ion{He}{2}~$\\lambda 1640$ line [$(2.5\\pm 0.3)\\times 10^{-13}$ \\ergcms; Collinge et~al. 2001] and a ground-based measurement of the \\ion{He}{2}~$\\lambda 4686$ line [$(1.3\\pm 0.4)\\times 10^{-14}$ \\ergcms; Peterson et~al. 2000]. Such a value is not supportive of a reddening interpretation which predicts the ratio to be $\\la 9$ (compare with Crenshaw et~al. 2002). If the continuum is reddened then the absorption at 1640~\\AA\\ is $\\sim 10$ times larger than at 4686~\\AA\\ (by comparing the observed local-continuum flux ratio in Figure~\\ref{sed} to what is seen in a normal AGN). Such absorption would be expected to produce a \\ion{He}{2}~$\\lambda 1640/\\lambda 4686$ ratio of $\\sim 1$. While our measured line ratio should be treated with caution since variability can add a factor of $\\sim 2$ to the uncertainty, it still seems significantly higher than the expected ratio from reddening. More accurate and simultaneous measurements of this and other line ratios are needed to confirm this result.  An alternative to models with significant reddening is that, perhaps due to its low black hole mass (Shemmer et al. 2003), NGC\\,4051 may have a high-temperature big blue bump that is not apparent even in the FUV (compare our Figure~\\ref{sed} with Figure~3 of Puchnarewicz et~al. 1995). Discrimination between significant reddening and a high-temperature big blue bump is challenging, and we leave this task to a future study. For example, it will be necessary to determine unambiguously if the observed excess soft X-ray emission is due to a continuum soft X-ray excess (e.g., Collinge et~al. 2001) or a relativistically broadened \\ion{O}{8} recombination spectrum (Ogle et~al. 2004)."
    },
    "0402/astro-ph0402404_arXiv.txt": {
        "abstract": "We study the consequences  of the hadron-quark deconfinement phase transition in stellar compact objects when finite size effects between the deconfined quark phase and the hadronic phase are taken into account. We show that above a threshold value of the central pressure  (gravitational mass) a neutron star is metastable to the decay (conversion) to a hybrid neutron star or to a strange star.  The {\\it mean-life time}  of the metastable configuration dramatically depends on the value of the stellar central pressure. We explore the consequences of the metastability of ``massive'' neutron stars and of the existence of stable compact  quark  stars (hybrid neutron stars or strange stars) on the concept of limiting  mass of compact stars. We discuss the implications of our scenario on the interpretation of the stellar mass and radius extracted from the spectra of several X-ray compact sources. Finally, we show that our scenario implies, as a natural consequence a two step-process which is able to explain the inferred ``delayed'' connection between supernova explosions  and GRBs, giving also the correct energy to power GRBs. ",
        "introduction": "One of the most fascinating enigma in modern astrophysics concerns  the true nature of the ultra-dense compact objects called {\\it neutron stars} (NS). The bulk properties and the internal structure of these stars chiefly depends upon the equation of state (EOS) of dense hadronic matter. Different models for the EOS of dense matter  predict  a neutron star maximum mass ($M_{max}$) in the range of   1.4 -- 2.2 $M_\\odot$, and a corresponding central density  in range of 4 -- 8 times the saturation density ($\\rho_0 \\sim 2.8 \\times 10^{14}$g/cm$^3$)  of nuclear matter (e.g.  Shapiro \\& Teukolsky 1983; Haensel 2003). In the case  of a star with $ M \\sim 1.4~M_\\odot$, different EOS models predict a radius  in the range of 7 -- 16 km   (Shapiro \\& Teukolsky 1983; Haensel 2003; Dey et al. 1998).  In a simplistic and conservative picture  the core of a neutron star  is modeled as a uniform fluid of neutron rich nuclear matter in equilibrium with respect to the weak interaction ($\\beta$-stable nuclear matter). However, due to the large value of the stellar central density and to the rapid increase of the nucleon chemical potentials with density, hyperons ($\\Lambda$, $\\Sigma^{-}$, $\\Sigma^{0}$, $\\Sigma^{+}$, $\\Xi^{-}$ and $\\Xi^{0}$ particles) are expected to appear in the inner core of the star. Other {\\it exotic} phases of hadronic matter such as  a Bose-Einstein condensate of negative pion ($\\pi^-$) or negative kaon ($K^-$)  could be present in the inner part of the star.  According to Quantum Chromodynamics (QCD)  a phase transition from hadronic matter to a deconfined quark phase  should  occur at a density of a few times nuclear matter saturation density.  Consequently, the core of the more massive neutron stars is one of the best candidates in the Universe where such deconfined phase of quark matter (QM) could be found. This possibility was realized by several researchers soon after the introduction of quarks as the fundamentals building blocks of hadrons (Ivanenko \\& Kurdgelaidze 1969; Itoh 1970; Iachello et al. 1974; Collins \\& Perry 1975; Baym \\& Chin 1976; Keister \\& Kisslinger 1976). Since $\\beta$-stable hadronic matter posses two conserved ``charges'' ({\\it i.e.,} electric charge and baryon number) the quark-deconfinement phase transition proceeds through a mixed phase over a finite range of pressures and densities according to the Gibbs' criterion for phase equilibrium (Glendenning 1992; M\\\"{u}ller \\& Serot 1995). At the onset of the mixed phase,  quark matter droplets form a Coulomb lattice embedded in a sea of hadrons and in a roughly uniform sea of electrons and muons. As the pressure increases various geometrical shapes (rods, plates) of the less abundant phase immersed in the dominant one are expected. Finally the system turns into uniform quark matter  at the highest pressure of the mixed phase (Heiselberg et al. 1993; Voskresensky et al. 2003). Compact stars  which possess a  ``quark matter core'' either as a mixed phase of deconfined quarks and hadrons or as a pure quark matter phase are called {\\it Hybrid Neutron Stars} or shortly {\\it Hybrid Stars} (HyS) (Glendenning 1996; Drago \\& Lavagno 2001). In the following of this paper, the  more {\\it conventional} neutron stars in which no fraction of quark matter is present, will be referred to as {\\it pure Hadronic Stars} (HS).  Even more intriguing than the existence of a quark core in a neutron star, is the possible existence of a new family of compact stars consisting completely of a deconfined mixture of {\\it up} ({\\it u}),  {\\it down} ({\\it d})  and  {\\it strange } ({\\it s}) quarks  (together with an appropriate number of electrons to guarantee electrical neutrality)  satisfying the Bodmer--Witten hypothesis (Bodmer 1971; Witten 1984; see also Terazawa 1979). Such compact stars have been called {\\it strange quark stars} or shortly {\\it strange stars} (SS) (Alcock et al. 1986; Haensel et al. 1986) and their constituent matter as {\\it strange quark matter} (SQM) (Farhi \\& Jaffe 1984; Madsen 1999). Presently there is no unambiguous proof about the existence of strange stars, however, a  sizable amount of observational data collected by the new generations of X-ray satellites, is providing a growing body of evidence for their  possible existence  (Bombaci 1997; Cheng et al. 1998; Li et al. 1999a; Li et al. 1999b; Xu 2002;  Drake et al. 2002). It is generally believed that the unambiguous identification of a strange star will imply that all pulsars must be strange stars. In the present work we argue that the possible existence of strange stars does not conflict with the existence of  {\\it conventional} neutron stars (pure Hadronic Stars).  Present accurate determinations of compact star masses in radio pulsar binaries (Thorsett \\& Chakrabarty 1999)  permit to rule out only {\\it extremely soft} EOS, {\\it i.e.} those giving $M_{max}$ less than about 1.45~$M_\\odot$. However, in at least two  accreting X-ray binaries it has been found evidence for compact stars with higher  masses. The first of these star is Vela X-1, with a reported mass $1.88 \\pm 0.13 M_\\odot$ (Quaintrell et al.  2003), the second is  Cygnus X-2, with a reported mass of $1.78 \\pm 0.23 M_\\odot$ (Orosz \\& Kuulkers 1999). Unfortunately, mass determinations in X-ray binaries are affected by large uncertainties (van Kerkwijk et al. 1995), therefore the previous quoted ``high mass values'' should always be handled with care. In addition to mass determination, existing observational data on the spin frequency of millisecond pulsars and on the thermal evolution of neutron stars do not put severe constraints  on  the EOS of dense matter. Fortunately, this situation is improving in the last few years. In fact, the extraordinary spectroscopic capabilities of the instruments on board Chandra X-ray and XMM-Newton satellites, are giving the unique possibility to perform accurate measurements of the gravitational red-shift in the spectral lines of a few compact X-ray sources. This provide informations on the mass to radius ratio for compact stars and will help us to understand the true nature of these compact objects.  In this work we study the effects of the hadron-quark deconfinement phase transition in stellar compact objects. We show that when finite size effects at the interface between the quark- and the hadron-phase are taken into account, pure Hadronic Stars, above a threshold value of the central pressure (gravitational mass),   are  metastable to the {\\it decay}  ({\\it conversion})  to  hybrid neutron stars or to  strange stars (depending on the properties of EOS for  quark matter). The {\\it mean-life time} of the metastable stellar configuration is related to the quantum nucleation time to form a drop of quark matter in the stellar center, and  dramatically depends on the value of the stellar central pressure. We explore the consequences of the metastability of ``massive'' pure Hadronic Stars and the existence of stable compact quark  stars  (hybrid neutron stars or strange stars)  on the concept of limiting mass of compact stars. Next,  we discuss the implications of our scenario in the interpretation of the mass-radius  constraints extracted from the spectra of several X-ray compact sources. Finally, we discuss the implications of our scenario for Gamma Ray Bursts. ",
        "conclusions": "In the present work, we have investigated the consequences  of the hadron-quark deconfinement phase transition in stellar compact objects when finite size effects between the deconfined quark phase and the hadronic phase are taken into account. We have found that above a threshold value of the gravitational mass a pure hadronic star is metastable to the decay (conversion) to a hybrid neutron star or to a strange star {\\footnote{~The particular type of quark star ({it i.e.} hybrid star or strange star) formed at the end of the stellar conversion, will depend on the details of the quark matter EOS (see sect. 3). Here we want to stress that our scenario does not require as a necessary condition the fulfilment of the Bodmer-Witten hypothesis on the absolute stability of strange matter and, thus the existence of strange stars. The delayed stellar conversion process described in this paper takes place also in the case a more ``traditional'' hybrid star is formed.}}.   % We have calculated the {\\it mean-life time}  of these metastable stellar configurations, the critical mass for the hadronic star sequence, and have explored how these quantities depend on the details of the EOS for dense matter. We have introduced an extension of the concept of limiting mass of  compact stars, with respect to the classical one given by Oppenheimer \\& Volkov (1939). We have demonstrated that, within the astrophysical scenario proposed in the present work, the existence of compact stars with ``small'' radii (quark stars) does not exclude  the  existence of  compact stars with ``large'' radii (pure hadronic stars),  and {\\it vice versa}.  Finally, we have shown that our scenario implies, as a natural consequence a two step-process which is able to explain the inferred ``delayed'' connection between supernova explosions  and GRBs, giving also the correct energy to power GRBs.  There are various specific  features and predictions of the present model, which we briefly mention in the following. The second explosion ({\\it Quark Deconfinement Nova})  take place in a ``baryon-clean'' enviroment due to the  previous SN explosion. Is is possible to have different time delays between the two events since the {\\it mean-life time} of the metastable hadronic star depends  on the value of the stellar central pressure. Thus the present model   is able to interpret  a time delay of a few years (as observed in GRB990705 (Amati et al. 2000;  Lazzati et al. 2001)), of a few months (as in the case of GRB020813 (Butler et al. 2003)), of a few days (as deduced for  GRB011211 (Reeves et al. 2002)), or the nearly simultaneity of the two events (as in the case of SN2003dh and GRB030329 (Hjorth et al. 2003))."
    },
    "0402/astro-ph0402297_arXiv.txt": {
        "abstract": "{% L'ensemble des observations du fond cosmologique repr\\'esentent une source d'informations remarquable pour la cosmologie moderne. Non seulement elle conforte le mod\\`ele du ``Big Bang'', mais pr\\'ecise notablement le cadre dans lequel les structures de l'univers se sont form\\'ees. Mais ce qui est sans doute le plus fascinant est que le fond cosmologique est la voie privil\\'egi\\'ee d'acc\\`es \\`a la physique des tr\\`es hautes \\'energies qui r\\'egnaient dans les tous premiers instants de l'univers et qui pourraient demeurer \\`a jamais inaccessibles de fa\\c con directe aux exp\\'eriences de laboratoire.  Les futures exp\\'eriences, dont Planck en particulier, sont donc une porte ouverte sur la physique du troisi\\`eme mill\\'enaire.\\\\ }{% Observations of the cosmic microwave background represent a remarkable source of information for modern cosmology.  Besides providing impressive support for the Big Bang model itself, they quantify the overall framework, or background, for the formation of large scale structure.  Most exciting, however, is the potential access these observations give to the first moments of cosmic history and to the physics reigning at such exceptionally high energies, which will remain beyond the reach of the laboratory in any foreseeable future.  Upcoming experiments, such as the Planck mission, thus offer a window onto the Physics of the Third Millennium. } ",
        "introduction": "The program of modern cosmology was born with Lema\\^{\\i}tre's 1927 paper in which he proposed a cosmological model primarily motivated by the desire of accounting for what he believed to be the two astronomical facts of major significance for the description of the Universe: its non--zero matter content and the apparent recession of galaxies that he interpreted as a direct evidence for the expansion of the Universe. A few years later, after the clear evidence for an expanding universe obtained by Hubble, Lema\\^{\\i}tre initiated a program whose basic questions still represent fundamental lines of research in modern Cosmology: the very early history of the Universe, including the nature of the initial singularity and its connection to quantum mechanics, and the question of the history of structure formation.  During the rest of the XXth century, cosmology underwent remarkable progress, by the continuation of this confrontation of some of the most recent, often regarded as the most exotic, theories in physics, with hard  astronomical data.  The determination of the values of the cosmological parameters, with a moderate error, has naturally always being one of the central goals of cosmology, although the strength of efforts in this direction has varied over time.  The concept of inflation (A.~Guth, 1981), introduced more than twenty years ago, revolutionized the field, pointing out how Cosmology contained deep connections between high energy physics and some astronomical observations.  Moreover, Inflation suggests that the actual value of several quantities could have a {\\em physical origin}, rather than being just constants that have to be determined.  In addition, the need for a {\\em physical origin} of the fluctuations that seed structure formation reinforced the link between the question of large--scale structure and the physics of the Big Bang. It was recognised during the last twenty years that the properties of large--scale structure, as revealed by the galaxy distribution, was a potential source of key information for understanding the physics that occurred during the very first instants of the universe. For these reasons, the determination of cosmological parameters has become a scientific program which significance goes far beyond the question of establishing the numerical values of the few parameters describing the universe within the framework of general relativity.  Establishing the precise abundance of light elements, which requires the modelling of the chemical evolution of galaxies and therefore the precise understanding of the physical process occurring in stellar interiors of stars, has been a fundamental test of the Big Bang during its first minutes, and is a good example of connections between modern cosmology, some fundamental physics (nuclear physics in this example) and classical astrophysics. There is now good convergence of data to a rather restricted range of possible values for the baryonic content (Charbonnel 2002) of the Universe: $$\\Omega_b \\sim 0.022 h^{-2} \\pm 10\\%$$ This convergence makes Big Bang nucleosynthesis one of the piller of modern cosmology.  The discovery of the CMB provided the third fundamental piller on which the standard Big Bang is build.  The verification of its remarkable black body spectrum by FIRAS/COBE represents the essential last achievement of the \"classical cosmology\" program, allowing a reliable description of the major points of the history of the universe between the first billionth second and the present epoch. However the discovery by DMR/COBE of the fluctuations of the microwave sky has brought an essential observational fact that requires physical explanation beyond the physics well--established in laboratories. Whether inflation is the correct explanation of the origin of the fluctuations in the observed spectrum of the angular fluctuations in the microwave sky is still a matter of debate, although it is remarkable that this theory proposed more than twenty years ago has passed remarkably well several observational tests.  But the need for new physics is increasingly evident. It has also become clear that Cosmology will provide a test bed for this high energy physics that may well remain unattainable otherwise.  In this respect, the observed properties of the CMB fluctuations appear as a remarkably clean tool for investigating early high energy physics.  In fact, the possibility of constraining some parameters to the percent level with Planck, an extraordinary challenge, naturally leads to the idea of \"high precision cosmology\" in a scientific domain where order of magnitudes were the only realistic perspective few years ago! ",
        "conclusions": "The so-called concordance model provides a remarkable simple cosmological model which reproduces well the WMAP results and which is in agreement with a number of astrophysical observations of cosmological relevance. Despite of this success, it should be realised that the WMAP data by themselve do not require the introduction of a cosmological constant. Actual direct evidence for the existence of a non--zero cosmological constant that dominates the density of the universe are rather limited: the Hubble diagram of distant SNIa and the possible detection of the correlation between deep galaxy surveys and CMB. Therefore, it is essential to confirm the actual non-zero value of the cosmological constant (or one of its generalisation like quintessence) by other data. The next generation of large projects dedicated to cosmology will undoubtedly allow the reliable establishment of a non-zero cosmological constant (if this is actually the case...).  This will allow cosmologists to work within the robust framework of a standard model. The high precision that should be obtained from satellite CMB experiments, typically 1\\% in the Planck experiment (see Bouchet et al., this issue), will open the possibility of determining the cosmological parameters with a precision of the same order. This will be possible by combining different data that will provide accurate, complementary information, including those on the power spectrum of matter fluctuations. In order to take full advantage of the accuracy of CMB data, the precision of data with which they are combined should be similar and therefore systematic uncertainties should be controlled with a similar precision. This is the great challenge for precision cosmology but the rewards will be the establishment of a standard model of cosmology to a high precision and probably unique access to physics at energies much beyond what would be attained directly from laboratory experiments."
    },
    "0402/astro-ph0402068_arXiv.txt": {
        "abstract": "{ Infrared spectroscopy is a powerful tool to probe the inventory of solid state and molecular species in circumstellar ejecta.  Here we analyse the infrared spectrum of the Asymptotic Giant Branch star W Hya, obtained by the Short and Long Wavelength Spectrometers on board of the Infrared Satellite Observatory.  These spectra show evidence for the presence of amorphous silicates, aluminum oxide, and magnesium-iron oxide grains.  We have modelled the spectral energy distribution using laboratory measured optical properties of these compounds and derive a total dust mass loss rate of $3\\times 10^{-10}$ M$_{\\odot}$ yr$^{-1}$.  We find no satisfactory fit to the 13$\\mu$m dust emission feature and the identification of its carrier is still an open issue. We have also modelled the molecular absorption bands due to H$_{2}$O, OH, CO, CO$_{2}$, SiO, and SO$_{2}$ and estimated the excitation temperatures for different bands which range from 300 to 3\\,000K. It is clear that different molecules giving rise to these absorption bands originate from different gas layers. We present and analyse high resolution Fabry-Perot spectra of the three CO$_{2}$ bands in the 15$\\mu$m region.  In these data, the bands are resolved into individual Q-lines in emission, which allows the direct determination of the excitation temperature and column density of the emitting gas.  This reveals the presence of a warm ($\\simeq$ 450K) extended layer of CO$_{2}$, somewhere between the photosphere and the dust formation zone. The gas in this layer is cooler than the 1\\,000K CO$_{2}$ gas responsible for the low-resolution absorption bands at 4.25 and 15$\\mu$m. The rotational and vibrational excitation temperatures derived from the individual Q-branch lines of CO$_{2}$ are different ($\\sim$ 450K and 150K, respectively) so that the CO$_{2}$ level population is not in LTE. ",
        "introduction": "One of the characteristics of Asymptotic Giant Branch (AGB) stars is their mass loss. As a star evolves up the AGB, its mass loss rate increases. Just before it evolves off to become a planetary nebula, it sheds most of its mass in a short time, i.e., superwind. Signatures of mass loss can be studied in the infrared by tracing the dust (e.g., Bedijn \\cite{bedijn}, % Justtanont \\& Tielens \\cite{jt92}), and in the sub-millimetre by using molecular rotational lines such as CO (e.g., Knapp \\& Morris \\cite{kn85}, Groenewegen \\cite{groen}, Skinner et al. \\cite{cjs}).  An example of an early AGB star is W Hya. It has a relatively low mass loss rate as deduced from the weak CO millimeter observations (Wannier \\& Sahai \\cite{wannier}; Bujarrabal et al. \\cite{bujar}; Cernicharo et al. \\cite{cerni}). It is classified as a semiregular variable (SRa) from its light curve. The visual magnitude ranges from 6 to 10 over a period of 382 days (Lebzelter et al. \\cite{leb}). It exhibits SiO (e.g., Schwartz et al. \\cite{schwartz}; Clark et al. \\cite{clark}), H$_{2}$O (e.g., Spencer et al. \\cite{spencer}; Lane et al. \\cite{lane}) and OH main line masers (e.g., Szymczak et al. \\cite{szym}; Chapman et al. \\cite{chapman}). It has a weak silicate dust emission at 10$\\mu$m along with a smaller emission feature at 13$\\mu$m. The carrier of the latter is still being debated. W Hya is one of the first AGB stars for which H$_{2}$O rotational emission was reported after the launch of the Infrared Space Observatory (ISO, Kessler et al. \\cite{kessler}). Neufeld et al. (\\cite{neufeld}) and Barlow et al. (\\cite{barlow}) reported that these lines were detected in Short and Long Wavelength Spectrometers (SWS and LWS), respectively. Both authors presented models accounting for the emission. Recently, Zubko \\& Zuckerman (\\cite{zubko}) also reported a model predicting H$_{2}$O line fluxes observed by ISO. Also, W Hya has been detected in the ortho ground state line by SWAS at 557GHz (Harwit \\& Bergin \\cite{harwit}) In our paper, we mainly discuss the SWS spectra taken during the ISO mission. We also present the SWS Fabry-Perot spectrum of W Hya centered on CO$_{2}$ bands. This band has been reported by Ryde (\\cite{ryde}) and Justtanont et al. (\\cite{jus98}). ",
        "conclusions": "By fitting the SED of W Hya from the near- to far-IR, we derive a dust mass loss rate of 3\\,10$^{-10}$ M$_{\\odot}$ yr$^{-1}$ using a combination of three different dust species of amorphous silicate, Al$_{2}$O$_{3}$ and MgFeO. We also tried adding a fourth component of spinel to account for the 13$\\mu$m dust feature but in order to fit it, we also expect to see another strong peak at 16$\\mu$m which is not present in the observed spectrum. From this, we conclude that spinel cannot be the carrier of the 13$\\mu$m dust emission feature. Hence, the debate on the carrier of the 13$\\mu$m feature is still open. With the estimated value of the our dust mass loss rate, we get a gas mass loss rate range of (3.5-8)\\,10$^{-8}$ M$_{\\odot}$ yr$^{-1}$ from simple consideration of momentum dust driven wind. Our derived mass loss rate agrees well with those previously derived to explain the CO emission, taking into account the differences in the adopted distance and velocity. However, the discrepancy still exists between these mass loss rates and those derived from the H$_{2}$O emission, the latter being consistently higher. The difference in derived mass loss rates from CO and H$_{2}$O may simply be due to the abundances assumed for each species.   We obtained the full scan SWS spectrum of W Hya in which we modelled the absorption bands due to various molecules. The derived excitation temperatures of different bands suggest that they originate from different layers of gas, from 300K to 3\\,000K. However, some molecules exist over a wide range of temperatures and cannot really be stratified within a particular zone, particularly those which show multiple absorption bands. The exception may be OH which is a product of water molecules which are dissociated by shocks in the inner part of the envelope.  We also observed CO$_{2}$ bands in emission using SWS FP and were able to resolve the three strongest Q branches in the mid-IR. The resulting rotation diagram yields an excitation temperature of $\\sim$450K while the CO$_{2}$ absorption at 14.97$\\mu$m requires a gas of 1\\,000K to explain the shift of the peak absorption bluewards. The stretching mode at 4.25$\\mu$m also requires a layer of hot gas to explain its width. We find from simple calculation based on the rotation diagram that the vibrationally excited levels may not be in LTE. Whether the CO$_{2}$ absorption and emission lines can indeed be explained by temperature stratification in an extended dynamical atmosphere needs to be investigated in a more detailed radiation transfer model.   \\begin{acknowledgement}  We are grateful to our referee, Dr. J Alcolea for constructive comments which improves the discussion in our paper. This research has made use of the SIMBAD database, operated at CDS Strasbourg, France.  \\end{acknowledgement}"
    },
    "0402/astro-ph0402033_arXiv.txt": {
        "abstract": "{   The most popular candidate for non baryonic dark matter is the neutralino. More than twenty experiments are dedicated to its direct detection. This review describes the most competitive and promising   experiments with different detection techniques. The most recent results are presented with some prospects for the near future. }]   \\baselineskip=13.07pt ",
        "introduction": "The existence of dark matter in the Universe is now  well established  in the astro-particle community reenforced by the recent astrophysical observations of the satellite experiment WMAP\\cite{wmap} : about 27\\% of the mass-energy of the Universe is composed of matter. Ordinary matter (baryons) contributes to about 4\\% of this total mass density  of the Universe and only $\\simeq$ 1 \\% is  visible according to the most recent measurements of the amount of deuterium in high red-shift clouds of gas and of the CMB\\cite{gas}. Hence about  90\\% of this dark matter is not baryonic.  We have to distinguish two categories, hot and cold dark matter particles refering to their velocity at the matter-radiation decoupling time in the early Universe.  Hot dark matter implies moving relativistically and cold  moving non-relativistically. Neutrinos with  non-zero masses are hot dark matter candidates, however WMAP\\cite{wmap} results  combined with other experiments and observations lead to a contribution $<1.5 \\%$ for light neutrino species.  So the bulk of the non-baryonic dark matter is cold dark matter (CDM). Among the numerous solutions proposed by theorists axions and neutralinos are favorites. Neutralinos are candidates of the generic class of Weakly Interacting Massive Particles (WIMP). Axion  are particles  proposed to solve the strong CP violation problem in the Peccei-Quinn theory \\cite{peccei}. Astrophysical considerations combined with experimental constraints require an axion mass in the range $10^{-3}$ to $10^{-6} eV/c^2$ . For a more detailed discussion about axions see reference\\cite{axion1,axion2,axion3}, since this paper will be dedicated to the WIMP/neutralino detection. The neutralino is the lightest supersymmetric particle, a linear combination of the supersymmetric partners of the photon, Z and Higgs bosons, in the minimal~supersymmetric~extension~of~the~standard~model : \\begin{equation} \\chi^0 = a \\tilde \\gamma + b \\tilde Z + c \\tilde{H_1^0} + d \\tilde{H_2^0} \\label{eq:neutralino} \\end{equation}  Its mass is constrained to lie in the range 45~GeV~$< m_{\\chi} <$~3~TeV, where the lower bound comes from accelerator results from LEP and the upper bound is given by  astrophysical constraints such as the age of the Universe or unitarity. Locally our galaxy is supposed to be imbedded in a WIMP halo.  Many experiments are  dedicated to direct and indirect detection of WIMPs, two complementary techniques. Direct detection experiments measure the energy deposited by elastic scattering of a neutralino of our own galaxy off a target nucleus.  For  masses larger than $\\simeq$ 200 GeV, indirect detection of dark matter particles through their annihilation products may be more suitable. In this paper we will concentrate on the case of direct detection techniques, and for a complete description of  indirect approaches we refer to the papers\\cite{indirect1}$^-$\\cite{indirect6}.  In the direct detection approach the expected event rate depends on various parameters coming from astrophysics, particle physics and nuclear physics; it can range from 1 to $10^{-5}$ events/kg/day. The measured signal is very low (few keV) depending  on the masses of the incident particle and of the scattered nucleus, but also on the nuclear recoil relative efficiency (quenching factor) in producing charges, light or heat. Hence  WIMP direct searches put strong constraints on  experimental background environments, and require detectors with very low energy thresholds. In this review we present the different possible signatures for disentangling  a WIMP signal from the background. Different experimental approaches are described and illustrated  by a few experiments. The current limits in the exclusion plot and near future prospects will be also presented.    \\section {WIMP/neutralino direct detection physics principles}  \\begin{figure} \\center \\psfig{figure=figure1.eps,width=10.0truecm} \\caption{Annual modulation.} \\label{fig:modann} \\end{figure}  As mentionned previously, WIMP-nucleus interaction rate depends on various parameters. First we have to define a WIMP halo model. For simplicity the approximation of a maxwellian velocity distribution in the galactic frame is made (see \\cite{kamionkowski} for a review on alternative halo models). Next, a supersymmetric model is chosen for predicting the WIMP interaction with quarks of nucleons inside the target nucleus. Depending on the chosen model  the WIMP-nucleus cross-section has two components\\cite{smith,jungman} :  spin-dependent and  spin-independent. The spin-independent term couples to the mass of the nucleus and the spin-dependent couples to its spin. The nuclear form factor depends on  the nature of the interaction. The spin-dependent case is the most complicated one, requiring detailed nuclear models (for more details see dedicated papers\\cite{jungman,ressel}). In the following we will restrict this review to the simplest spin-independent case which is supposed to dominate in most models for massive target nuclei. Taking into account these previous considerations the interaction rate can be expressed as follows :  \\begin{equation} {{dR}\\over {dQ}} = {{\\sigma_0 \\rho_h}\\over {2 m^2_r m_\\chi}} F^2(Q) {\\int_{v_{min}}^\\infty} {f(v)\\over v} dv \\label{eq:rate} \\end{equation}  where $m_r$ is the WIMP-nucleus reduced mass $m_\\chi m_N/(m_\\chi+m_N)$, $m_\\chi$ is the WIMP mass, $m_N$ is the nucleus mass.  $\\rho_h = 0.3 GeV/c^2/cm^3$ is the assumed halo WIMP density at the position of the solar system, $f(v)$ is the dark matter velocity distribution,  with an average rms velocity $v_0 = 220 km/s$,  truncated above the escape velocity of the galaxie $v_{esc}\\simeq 575 km/s$, $\\sigma_0$ is the total nucleus-WIMP interaction cross section and $F(Q)$ is the nuclear form factor.  \\subsection{Exclusion plot $\\sigma (m_{\\chi})$}  In order to reliably compare supersymmetric models with results obtained by different experiments using different techniques a $\\sigma_0 (m_{\\chi})$ plot is built in the following way. The cross-sections $\\sigma_0 (m_{\\chi})$ are normalized to a single nucleon $\\sigma (m_{\\chi})$ to allow comparisons between  different target nuclei. The measured nuclear recoil event rate is compared to a theoretical spectrum calculated for a given  WIMP mass and cross-section. If an experiment observes a signal then we build a $\\sigma (m_{\\chi})$ contour plot. If the observed events cannot be unambigously associated with a WIMP signal an exclusion limit is calculated. WIMP signals have distinctive signatures that backgrounds are not supposed  able to mimic. Three different signatures are proposed.  \\begin{figure} \\center \\psfig{figure=figure2.eps,width=8.0truecm} \\caption{Diurnal modulation.} \\label{fig:moddiu} \\end{figure}  \\subsection{Annual modulation }\\label{subsec:modann}  As a result of the   Earth motion around the Sun the count rate in detectors should show an annual modulation (fig.~\\ref {fig:modann}). Along the year the Earth's velocity relative to the galaxy varies, as in June the Earth and the Sun velocities add up  whereas in December they subtract\\cite{hasen98}. The maximum amplitude of this effect in the signal is about 7\\%.  We will report  later that the DAMA collaboration using NaI scintillating crystals is the first experiment  and at the moment the only one, claiming for an evidence of a WIMP annual modulation signal.   \\subsection{Diurnal modulation and directionnality}\\label{subsec:moddiu}  Another possible modulation in the WIMP signal is the night and day variation, this effect is due to the shielding of the detector by the Earth of the incident flux. For masses close to 50 GeV and under certain assumptions the diurnal modulation can be larger than the annual one\\cite{spergel88}. However the most interesting daily signature coupled with the annual one  is the directionnality of the WIMP wind as illustrated in figure~\\ref{fig:moddiu}. This  effect is also larger than the annual one. The validation of the principle has been performed by the DRIFT-I experiment with a 1 m$^3$ low pressure TPC\\cite{martoff00} prototype.   \\subsection{Target atomic mass effect}  Observed together annual and diurnal modulations are unambigous methods to distinguish WIMP and background signals, but they are very difficult to operate. In the spin-independent case, an easiest method is to use different target materials as the event rate depends on the target atomic mass. To give an estimate of this effect we can use the Smith and Lewin\\cite{smith} calculated integral rate $R_0$ with no form-factor correction and an average recoil energy $E_{R}$  \\begin{equation} R_0 \\simeq 5.87\\times  v_0 \\sigma_{m_\\chi } \\rho_h A^3 {{m_\\chi }\\over{(m_A+m_\\chi)^2}} /kg/d \\label{eq:rzero} \\end{equation}  \\begin{equation} E_R \\simeq 2\\times 10^{-6} m_\\chi^2 ({v_0 \\over c})^2  {{m_A }\\over{(m_A+m_\\chi)^2}} keV \\label{eq:er} \\end{equation}   Table ~\\ref{tab:targets} reports on  $R_0$ and $E_{R}$ values for different targets, and for a given WIMP mass of $m_\\chi \\simeq 50$ GeV, $\\sigma_{m_\\chi } \\simeq 7.10^{-6}pb$. Naively if we consider the event rate it seems to be more advantageous to use high mass nuclei, but if we look to the recoil energy as the target atomic number A increases, the average deposited energy tends to decrease. So the choice of a target is a compromise between these two quantities. Moreover we can see for example that germanium is more efficient than silicon for WIMPs detection while they have similar cross-sections for neutrons.\\par Another important point is the possible neutron multiple scattering in the detector, which is  impossible for a WIMP. We will see hereafter this  method is used by the CDMS collaboration\\cite{akerib1}, with germanium and silicon targets as illustrated in figure~\\ref{fig:cdms}.   \\begin{table} \\begin{center} \\caption{Integrated event rate $R_0$ and average energy deposition for different target atomic masses, no form-factor correction and $m_\\chi = 50$ GeV, $v_0 = 220 km/s$ and $\\sigma _{n\\chi} = 7\\times 10^{-6}$ pb.\\label{tab:targets}} \\vspace{0.2cm} \\begin{tabular}{|c|c|c|c|} \\hline \\raisebox{0pt}[12pt][6pt]{} & \\raisebox{0pt}[12pt][6pt]{$A$} & \\raisebox{0pt}[12pt][6pt]{$R_0$} & \\raisebox{0pt}[12pt][6pt]{$<E_R>$} \\\\ \\hline \\raisebox{0pt}[12pt][6pt]{H} & \\raisebox{0pt}[12pt][6pt]{1} & \\raisebox{0pt}[12pt][6pt]{5.10$^{-5}$} & \\raisebox{0pt}[12pt][6pt]{1} \\\\ \\hline \\raisebox{0pt}[12pt][6pt]{Na} & \\raisebox{0pt}[12pt][6pt]{23} & \\raisebox{0pt}[12pt][6pt]{0.3} & \\raisebox{0pt}[12pt][6pt]{11} \\\\ \\hline \\raisebox{0pt}[12pt][6pt]{Si} & \\raisebox{0pt}[12pt][6pt]{28} & \\raisebox{0pt}[12pt][6pt]{0.5} & \\raisebox{0pt}[12pt][6pt]{12} \\\\ \\hline \\raisebox{0pt}[12pt][6pt]{Ge} & \\raisebox{0pt}[12pt][6pt]{73} & \\raisebox{0pt}[12pt][6pt]{3} & \\raisebox{0pt}[12pt][6pt]{13} \\\\ \\hline \\raisebox{0pt}[12pt][6pt]{I} & \\raisebox{0pt}[12pt][6pt]{127} & \\raisebox{0pt}[12pt][6pt]{8} & \\raisebox{0pt}[12pt][6pt]{11} \\\\ \\hline \\raisebox{0pt}[12pt][6pt]{Xe} & \\raisebox{0pt}[12pt][6pt]{131} & \\raisebox{0pt}[12pt][6pt]{9} & \\raisebox{0pt}[12pt][6pt]{11} \\\\ \\hline \\raisebox{0pt}[12pt][6pt]{Pb} & \\raisebox{0pt}[12pt][6pt]{210} & \\raisebox{0pt}[12pt][6pt]{18} & \\raisebox{0pt}[12pt][6pt]{8} \\\\ \\hline \\end{tabular} \\end{center} \\end{table}  \\section {WIMP/neutralino direct detection techniques}  WIMP  detectors  are constrained by  three important requirements : low threshold, ultra low background and high mass detector. When a WIMP interacts with a nucleus, the nuclear recoil can induce different signals (fig. ~\\ref{fig:dephys}) : heat, ionization and scintillation. During the last decade important technical developments were based on one or two of  these different physics processes.  \\begin{figure*} \\psfig{figure=figure3.eps,width=6truein} \\caption{Illustration of the different techniques developped for the WIMP direct detection.} \\label{fig:dephys} \\end{figure*}  \\subsection{Quenching factor}  A relevant parameter in WIMP direct detection is the relative efficiency of nuclear recoil   called quenching factor. It is the ratio of the number of charge carriers produced by a nuclear recoil due to the WIMP interaction over an electron recoil of the same kinetic energy (electron equivalent energy or \"eee\"). For scintillating materials  the quenching factor is defined as the ratio between the light produced by a nuclear  recoil and by an electron recoil. \\par \\noindent While in conventional detectors this factor is usually below 30\\% (measured, e.g. to be $\\simeq$ 0.3 for germanium\\cite{messous}, $\\simeq$ 0.25 for sodium and $\\simeq$0.08 for iodine\\cite{gerbier} ), for cryogenic detectors described hereafter it has been measured to be around one for  recoiling nuclei independently on energy\\cite{zhou,alessandrello,sicane}.  \\subsection{Classical detectors : semiconductors and scintillators}   Germanium diodes initially used in double beta decay experiments were the first detectors  used to search  for WIMPs, since they have very low thresholds and very good resolutions. Experiments like IGEX\\cite{igex1,igex2} and HDMS\\cite{hdms}, with about 2 kg of enriched $^{76}Ge$,  achieved very low background count rates ($<$0.2 evt/kg/day in the interval 10-40 keV) and $E_{thr} \\simeq 4-10$ keV-ee (equivalent to $\\simeq 15-30$ keV recoil).   \\begin{figure} \\center \\psfig{figure=figure4.eps,width=7.0truecm} \\caption{DAMA model independent residual count rates as a function of time for 7 years and three energy intervals (2-4), (2-5) and (2-6) keV-ee. \\label{fig:damamod}} \\end{figure}   Large masses were easily achievable with scintillators like NaI or liquid xenon in a very pure environment. The DAMA experiment has operated more than 100 kg of NaI (each crystal weighting about 9.7 kg with energy threshold of $\\simeq 2 keV-ee$ ie 22 keV recoil) for several years in the Gran Sasso underground laboratory. They accumulated data during 7 years and since 1997 they announce evidence for an annual modulated WIMP signal. The DAMA group claim their observation  is compatible with a signal induced by a WIMP of $\\simeq 52$ GeV mass ad a WIMP-nucleon cross section of $\\simeq 7.2 $ pb. The DAMA collaboration has published\\cite{dama03} this last summer the last 3 years campaign totalizing 7 years and confirms their observation of an annual modulation signal as illustrated in figs.~\\ref{fig:damamod},\\ref{fig:damalimit}. Right now none of the currently running dark matter  experiments confirms this signal as we can see in the current exclusion plot in fig.~\\ref{fig:limit03}. Independent experiments with NaI detectors (NAIAD\\cite{naiad} in the Boulby mine, ANAIS\\cite{anais} in Canfranc, ELEGANT\\cite{elegant} in Oto Cosmo Observatory) are currently running. The NAIAD\\cite{naiad1} experiment most recent results begin to exclude   the DAMA $\\sigma(m_\\chi)$ region in the spin-independent exclusion plot.    As we have seen previously despite the very high purity level of classical detectors,  they suffer ultimately from a lack of power discrimination between electron and nuclear recoils.  The first discrimination method used   is based on a pulse shape analysis. It is a statistical method where the measured quantity is the rise-time of the light signal which depends on the nature of the recoiling particle. This discrimination method is used with sodium iodide crystals (DAMA, NAIAD) but is also successfully used with liquid scintillators like liquid xenon. \\par With a 3.1 kg liquid-Xenon detector the ZEPLIN-I\\cite{sumner} collaboration has reached preliminary sensitivities which could exclude the DAMA zone. However some problems remain : a relativily high electronic background rate has to be understood, there's no nuclear recoil calibration for the low energy part of the spectrum ($<$50 keV-ee), a poor energy resolution compared to bolometers. Some of these points should be answered in the next few months as the experiment in now  currently running deep underground in the BOULBY mine\\cite{smith}. \\par The DAMA/LIBRA collaboration is currently running a new NaI detector mith a larger mass ($\\approx$ 250~kg) as well as  a liquid-Xenon detector.  The future projects ZEPLIN-II  and -III aimed to be able to develop a discrimination technique with a two phase  liquid-gas Xenon detector with charge and light signals.  \\begin{figure} \\center \\psfig{figure=figure5.eps,width=8.0truecm} \\caption{Dama limit for the dominant spin-independent case obtain with 7 years of data taking. This contour plot is obtain with different WIMP-halo models, see ref$^{29}$ for a detailed discussion.} \\label{fig:damalimit} \\end{figure}  \\begin{figure} \\center \\psfig{figure=figure6.eps,width=5.0truecm} \\caption{CDMS detector tower.} \\label{fig:cdms} \\end{figure}   \\subsection{Cryogenic detectors}  Since the beginning of the 90's important developments were also made in new directions like cryogenic detectors. They are made of a crystal with a thermometer glued on it, operating  at very low temperature (few tens of millikelvin). Very low thresholds were reached  by the CRESST-I experiment\\cite{cresst1} with a 262 g sapphire calorimeter (resolutions of $\\simeq 133$ eV at 1.5 keV and thresholds $\\simeq 500$ eV).   \\par But most impressive results were obtained with mixed techniques allowing the simultaneous measurement of two components heat-light or heat-charge.  The two combined informations are a powerful tool to distinguish a nuclear recoil induced by a WIMP or a neutron interaction from electron recoils induced by a gamma or an electron interaction (quenching factor described previously). It is an event by event discrimination method. Again different approaches were explored by different worldwide collaborations. For cryogenic detectors the CDMS and EDELWEISS collaborations investigate the heat-ionization way, and  the CRESST and ROSEBUD collaborations explore the heat-light channels.   The CDMS collaboration was the first\\cite{cdms0,cdms01} to operate a detector giving simultaneously ionization and heat signals with a germanium crystal. Until 2002 the experiment was running in the shallow  site in Stanford with a poor muon shielding  inducing an important neutron background. Despite this limitation they derive competitive dark matter limits and were leaders for several years. They could  subtract the neutron backgroung using a monte carlo simulation but also taking advantage of the fact that they run simultaneously two different targets : germanium and silicon\\cite{akerib1,akerib2}. During the year 2003 the CDMS-II experiment is being installed in  the deep underground Soundan mine  where the muon flux is reduced by 5 orders of magnitude reducing the neutron background by a factor 400. They are currently operating 2 towers (fig.~\\ref{fig:cdms}) of 3x165 g Ge and 3x100 g Si detectors  and 18 more detectors are under fabrication totalizing 4 kg of germanium. The CDMS collaboration expects to improve its current sensitivity ($\\simeq 1 evt/kg/day$) by two orders of magnitude.    \\begin{figure} \\center \\psfig{figure=figure7.eps,width=7.0truecm} \\caption{EDELWEISS 320 g Ge detector.} \\label{fig:edw} \\end{figure}    \\begin{figure} \\center \\psfig{figure=figure8.eps,width=7.0truecm} \\caption{Discrimination between gammas and nuclear recoils in a 50 g sapphire bolometer at 20 mK by the  ROSEBUD collaboration .} \\label{fig:rosebud} \\end{figure}  \\begin{figure} \\center \\psfig{figure=figure9.eps,width=8.0truecm} \\caption{Current spin-independent limits for the most competitive experiments. The WIMP halo parameters used are $\\rho_h=0.3 GeV/c^{-2}cm^{-3}$, $v_0=220 km/s$. The closed contour corresponds to the 3$\\sigma$ allowed region of the DAMA first four years obtain with the same WIMP halo parameters.} \\label{fig:limit03} \\end{figure}   The currently best spin-independent published limit was obtained by the EDELWEISS collaboration cumulating 32 kg.d. The EDELWEISS experiment is installed in the underground laboratory of Modane in the French-Italian Alps. They operate similar detectors to those of CDMS germanium crystals (fig.~\\ref{fig:edw})  with different technologies for the electrodes\\cite{edw1} runing at $\\simeq 18 mK$. Three 320 g detectors are running simultaneously. During the last campaign in june 2003, 2 events were observed in the nuclear recoil zone which origin is under investigation. More data is being analysed, but the EDELWEISS-I  stage data taking will be soon  finished. For the next stage a larger cryostat with a detection volume of 100 litres is built  and is currently beeing tested. This cryostat benefits from an original technology developped at the CRTBT-Grenoble laboratory. The EDELWEISS-II installation will take a year from now. The first step will operate 21x320 g germanium detectors with NTD thermometers and 7x200 g NbSi thin film germanium detectors developped by the group of the CSNSM laboratory\\cite{csnsm}. A muon veto made of 140 $m^2$ plastic scintillator will be  added. It should reject the  neutron background induced by cosmic muons in the inner lead shielding, which has been  evaluated  two orders of magnitude below the present EDELWEISS-I sensitivity $\\simeq 0.2$ evt/kg/day.  Such  background has to be clearly identified and rejected since the expected  event rate for the EDELWEISS-II stage is about $10^{-2}$ evt/kg/day. In a second step up to 120 detectors will operate simultaneously   The CRESST-II\\cite{cresst2} and ROSEBUD\\cite{demarcillac} experiments involve scintillating crystals as cryogenic detectors. They operate in the same way ; the heat is measured with thermometer glued on the scintillator and the light is collected with a second thin but large surface crystal. The main advantage of such method is the large possibility for scintillating target materials : CaWO$_4$, PbWO$_4$, Al$_2$O$_3$, BaF, BGO, ... and for important volumes. A few years ago S.P\\'ecourt et al.\\cite{sophie} characterized the phonon channel of a 1 kg Al$_2$O$_3$ bolometer and recently the same team\\cite{demarcillac} has succeeded in measuring the light output of a 50 g Al$_2$O$_3$ bolometer(fig.~\\ref{fig:rosebud}). \\par The CRESST-II\\cite{cresst2} experiment should operate 33x300 g modules of CaWO$_4$ totalizing about 10 kg.     \\subsection{New promising techniques}   \\begin{figure} \\center \\psfig{figure=figure10.eps,width=6.0truecm} \\caption{PICASSO new 1liter module } \\label{fig:picasso} \\end{figure}    \\begin{figure} \\center \\psfig{figure=figure11.eps,width=8.0truecm} \\caption{DRIFT-1 ionization tracks for three different types of recoiling particles : argon, helium and electrons} \\label{fig:drift} \\end{figure}   In addition to the techniques described above, illustrated by currently running experiments and their near future, other promising techniques are under investigation.  The PICASSO\\cite{picasso1,picasso2} and SIMPLE\\cite{simple} experiments have choosen to adapt a well known technology used in neutron dosimetry, to develop a counter for WIMP induced nuclear recoils. The method is based on small superheated Freon droplets imbedded in a gel matrix at room temperature. The nuclear recoil of $^{19}F$ induces the explosion of a droplet , creating an acoustic shock wave measured with  piezoelectric transducers.   By varying the temperature of the gel the energy threshold can be triggered in such a way that the electron recoil induced by gamma background can be supressed. Calibration is made at different pressures and temperatures with monoenergetic neutrons produced by a Van de Graff Tandem . The use of $^{19}F$ (spin-$1/2$ isotope) is particularly interesting to search for spin-dependent neutralinos. A first generation of detectors, 16 modules of  8 ml, lead to the published limit of the PICASSO collaboration\\cite{picasso1,picasso2}. They are currently running the second generation of modules with a larger volume (fig.~\\ref{fig:picasso})) in an improved low background environnement in the SNO underground laboratory : PIC@SNO. New purification techniques were developped especially for the PICASSO experiment \\cite{picasso3}. Despite a very good backgroung discrimination the main disadvantage of such an integrating  detector is the necessity to run the experiment at different threshold energies in order to measure the deposited energy spectrum. \\par  To take advantage of the directionnality which appears as the clearest signature of WIMPs, the UKDMC collaboraton has developped and is currently running successfully, the DRIFT-I detector. It consists in  a 1 m$^3$ low pressure TPC filled with a $Xe-CS_2$ gas mixture. The principle of the TPC is well known, the innovation is the use of $CS_2$ negative ions instead of $e^-$ as charge carriers reducing the diffusion in order to achieve millimetric track resolution (fig.~\\ref{fig:drift}). Important  improvements on the read-out techniques such as MICROMEGAS\\cite{luscher}, in order to increase the pressure hence the target mass, are underway. Other possible target gases are also studied to prepare the next generations DRIFT-II and -III with a larger gas mass  for the TPC of the order of 100 kg.  \\begin{figure} \\psfig{figure=figure12.eps,width=3.5truein} \\caption{Projected limits for some of the next generation experiments. The colored regions represent different SuSy model calculations \\label{fig:prospect} } \\end{figure} ",
        "conclusions": "The current experimental spin-independent limit turns  around 10$^{-6}$~pb which corresponds to a count rate of about 0.2 to 1~evt/kg/day. To achieve this limit it took about 10 years for most of the currently running first generation experiments to develop these detectors. The next generation under construction and for most of them on the final stage, aim to improve this limit by two orders of magnitude, that means a count rate around 10$^{-2}$~evt/kg/day. This has a price : lowering the sensitivity by about two orders of magnitude  implies increasing the target mass by about the same factor (for example EDELWEISS-I worked with 3x320 g Ge and EDELWEISS-II should run at the end 120x320 g Ge detectors).\\par With this scaling the ultimate  neutron background induced by muons can no longer be neglected. It is the reason why experiments like EDELWEISS-II,CDMS-II and CRESST-II will use a muon veto. \\par The next five years are very promising : a clarification of the DAMA annual modulation signal is essential. Indirect Earth-based and Space experiments like Antares, HESS, AMS and GLAST should give  independent cross checks. Meanwhile accelerator physics will explore an important part of SuSy space parameters on the exclusion plot (fig.~\\ref{fig:prospect}). \\par Nevertheless the one-tonne scale experiment will probably involve larger  international collaborations. The technical challenge will be to build an experiment able to achieve the extremely low background necessary to cover most of the prediction mSUGRA models."
    },
    "0402/astro-ph0402519_arXiv.txt": {
        "abstract": "We present near-infrared spectra of ``A2'', the primary counter arc to the gravitationally lensed galaxy MS1512-cB58. The spectra show redshifted \\ha, [NII], [OIII], and \\hb\\ at $z=2.728\\pm 0.001$.  We observe the same \\ha/[OIII] ratio as cB58, which together with the redshift confirms that A2 is indeed another image of a single background galaxy. Published lensing reconstruction reports that A2 is a magnification of the entire source, while cB58 is an image of only a part. At marginal significance, A2 shows higher line to continuum ratios than cB58 (by a factor of $\\sim 2$), suggesting a non-uniform ratio of young to old stars across the galaxy. We observe a second emission line source in the slit.  This object, ``W5'', is predicted to be a lensed image of another galaxy at a redshift similar to cB58.  W5 is blueshifted from cB58 by $\\sim 400$\\ \\kms\\ and has a significantly lower \\ha/[OIII] ratio, confirming that it is an image of a different background galaxy in a group with cB58.  The \\ha\\ emission line in W5 implies a star formation rate of 6 $\\mbox{M}_{\\odot}\\ \\mbox{yr}^{-1}$\\ ($H_0 = 70$ km s$^{-1}$\\ Mpc$^{-1}$, $\\Omega_M=0.3, \\Omega_{\\Lambda}=0.7$), after correcting for lensing magnification. ",
        "introduction": "Most of the detailed spectral information on $z>2$~ galaxies has come from observation of the rest-frame ultraviolet redshifted into the optical passband.  The Lyman Break Galaxies (LBGs; Steidel et al. 1996, 2003) are strongly starbursting galaxies, and in principle may be the tracers of the global star formation history of the universe (Madau et al. 1998) if the effects of dust extinction on the UV continuum can be quantified.  The average attenuation from dust has been suggested to be between a factor of $<2$\\ and 10 (Pettini et al.\\ 1998; Trager et al. 1997).  Even with the advent of IR spectrographs on large telescopes, only a few dozen $z>2$\\ starbursts have been spectroscopically observed in the rest frame optical (e.g. Teplitz et al. 2000a, 2000b; Kobulnicky \\& Koo 2000; Pettini et al.\\ 1998, 2001; Lemoine-Busserolle et al.\\ 2003; Erb et al.\\ 2003).  MS1512-cB58 (hereafter cB58) is a strongly lensed, but otherwise typical, star-forming galaxy at z=2.729 (Yee et al.\\ 1996; Seitz et al.\\ 1998, hereafter S98).  Its magnification by a factor of $\\sim 30-50$\\ makes it the apparently brightest LBG known.  Spectra have been obtained in the rest-frame UV (Pettini et al.\\ 2000) and optical (Teplitz et al.\\ 2000b, hereafter T00).  Lensing reconstruction identifies another faint source in the same field as a counter arc to cB58 (S98). This object, designated ``A2'' by S98, is a magnification of a different part of the same background source.  Thus the gravitational lensing provides a means to spatially resolve the observations of this intrinsically faint source.  Reconstruction shows that cB58 is a magnification of only half of the source galaxy, including the nucleus, and that A2 is a magnification of  the entire source.  The lensing model also identifies several other faint objects as probable arclets that are magnified by the foreground cluster (see S98 and their Figure 1).  At least four background objects are likely to be lensed: the A source, magnified as cB58 and A2; the B source, possibly at $z\\sim 3$; the C source, most likely at a redshift slightly larger than cB58; and the W source, which should have a redshift similar to cB58.  The W system of consists of five arclets:  three sources with ``shrimp-like'' morphologies in the Hubble Space Telescope (HST) imaging; a fourth, very faint source; and a probable 5th image with an uncertain predicted location.  In this paper, we present near-infrared (NIR) spectra and \\kp\\ imaging of A2 and the probable fifth image of the W source, which we call W5. The data were obtained with the NIRSPEC instrument on the Keck II telescope.  We also compare our new observations with archival HST imaging of the objects.  Throughout the paper, we will adopt a flat, $\\Lambda$-dominated universe ($H_0 = 70$ km s$^{-1}$\\ Mpc$^{-1}$, $\\Omega_M=0.3, \\Omega_{\\Lambda}=0.7$). ",
        "conclusions": "The large magnification of cB58 makes it an important target for available instrumentation at many wavelengths.  For example, it is a planned target for Spitzer guaranteed time observations, which will probe the star-formation and dust content of such galaxies.  The addition of the counter arc, A2, and the companion galaxy, W, serve to increase the return on future observations of the field.  We expect that confirmation of the C and possibly B sources will soon show them to be members the same group.  We have shown that A2 is the counter image of MS1512-cB58.  The spatial resolution provided by differential magnification suggests an uneven distribution of star-forming regions in the source galaxy.  The rotational velocity limit obtained from the A2 spectrum supports other observations of LBGs at somewhat lower redshift.  We have also confirmed the existence of the companion W source, which is intrinsically fainter than the A source.  The probable lower metallicity in the W source might be evidence of the luminosity-abundance relationship at high redshift (Kobulnicky et al.\\ 2003).  These lensed LBGs are valuable examples of moderate luminosity galaxies.  The magnification enables detailed studies of average luminosity galaxies at high redshift.  A larger sample is still needed, however, to draw more general conclusions.  Other galaxy clusters offer magnified star-forming galaxies at high redshift as well (i.e.\\ Fosbury et al.\\ 2003), but so far none of comparable brightness have been discovered.  Future study of lensed LBGs will extend to higher redshifts (i.e.  Ellis et al.\\ 2001).  In fact, the observation of gravitationally lensed sources may reveal the highest redshift and intrinsically faintest sources (Pell\\'{o}\\ et al.\\ 2003)."
    },
    "0402/astro-ph0402205_arXiv.txt": {
        "abstract": "The quasar Q0353-383 has long been known to have extremely strong nitrogen intercombination lines at $\\lambda$1486 and $\\lambda$1750 \\AA, implying an anomalously high nitrogen abundance of $\\sim 15$ times solar.  A search for similar nitrogen-rich quasars in the Sloan Digital Sky Survey First Data Release (SDSS DR1) catalog has yielded 20 candidates, including four with nitrogen emission as strong or stronger than that seen in Q0353-383.  Our results indicate that only about 1 in 1700 of quasars have nitrogen abundances similar to Q0353-383, while up to 1 in 130 may be in the process of extreme nitrogen enrichment. ",
        "introduction": "The quasar Q0353-383 \\citep{OsmerSmith1980} is an unusual object, with prominent \\ion{N}{3}], \\ion{N}{4}] and \\ion{N}{5} emission lines and abnormally weak \\ion{C}{3}] and \\ion{C}{4} lines compared to other quasars. To illustrate this point, Figure 1a displays the spectrum of Q0353-383 (Baldwin, 1992, private communication) in comparison to a ``normal'' quasar spectrum given by the Sloan Digital Sky Survey \\citep{York2000} composite in Figure 1b \\citep{VandenBerk2001}. \\citet{Osmer1980} concluded that Q0353-383 has an anomalously high nitrogen abundance due to recent CNO processing in stars. \\citet{Baldwin2003} used improved data and models to confirm and refine these conclusions, and to determine that Q0353-383 has a metallicity of at least 5 times solar, and more likely 15 times solar. Simulations by \\citet{Hamann1999} show that this level of overabundance is expected to occur near the end of an era of rapid metal enrichment which can result in metallicities of as much as 10-20 times solar.  The scarcity of objects like Q0353-383 may be an indication of the amount of time a quasar spends in this state of extreme metal enrichment before the gas supply is exhausted and the quasar becomes inactive.  As the only object of its kind known, Q0353-383 raises some obvious questions: what percentage of the quasar population is nitrogen strong, and what are the global properties of the nitrogen-strong quasar population?  Until recently, anwering these questions was difficult due to the relatively nonstandard collection of quasar data in various wavelength regimes and the lack of spectra with high S/N (signal-to-noise ratio).  The advent of the Sloan Digital Sky Survey (SDSS) has remedied this situation by working to compile, in one database, approximately 100,000 high-quality quasar spectra as it scans 10,000 deg$^2$ of the north Galactic cap \\citep{York2000}. \\citet{Bentz2003} searched the SDSS Early Data Release (EDR, \\citealt{Stoughton2002}) for objects similar to Q0353-383 and determined that although several objects have nitrogen emission that is unusually strong, none of the objects in the EDR Quasar Catalog \\citep{Schneider2002} with $1.8 < z < 4.1$ have emission from both \\ion{N}{4}] $\\lambda$1486 and \\ion{N}{3}] $\\lambda$1750 with strengths that are comparable to Q0353-383.  In this paper, we have searched for nitrogen-rich quasars in the Quasar Catalog \\citep{Schneider2003} from the First Data Release (DR1, \\citealt{Abazajian2003}), which covers almost three times the area on the sky as the EDR, and has over four times as many quasars.  We present numerous objects with stronger nitrogen emission than is usually seen in quasars, including four objects that have emission as strong or stronger than that seen in Q0353-383. ",
        "conclusions": "We have searched 6650 quasars in the SDSS DR1 database for nitrogen-rich objects similar to Q0353-383.  Four candidates have nitrogen emission as strong or stronger than that seen in Q0353-383, and an additional 16 exhibit slightly weaker nitrogen emission. We have also identified 33 objects that may have visible \\ion{N}{4}] and \\ion{N}{3}] emission, although it is less clear from the quality of their spectra.  With the data available, we may set a lower limit of 0.06\\% (about 1 in 1700) on the number of nitrogen-rich objects that are truly similar to Q0353-383.  If we view Q0353-383 and its companion SDSS quasars as being the most extreme objects in a continuous phase of nitrogen enrichment, then it also appears that a lower limit of 0.2\\% - 0.7\\% of quasars (up to 1 in 130) are approaching the nitrogen enrichment levels of their more extreme counterparts.  If nitrogen-strong quasars are quasars viewed at the peak of metal enrichment, then the length of that phase is approximately 1/1700th of the typical quasar lifetime.  For example, for a quasar lifetime of 10$^7$ years, these objects would be viewed only in the last 6,000 years of their existence as quasars.  Alternatively, it may be that only 1 in 1700 quasars reaches the extremely high metallicities needed to generate strong nitrogen emission.  Further data and modeling outside the scope of this paper are needed to place these quasars in their correct context to the overall quasar population.  For example, if these quasars are near the end of their accretion phases, or are the most metal-rich because they are in the most massive bulges, their black holes should be biased towards higher masses than randomly selected quasars.  However, there is no evidence of such a bias in the widths of the emission lines, which range from narrow to average.  Additional observations should be undertaken in order to obtain higher S/N spectra for more accurate measurements of the emission lines, and also to push the observed wavelengths further into the blue in order to gain the Ly$\\alpha$ and \\ion{N}{5} emission lines for those objects with $z \\approx 2$.  With such additional data, metallicity estimates may be made using line ratios such as \\ion{N}{3}]/\\ion{O}{3}], \\ion{N}{3}]/\\ion{C}{3}], and \\ion{N}{5}/\\ion{He}{2}.  These estimates would allow us to test the hypothesis that nitrogen-rich quasars are exhausting their fuel supplies and approaching the metallicities predicted by numerical simulations for black holes as they end their quasar activity."
    },
    "0402/astro-ph0402496_arXiv.txt": {
        "abstract": "{We present a systematic study of the effect of metallicity on the stellar spectral energy distribution (SED) of O main sequence (dwarf) stars, focussing on the hydrogen and helium ionizing continua, and on the optical and near-IR lines used for spectral classification. The spectra are based on non-LTE line blanketed atmosphere models with stellar winds calculated using the {\\sc cmfgen} code of \\cite{hillier98}. We draw the following conclusions. First, we find that the total number of Lyman photons emitted is almost independent of line blanketing effects and metallicity for a given effective temperature. This is because the flux that is blocked by the forest of metal lines at $\\lambda < 600$ \\AA\\ is redistributed mainly within the Lyman continuum. Second, the spectral type, as defined by the ratio of the equivalent widths of \\HeI~$\\lambda$4471 and \\HeII~$\\lambda$4542, is shown to depend noticeably on the microturbulent velocity in the atmosphere, on metallicity and, within the luminosity class of dwarfs, on gravity. Third, we confirm the decrease in \\Teff\\ for a given spectral type due to the inclusion of line blanketing recently found by e.g. \\citet{martins02}. Finally, we find that the SED below $\\sim 450$ \\AA\\ is highly dependent on metallicity. This is reflected in the behaviour of nebular fine-structure line ratios such as \\neratio\\ 15.5/12.8 and \\arratio\\ 9.0/7.0 \\mum. This dependence complicates the use of these nebular ratios as diagnostic tools for the effective temperature determination of the ionizing stars in \\HII\\ regions and for age dating of starburst regions in galaxies. ",
        "introduction": "Massive stars are a dominant force in the evolution of the interstellar medium of galaxies. The extreme ultraviolet photons of massive stars can ionize surrounding atomic hydrogen gas. This ionized gas is heated by the photo electrons and cooled through forbidden line radiation of trace species such as oxygen, resulting in a temperature of some $10^4$~K. Besides this energetic coupling through the radiation field, massive stars can also influence their surroundings dynamically. The high pressure of the ionized gas will drive shock waves in their surroundings, sweeping up the ambient molecular gas. Moreover, these stars also have strong stellar winds and explode as type {\\sc ii} supernovae, both of which provide kinetic energy to the interstellar gas. Understanding the characteristics of massive stars and their interaction with their environment is therefore a key problem in astrophysics.  The stellar spectral energy distribution in the extreme ultraviolet (EUV) controls the ionization structure of \\HII~ regions. However, because of the high opacity of neutral gas in the EUV, this wavelength range cannot be observed directly. Spectral typing of stars is generally done through optical \\citep[e.g.][]{conti71,conti77,mathys88} or near-IR features \\citep[e.g.][]{hanson96,meyer98,hanson98,lenorzer02}. Alternatively, observed nebular ionization structures can be used to probe the ionizing fluxes of the stars energizing the medium. This technique is now widely used to determine the properties of stars in the EUV and has even evolved to a more general astronomical tool, particularly for the study of regions where individual stars cannot be directly typed \\cite[e.g.][]{oey00,takahashi00,okamoto01,morisset:paperiii}. Moreover, because the hardness of the EUV radiation field is a good measure of the spectral type of the ionizing star and because later spectral types live longer, the observed ionization structure can be used as an age indicator of a starburst region \\citep[e.g.][]{crowther99,thornley00,spoon00}.  In recent years, some problems with the latter technique have emerged. For example, for the well-studied \\HII~ region G29.96$-$0.02, the spectral type derived from direct observations of the stellar lines in the near-IR \\citep{watson97b,martin:G29} indicates a much earlier spectral type than that obtained from the ionization structure \\citep{morisset:paperiii}. This seems to be a general problem and may be related to the metallicity of the region. Indeed, there is a loose correlation between the ionization structure - as measured by for example the \\neratio~15.5/12.8 \\mum\\ line ratio - and the metallicity of \\HII~ regions \\cite[see][]{martin:metal}. Such a correlation may reflect the effects of metallicity on line blanketing or on the characteristics of the stellar wind.  Much theoretical effort has been dedicated to best describe the EUV spectra of massive stars. This is a formidable task because these stars have strong winds and extended atmospheres. This leads to strong non-LTE effects in the formation of spectral lines. Over the last ten years much progress has been made and current models include the effects of tens of thousands of lines on the energy balance and temperature structure of the stellar photosphere and wind. So far relatively modest effort has been investigated in systematic studies of the effects of metallicity on the stellar spectral energy distribution. In particular, there is no good theoretical understanding of the effects of metallicity on the ionizing fluxes of massive stars or on the optical and near-IR spectral characteristics used to type these stars.  Here, we study the influence of metallicity on the spectral energy distribution of O stars and determine its influence on the resulting ionization structure of \\HII\\ regions.  This paper is organized as follows. Sect.~\\ref{sect:seds:models} presents a set of main-sequence (dwarf) star models constructed using the {\\sc cmfgen} code by \\cite{hillier98} and compares the predicted EUV fluxes with those from other codes. Sect.~\\ref{sect:seds:description} presents a detailed analysis of the variations of the EUV spectral appareance and ionizing fluxes with effective temperature and metallicity.  Sect.~\\ref{sect:seds:sptype} investigates the influence of metallicity and other stellar parameters on the optical and near-IR spectral calibration. In Sect.~\\ref{sect:seds:hii}, the ionizing structure of single star \\HII\\ regions is studied. Finally, Sect.~\\ref{sect:seds:conclusions} presents the conclusions of this study. ",
        "conclusions": ""
    },
    "0402/astro-ph0402175_arXiv.txt": {
        "abstract": "I review of the observational properties of Soft Gamma Repeaters (SGRs) and Anomalous X-ray Pulsars (AXPs), two unusual manifestations of neutron stars.  I summarize the reasoning for SGRs being ``magnetars,'' neutron stars powered by the decay of a very large magnetic field, and the now compelling evidence that SGRs and AXPs are in fact members of the same source class, as predicted uniquely by the magnetar model.  I discuss some open issues in the magnetar model, and the prospects for future work. ",
        "introduction": "Since Baade \\& Zwicky made their now-famous 1934 prediction regarding the existence of neutron stars, these amazing objects have not ceased to surprise us in the variety of their observational manifestation.  Apart from thermal X-ray emission from initial cooling, predicted early on and now detected in a small handful of sources, the emission properties of neutron stars have been formally unpredicted, and informally unimagined.  The objects today being identified as ``magnetars'' are no exception. These sources literally exploded onto the astronomy scene in March 5, 1979, when the object today known as SGR~0525$-$66 emitted a soft-gamma-ray burst so intense that it saturated every gamma-ray detector that saw it (Mazets et al. 1979), likely measurably affected the Earth's ionosphere, and implied an awe-inspiring $>10^6$ Eddington luminosities.  This and the other handful of known ``Soft Gamma Repeaters'' (SGRs) prompted model explanations that ranged from vibrating neutron stars to strange star/pulsar phase transitions. Duncan \\& Thompson (1992), and quasi-simultaneously, Paczy\\'nski (1992), came up with the magnetar hypothesis, summarized below, which, particularly following seminal papers by Thompson \\& Duncan (1995, 1996), has uniquely stood the tests of increasingly constraining SGR observations.  They also identified ``Anomalous X-ray Pulsars'' (AXPs) as additional members of the magnetar club.  Though at the time having little in obviously common with SGRs, the AXPs, as we discuss below, have recently revealed themselves to be true siblings of the SGRs, with so many properties in common that the question to be answered today is ``what differentiates them from SGRs?'' ",
        "conclusions": "The discovery of SGR-like bursts from 1E 1048.1$-$5937 and especially 1E~2259+586 solidifies the common nature of AXPs and SGRs as predicted uniquely by the magnetar model.  This model has now made two major predictions, namely the spin-down of SGRs and the common nature of AXPs and SGRs, both of which have been unambiguously borne out by observations. The magnetar hypothesisthus appears to be very compelling.  However, there is still no direct evidence for the magnetar strength fields.  Such evidence could in principle be obtained from the detection of cyclotron lines in SGR or AXP spectra.  Spectral features detected in some SGR and AXP bursts may well be providing us with an important clue (Ibrahim et al. 2002; Gavriil et al. 2002) but their interpretation remains unclear as of yet.  Further, one expects a magnetar/radio pulsar connection,  This could come in two ways.  One is to detect radio pulsations from an AXP or SGR.  Such detections have been claimed but not confirmed (see paper by XXX, this volume).  However, detecting radio pulsations may be impossible; the long spin periods imply small polar caps, hence very narrow radio beams.  Alternatively, QED processes at high $B$, such as photon splitting, may preclude the electron/positron cascades necessary to produce radio emission (Baring \\& Harding 2001).  Another way to prove a magnetar/radio pulsar connection is to detect enhanced X-ray emission from a high-$B$ radio pulsar.  This has not yet been done even though several radio pulsars (see McLaughlin et al., this volume) have now been found having inferred $B$ comparable to or higher than that of 1E~2259+586, yet with no evidence for excess X-ray emission.  This is puzzling, but may simply reflect that our $B$ estimate is, in reality, not very accurate.  Continued discoveries of high-$B$ radio pulsars should prove interesting."
    },
    "0402/astro-ph0402343_arXiv.txt": {
        "abstract": "We have explored the relationship between the [O~III] $\\lambda$5007 and the 2--10 keV luminosities for a sample of Broad- and Narrow-Line Seyfert 1 galaxies (BLSy1 and NLSy1, respectively). We find that both types of Seyferts span the same range in luminosity and possess similar [O~III]/X-ray ratios. The NLSy1s are more luminous than BLSy1s, when normalized to their central black hole masses, which is attributed to higher mass accretion rates. However, we find no evidence for elevated [O~III]/X-ray ratios in NLSy1s, which would have been expected if they had excess EUV continuum emission compared to BLSy1s. Also, other studies suggest that the gas in narrow-line regions (NLR) of NLSy1s and NLSy1s span a similar range in ionization, contrary to what is expected if those of the former are exposed to a stronger flux of EUV radiation. The simplest interpretation is that, like BLSy1s, a large EUV bump is not present in NLSy1s. However, we show that the [OIII]/X-ray ratio can be lowered as a result of absorption of the ionizing continuum by gas close to the central source, although there is no evidence that intrinsic line-of-sight absorption is more common among NLSy1s, as would be expected if there were a larger amount of circumnuclear gas. Other possible explanations include: 1) anisotropic emission of the ionizing radiation, 2) higher gas densities in the NLR of NLSy1s, resulting in lower average ionization, or 3) the presence of strong winds in the the nuclei of NLSy1s which may drive off much of the gas in the narrow-line region, resulting in lower cover fraction and weaker [O III] emission. ",
        "introduction": "Osterbrock \\& Pogge (1985) discovered a class of Seyfert 1 galaxies characterized by relatively narrow ($FWHM$ $\\leq$ 2000 km s$^{-1}$) permitted lines, which they dubbed ``Narrow-Line Seyfert 1s''(NLSy1s). Compared to their broad-line counterparts (BLSy1s), the X-ray continua of NLSy1s are characterized by steeper slopes in the soft and hard bands (Boller, Brandt, \\& Fink 1996; Brandt, Mathur, \\& Elvis 1997) and more rapid variability (Turner et al. 1999). In a {\\it ROSAT} sample of active galaxies (AGN), Grupe et al. (1998) found similar optical to X-rays spectral indices for BL- and NLSy1s, which suggests that the soft X-ray steepness is indicative of excess emission, rather than weakness in the hard X-ray band. They also found that the optical and soft-X-ray spectral indices were anti-correlated suggesting that their continua peak in the EUV. There have been several models proposed to explain the more extreme properties of NLSy1s. For example, Wandel, Peterson, \\& Malkan (1999) suggested that NLSy1s possess a more distant Broad Line Region, due to over-ionization. Another possibility is that the BLR is flattened and viewed roughly face-on, with narrow line profiles as a result (e.g. Boller at al. 1996). The currently popular paradigm posits that, while AGN in general are powered by accretion of material onto a supermassive central black hole, NLSy1s possess black holes of relatively modest mass ($\\leq$ 10$^{7}$ M$_\\odot$), accreting matter at or above its Eddington limit (Pounds, Done, \\& Osborne 1995). The narrow line widths are then due to clouds in motion around the small-mass black hole, while the steep soft X-ray continuum is the high energy tail of the ``Big Blue Bump'', which is presumably emission from the accretion disk. Due to the high accretion rates (\\.{M}) , the disks in NLSy1s are likely much hotter than those in BLSy1s, hence the emission is peaked at higher energies. Wang \\& Netzer (2003) showed that high \\.{M} results in an extremely thin disk which emits a double-peaked continuum, with a peak in the EUV-soft X-ray due to the disk itself, and a high energy peak due to a hot corona, and, further, that most NLSy1s are super-Eddington accretors.  Wills et al. (1999) found unusually high N~V $\\lambda$1240/ C~IV $\\lambda$1550 flux ratios in NLSy1s, which may be indicative of high nitrogen abundances, as would result from metal enrichment due to a recent burst of star formation (see, also, Shemmer \\& Netzer 2002). This may indicate a different circumnuclear environment in NLSy1s compared to BLSy1s. Kuraszkiewicz et al. (2002) compared the UV spectra of NLSy1s and BLSy1s and determined that the former are characterized by weaker C~IV $\\lambda$ 1549 and C~III] $\\lambda$1909 and stronger Al~III $\\lambda$1857 emission. They suggested that the emission-line gas is of higher density and lower ionization than BLR gas in BLSy1s (see, also, Marziani et al. 2001). It has been noted (Mathur 2000, and references therein) that NLSy1s possess strong Fe~II and weak [O~III] $\\lambda$5007 emission relative to H$\\beta$. This puts them at one extreme of eigenvector 1 of Boroson \\& Green (1992), derived from principal component analysis of a large sample of low-redshift QSOs. In this regard, the NLSy1s do tend to show stronger optical Fe~II emission than BLSy1s (e.g., Gaskell 1985; Zheng \\& Keel 1991), however it has recently been suggested that this is a result of weaker H$\\beta$ emission (Veron-Cetty, Veron,\\& Goncalves 2001). On the other hand, by deconvolving the broad and narrow components of the H$\\beta$ profiles, Veron-Cetty et al. (2001) argued that the the ratio of [O~III] to narrow H$\\beta$ was similar to that of BLSy1s, which suggests that conditions in the Narrow-Line Regions (NLR), i.e, where the forbidden emission lines arise, are similar among the two groups of Seyferts. If the ionizing continuum in NLSy1s peaks more strongly in the EUV compared to BLSy1s, one would expect stronger emission from ions with high ionization potentials, such as [Fe~VII] $\\lambda$6087 and [Fe~X] $\\lambda$6374, relative to [O~III] $\\lambda$5007, however there is no strong evidence for this effect (Nagao et al. 2000). Furthermore, although Rodriguez-Ardila et al. (2000) found that the NLR of NLSy1s may be more highly ionized (or lacking some low-ionization gas) and that the [O~III] emission-line gas is somewhat denser than in BLSy1's, the two classes span the same range in [O~III] luminosity. Note, however, that these studies did not fully explore the relationship between the emission lines and the luminosity of the AGN. In this paper, we compare the luminosities of the [O~III] $\\lambda$5007 (L$_{[OIII]}$) to that in the 2--10keV band (L$_{2-10}$) for a sample of BLSy1s and NLSy1s, in order to determine if there is evidence for a stronger EUV flux in the latter. Specifically, if the conditions (density, cover factor, etc.) in the NLRs of BL- and NLSy1s are not significantly different, the [O~III] emission will be a reliable probe of their luminosities in the EUV, with higher EUV luminosities producing more [O~III] per hard X-ray luminosity. ",
        "conclusions": "As illustrated by the K-S tests, our sample of BLSy1 and NLSy1s exhibit similar distributions in X-ray luminosity (which, of course, is partly a selection effect, since we did not include LINERS and low-luminosity AGN in the sample), and we find a similar relationship between the [O~III] and hard X-ray luminosities. If the ionizing SEDs of NLSy1s were significantly different than BLSy1s, in the sense that they peaked in the EUV (e.g., Kuraszkiewicz et al. 2000), one would expect that they would have larger luminosities in ionizing photons compared to their hard X-ray luminosities. As a result, one would also expect higher emission-line luminosities relative to the hard X-ray, which is not the case. Hence, the simplest interpretation of these results is that, like BLsy1s (Laor et al. 1997), NLSy1s lack a strong EUV bump, contrary to the observational evidence (Boller et al. 1996). However, the effect of an EUV excess could be masked by other differences between BLSy1s and NLSy1s. For example, we have demonstrated that the presence of an intervening absorber which is optically thick above the He~II Lyman limit could modfiy the ionizing continuum sufficiently to drop the average ionization in the NLR of NLSy1s. This could explain why the NLR in NLSy1s does not appear to more ionized than that of BLSy1s (e.g. Nagao et al. 2000), and could be a consequence of high mass accretion rates. However, there is no direct evidence for a greater amount or distribution of circumnuclear gas in NLSy1s required by this scenario. Furthermore, although some unabsorbed NLSy1s lie close to the [O~III]/X-ray line from the high ionization NLSy1 model in Figure 1 (e.g. Mrk 478; Crenshaw et al. [1999]), Ton S180, which shows little evidence for intrinsic absorption (Turner et al. 2001), lies well below the BLSy1 model line.  A denser or more distant NLR in NLSy1s could also effectively mask the effects of a stronger EUV flux, however, there is no strong evidence for either of these conditions. A number of NLSy1s show blueshifted [O~III] relative to H$\\beta$ (e.g., I Zw 1; Phillips 1976) which may indicate strong winds in the NLR. This may result in lower NLR covering factors, which could drive the L$_{[OIII]}$/L$_{2-10}$ ratio down, but it is not clear how this might affect the average ionization of the gas, since the lower ionization cloud only occur if the emission-line clouds were confined such that their densities did not drop too rapidly with increasing radial distance. Another possible explanation is that the NLR gas in the NLSy1s is much more highly ionized that than in BLSy1s, with a smaller fraction of doubly-ionized oxygen. However, if this were the case, one would expect stronger high ionization lines, such as [Fe~X] $\\lambda$6374, and Nagao et al. (2000) did not find strong evidence of this. A number of NLSy1s show blueshifted [O~III] relative to H$\\beta$ (e.g., I Zw 1; Phillips 1976) which may indicate strong winds in the NLR. This may result in lower NLR covering factors, which could drive the L$_{[OIII]}$/L$_{2-10}$ ratio down. Finally, there is the possibility that the ionizing radiation is emitted anisotropically, hence the NLR gas is not exposed to the excess EUV radiation. However, there is no indication that the characteristics of NLSy1s can be the result of preferential viewing angle (Boller et al. 1996)  Based on this sample, the strength of L$_{[OIII]}$ relative to M$_{bh}$ supports the idea that there is a range of Eddington ratios among Seyfert 1s, with some of the NLSy1s occupying one extreme. While these results strongly suggest that the [O~III] emission is not unusual in NLSy1s (see, also, Veron-Cetty et al. 2001), perhaps other lines can be used to constrain the EUV continuum and, ultimately, the structure of the central engine. In fact, Veron-Cetty et al. (2001) suggest that it might be more advantageous to classify Seyferts based on the Fe~II emission. Clearly, more data are needed to probe the nature of the NLSy1 phenomenon. For example, a larger set optical and UV spectra will be invaluable for probing any differences in the emission-line gas that may not be apparent from [O~III] alone, and high resolution UV and X-ray spectra are required to help determine if the NLSy1 possess stronger intrinsic absorption."
    },
    "0402/astro-ph0402669_arXiv.txt": {
        "abstract": "This Letter explores influences of intracluster magnetic fields ($\\gsim 1\\mu$G) submerged in the hot electron gas on classic Sunyaev-Zel'dovich effect (SZE) and thermal bremsstrahlung in X-ray emissions. As the Larmor frequency is much higher than all collision frequencies, the presence of magnetic field may lead to an anisotropic velocity distribution of hot electrons. For the two-temperature relativistic Maxwell-Boltzmann distribution, we compute modifications to the classical thermal SZE. Intracluster magnetic fields tend to enhance the SZE with steeper radial variations, which bear important consequences for cluster-based estimates of cosmological parameters. By applying the magnetic SZE theory to spectral observations of SZ and Chandra X-ray emissions from the galaxy cluster Abell 2163, a $\\sim 30-40\\mu$G central core magnetic field $B_0$ is predicted. For the SZ and Chandra X-ray spectral observations of the Coma cluster, our theoretical analysis is also consistent with an observationally inferred $B_0\\lsim 10\\mu$G. As the magnetic SZE is redshift $z$ independent, this mechanism might offer a potentially important and unique way of probing intracluster magnetic fields in the expanding universe. ",
        "introduction": "The Sunyaev-Zel'dovich effect (SZE) in galaxy clusters offers a unique and powerful observational tool for cosmological studies. There has been persistent progress in detecting and imaging the SZE in clusters. In view of this rapid development in SZE observations, several important physical effects associated with the classical thermal and kinetic SZEs (Sunyaev \\& Zel'dovich 1969, 1980) have been further explored for their diagnostic potentials, such as relativistic effects (Rephaeli 1995), the shape and finite extension of a galaxy cluster with a polytropic temperature (Puy et al. 2000), halo rotation SZE (Cooray \\& Chen 2002; Chluba \\& Mannheim 2002), Brillouin scattering (Sandoval-Villalbazo \\& Maartens 2001), early galactic winds (Majumdar et al. 2001) and cooling flows (Schlickeiser 1991; Majumdar et al. 2001; Koch et al. 2002).  Intracluster magnetic fields have been measured using a variety of techniques and diagnostics, including synchrotron relics and halo radio sources within clusters, inverse Compton X-ray emissions from clusters, Faraday rotation measures of polarized radio sources either within or behind clusters, and cluster cold fronts in X-ray images (Clarke et al. 2001; see Carrili \\& Taylor 2002 for a recent review). These observations reveal that most cluster atmospheres are substantially magnetized with typical field strengths of $\\gsim 1\\mu$G and with high areal filling factors out to Mpc radii. In the cores of `cooling flow' clusters (Eilek \\& Owen 2002; Taylor et al. 2001) and at cold fronts of merging clusters (Vikhlinin et al. 2001), magnetic fields may gain intensities of $\\sim 10-40\\mu$G and thus become dynamically important.  Magnetic fields in the intracluster gas allows for particle acceleration processes which modify specifics of heating processes, such that the electron energy distribution differs from the Maxwell-Boltzmann distribution. Such stochastic acceleration processes include shocks and magnetohydrodynamic (MHD) waves, etc. The bremsstrahlung from a modified Maxwell-Boltzmann electron gas might account for the observed X-ray spectra up to highest energies of current X-ray observations (Ensslin et al. 1999; Blasi 2000a). If energy injections by MHD waves are turned off, a galaxy cluster gradually thermalizes with electrons approaching a Maxwell-Boltzmann distribution on a rough timescale of $\\sim 10^7-10^8$ yrs. As all collision frequencies (Nicholson 1983) are much lower than the electron Larmor frequency for the magnetized intracluster gas (Sarazin 1988), the electron velocity distribution is likely to be anisotropic as long as the parallel (relative to magnetic field ${\\bf B}$) pressure is not too much higher than the perpendicular pressure (Parker 1958; Hasegawa 1975).  We presume the result of a partial electron thermalization is a two-temperature relativistic Maxwell-Boltzmann distribution, i.e. an anisotropic velocity distribution. This two-temperature does not mean an electron gas having two components with different temperatures, but refers to the same population with different average kinetic energies along and perpendicular to the local magnetic field. The main thrust of this Letter is to advance magnetic SZE theory in contexts of Chandra X-ray and radio SZE spectral observations for galaxy clusters.  In Section 2, we calculate the X-ray emission and SZE spectra using the two-temperature relativistic Maxwell-Boltzmann distribution for electron velocity. Based on both Chandra X-ray and SZ spectral observations, we offer a specific prediction for the galaxy cluster A2163. Finally, we discuss cosmological implications of our magnetic SZE theory in Section 3. ",
        "conclusions": "Contrary to recent results (Koch et al. 2003; Zhang 2003), we find that the anisotropic velocity distribution of electrons caused by magnetic field $B$ enhances the SZE. Our model results of Figs. $1-3$ can be critically tested against more precise spectral SZE measurements of A2163 in the frequency bands of $\\sim 50-130$GHz and $\\sim 300-600$GHz by MAX, MSAM and SuZIE types of experiments in the frequency passbands of $90-670$GHz and by AMiBA in the band $84-104$GHz, Nobeyama at 21 and 43GHz, JCMT at 350 and 650GHz, SZA in the bands of $26-36$GHz and $85-115$GHz, BIMA and OVRO in the band of $26-36$GHz, MINT at 150GHz and ACT at 150, 220, and 270 GHz. Multi-frequency projects such as the upgraded MITO (Lamagna et al. 2002; De Petris et al. 2002) and the OLIMPO (Masi et al. 2003) experiments are very promising to provide some results.  A2163 may involve merger shocks that could amplify $B$ (e.g. Markevitch \\& Vikhlinin 2001). The spectral index maps of A2163 show a spectral steepening from the central to peripheral radio halo regions, implying a radial decrease of $B$ in reacceleration models (e.g. Feretti et al. 2003). It was attempted to fit the SZE spectrum of A2163 with a combination of thermal and non-thermal electrons (e.g. Colafrancesco et al. 2003), but no evidence was found for hard X-ray excess due to the non-thermal component in the BeppoSAX data (e.g. Feretti et al. 2001).  Based on X-ray and SZE measurements, 41 galaxy clusters were used to independently estimate Hubble constant $h_{100}=0.61\\pm 0.03\\pm 0.18$, where the uncertainties are statistical and systematic at 68\\% confidence level for $\\Omega_M=0.3$ and $\\Omega_{\\Lambda}=0.7$ cosmology (Carlstrom et al. 2002; Reese 2003). Our analysis of A2163 shows that intracluster magnetic field induces microscopic anisotropies in electron velocity distribution to enhance the SZE. It appears that inferences from cluster models without magnetic field would systematically underestimate $h_0$ as in the case of A2163 for which Holzapfel et al. (1997) inferred a lower $h_{100}=0.60\\pm 0.04$ against the current WMAP result of $h_{100}=0.71\\pm 0.04$. As the cluster asphericity and orientation in the sky are random and the average cluster peculiar velocity is zero, these factors should contribute to the systematic uncertainty with the Hubble constant being statistically unaltered. The underestimation of Hubble constant may be explained by the generic presence of core magnetic field $B_0\\sim 10-40\\mu$G in this sample of galaxy clusters.  Another important cosmological effect of the ubiquitous enhancement of magnetic SZE due to the prevalence of $\\gsim 1\\mu$G magnetic fields in galaxy clusters would be observable in the CMB angular spectrum especially at high $l\\gsim 3000-4000$. This contribution to CMB fluctuations may be estimated and tested by CMB experiments such as ACT, Planck, SZA, etc. Details of these two cosmological effects will be pursued in forthcoming papers.  For X-ray (Arnaud et al. 2001) and SZE (De Petris et al. 2002) spectral observations of Coma cluster (Abell 1656), our magnetic SZE analysis is consistent with the currently inferred $B_0\\lsim 10\\mu$G (Carilli \\& Taylor 2002). Likewise, magnetic SZE can be utilized in other galaxy clusters with high-resolution and high-sensitivity X-ray and SZE spectral observations to estimate the lower limit of $B_0$ as well as SZE spatial features. While it is necessary to estimate all possible corrections to the classic SZE in order to isolate the magnetic contribution, this may be a unique procedure to probe intracluster magnetic fields at high redshifts $z$, at least statistically. Finally, anisotropic distributions of nonthermal electrons should lead to distinct magnetic SZE in radio lobes of extragalactic jets. \\\\  We thank the referee Dr. D. Puy for useful comments. This research has been supported in part by the ASCI Center for Astrophysical Thermonuclear Flashes at the U. of Chicago under DOE contract B341495, by the Special Funds for Major State Basic Science Research Projects of China, by the THCA, by the Collaborative Research Fund from the NSF of China for Outstanding Young Overseas Chinese Scholars (NSFC 10028306) at the National Astronomical Observatory, CAS, by NSFC grant 10373009 at the Tsinghua U., and by the Yangtze Endowment from the Ministry of Education through the Tsinghua U. Affiliated institutions of YQL share this contribution."
    },
    "0402/astro-ph0402613_arXiv.txt": {
        "abstract": "We report on four years of multiple wavelength observations of the RS~CVn system V711~Tau (HR~1099) from 1993, 1994, 1996 and 1998.  This combination of radio, ultraviolet (UV), extreme ultraviolet (EUV), and X-ray observations allows us to view, in the most comprehensive manner currently possible, the coronal and upper atmospheric variability of this active binary system.  We report on the changing activity state of the system as recorded in the EUV and radio across the four years of observations, and study the high energy variability using an assemblage of X-ray telescopes. We find: \\\\ (1) evidence for coherent emission at low radio frequencies ($\\leq$ 3 GHz) which appears to be both highly time variable and persistent for several hours. Such phenomena are relatively common, occurring $\\approx$ 30\\% of the time HR~1099 was observed at L-band. The measured polarizations of these bursts are left circularly polarized, in contrast with behavior at higher frequencies which has the opposite helicity.  The polarizations are consistent with a variable source that is 100\\% left circularly polarized, along with a steady level of flux and polarization which is 0 or slightly right circularly polarized.  There appears to be a low degree of correlation between bursts at 20 cm and higher frequency gyrosynchrotron flares, and also between 20 cm bursts and large EUV/soft X-ray (SXR) outbursts. \\\\ (2) Higher frequency (5--8 GHz) flares show an inverse relationship between flux and polarization levels as the flare evolves; this behavior is consistent with flare emission which is initially unpolarized. Large variations in spectral index are observed, suggesting changes in optical depths of the flaring plasma as the burst progresses. Quiescent polarization spectra show an increase of polarization with frequency, a pattern typically seen in active binary systems but still not understood. \\\\ (3) EUV observations reveal several large flares, in addition to numerous smaller enhancements.  The total range of variability as gleaned from light curve variations is only a factor of 7, however. Observations in different years provide evidence of a change in the quiescent, not obviously flaring, luminosity, by a factor of up to 2.  From an analysis of time-resolved spectral variations, we are able to infer evidence for the creation of high-temperature plasma during flare intervals compared with quiescent intervals.  Interpretation of EUV spectral variations is hindered by the lack of ability to diagnose continuum levels and activity-related abundance changes, which are known from higher energy observations. Electron densities determined by line ratios of density-sensitive emission lines are high (10$^{12}$--10$^{13}$ cm$^{-3}$) and there is no evidence for large density enhancements during flare intervals, compared with quiescent intervals.  \\\\ (4) X-ray observations reveal several flares, and provide evidence of energy-dependent flare evolution:  harder X-ray energies show faster temporal evolution than at softer energies.  Time-resolved X-ray spectral analysis shows the presence of hot plasma,  T$_{e}\\sim$ 30 MK, during flares compared to quiescent intervals, as well as evidence for changing abundances during flares.  The abundance of iron (which is subsolar) shows an enhancement of a factor of three at the peak of a moderate flare seen by {\\it ASCA} relative to the pre-flare level; abundances decrease during the flare decay. No hard ($> 15$ keV) emission is detected by either {\\it RXTE} or {\\it BeppoSAX}.\\\\ (5) The luminosity ratios L$_{EUV}$/L$_{R}$ in quiescence determined from several time intervals during the four campaigns are consistent with previously determined ratios from a sample of active stars and solar flares.  The range of L$_{EUV}$/L$_{R}$ from three EUV/radio (3.6 cm) flares is the same as the values obtained during quiescence, which points to a common mechanism for producing both flaring and not flaring emission.\\\\ (6) Seventeen flares were observed in the EUV and$/$or SXR during the four campaigns; of the eight flares that had radio coverage, three show 3.6 cm radio flares, which are generally consistent with the Neupert effect.  Five EUV$/$SXR flares had partial UV coverage; all show UV responses, particularly in the C~IV transition. The UV flare enhancements can occur at the same time as the 3.6 cm radio flares, in two cases where radio, UV, and EUV/SXR flare coverage overlapped.\\\\ (7) For SXR flares, we find that the contrast between flare emission and quiescent emission increases as expected towards higher energies, making flare detections easier at harder X-ray energies.  This is due to the creation of high temperature plasma during flares, which shows up predominantly in high energy continuum emission. We find a discrepancy between the implied flaring rate based on EUV observations, and higher energy observations; the lower energies tend to miss many of the flares, due to the lack of sufficient contrast with quiescent emission. \\\\ ",
        "introduction": "As a bright and variable system in almost all wavelength regions, HR~1099 (V711~Tau; HD~22468) has been the subject of many past observations from radio to X-ray wavelengths. HR~1099 is a short-period (P$_{\\rm orb} =$ 2.83774 d) binary composed of a G5 dwarf and a K1 subgiant. The system lies at a Hipparcos distance of 29 pc \\citep{hipparcos}. The orbital and rotational periods of the stars are tidally synchronized. The mass ratio is 1.3, yet the K subgiant outsizes the G subgiant by a factor of three in radius.  The orbital inclination is $\\approx$ 33$^{\\circ}$ \\citep[][and references therein]{cabs} and there are no eclipses. Because of its spectroscopic characteristics and highly active chromospheric, transition region, and coronal emissions, HR~1099 is a member of the RS~CVn class of binary systems. Many of the phenomena seen on RS~CVn systems have solar counterparts, and invite the comparison between ``hyper-active'' stars such as these and less active stars like the Sun.  A schematic of the generally accepted model for solar flares proceeds as follows: A flare begins with a source of free energy, thought to originate from magnetic reconnection high in the solar corona.  Some of this energy is used to accelerate electrons to moderately relativistic speeds. Electron beams can be generated, and propagate outward or into the atmosphere; the electron beams emit plasma radiation (type III bursts), whose frequencies trace the ambient plasma density the beams encounter as they propagate either outward or into the atmosphere \\citep{aschwandenbenz1997}.  Downward-directed accelerated electrons can become magnetically trapped in the coronal loop or arcade, and emit gyrosynchrotron emission at microwave frequencies. As the electrons impact the denser regions of the lower atmosphere, they deposit their energy, emitting hard X-rays via nonthermal bremsstrahlung emission and white light continuum emission at the footpoints of the loops. The energy deposited by the electrons in the lower atmosphere heats plasma to coronal temperatures on a timescale short compared to the hydrodynamic expansion time, ablating material at 10$^{6}$--10$^{7}$ K up the coronal loops, where it emits soft X-ray radiation.  As the coronal density increases coronal material can stop the energetic electrons higher up in the atmosphere thereby heating the corona directly. Once the nonthermal energy input has ceased, the material in the loop condenses into the chromosphere and the soft X-ray emission returns to its preflare state.  This scenario implies a set of related and correlated variations that should be observable in different spectral regions.  Despite the relatively advanced state of understanding of multi-wavelength solar flare emissions, there is a paucity of observations on the stellar side.  Part of this may result from the difficulties inherent in organizing such an observing campaign for stellar flares; but the potential benefits to understanding stellar flares must surely outweigh the enormous effort needed to coordinate such observations.  \\citet{gudelbenz} showed that there existed an almost linear relationship between stellar quiescent radio and soft X-ray emission, suggesting that coronal heating and particle acceleration are closely linked in a way that may be similar to solar flares.  In one of the earliest studies of multi-wavelength emission, \\citet{weileretal} examined optical, UV, and radio observations of two RS~CVn systems (one of them HR~1099) and said ``there are suggestive coincidences between peak radio flux density and optical-UV emission activity from both systems.''  Other studies of multi-wavelength stellar flares have revealed a zoo of phenomena, leading to split opinions as to whether such correlations even exist.  In one of the earliest simultaneous studies, \\citet{kundu1988} examined radio and X-ray observations of four flare stars.  While they noted some X-ray bursts coincided in time or were preceded by 10--15 minutes by 20 cm radio flares, they still claimed the degree of correlation was low.  \\citet{foxetal1994} observed the decay of a radio flare on the RS~CVn system EI~Eri but did not see an accompanying X-ray or EUV flare. Yet \\citet{stern1992} found UV and microwave flaring occuring during an X-ray outburst on the RS~CVn system $\\sigma^{2}$~CrB.  \\citet{gagne} investigated the M dwarf binary system EQ~Peg with radio, optical, EUV and X-ray wavelengths and found two populations of X-band (3.6 cm) flares:  highly polarized flares with no counterparts at shorter wavelengths, and moderately polarized flares which do have shorter wavelength counterparts.  Recently, \\citet{ostenetal2000} established a correlation of radio flares with X-ray flares on $\\sigma^{2}$~CrB, with X-ray flares peaking up to 1.4 hours before the radio peak.  And most recently, \\citet{ayres2001} conducted coordinated UV, EUV, X-ray and radio observations on HR~1099 and detected a large UV flare not seen in higher energy emissions.  The flare mechanism is a poorly understood phenomenon, even when considered in a single wavelength region (or on a very well-studied star like the Sun). Expectations that stellar flares mimic the behavior of solar flares can introduce biases into the interpretation of multi-wavelength studies.  Also, it is necessary to recognize the importance of time delays between different wavelength regions when studying flares in a multi-wavelength context and the necessity of long exposure times to catch and observe flares in their entirety.  This paper describes the results of four campaigns which observed HR~1099 in multiple wavelength regions, in 1993, 1994, 1996, and 1998.  A summary of these campaigns is given in Table~\\ref{table1}.  Section~\\ref{sec:prevobs} summarizes previous observations of HR~1099; section~\\ref{sec:ch5datred} describes the observations and initital data reduction; Section~\\ref{sec:anal} examines the wavelength regions individually across the time span of these four campaigns; \\S~\\ref{sec:multi} investigates comparisons between the different wavelength regions investigated; and \\S~\\ref{sec:ch5conc} concludes. ",
        "conclusions": "There are several conclusions to be drawn from such a broad investigation of flaring on the active binary system HR~1099.  One of the surprising results obtained from radio observations is how common bursts at low frequencies are on HR~1099 -- occurring roughly one third of the time. We have investigated the nature of L-band bursts, and find that they are characterized by emission that can achieve large values of percent left circular polarization ($\\rightarrow$ 100\\%), is highly variable (possibly varying on timescales less than the 10 second integration time of the VLA observations) and long-lasting (attaining high flux levels and high degrees of left percent circular polarization for hours). Due to an inability to constrain the source size of these bursts, only a moderate lower limit on the brightness temperature, T$_{b} \\geq$ 10$^{10}$K, can be determined.  Such brightness temperatures could be consistent with incoherent gyrosynchrotron emission; coupled with large values of circular polarization, short timescales and apparent lack of similarity with higher frequency variations, however, it is more likely that this phenomenon is an example of a coherent mechanism.  Determining what mechanism might be operating in the coronae of HR~1099 requires observations with better time resolution (to constrain the source size), and larger frequency coverage (to determine the bandwidth of the emission).  The characteristics of X-ray and microwave emission from active binary systems and dMe stars share similarities. The thermal coronal emission from both is characterized by electron temperatures T$_{e}$ of 5--30 MK, high densities (n$_{e} \\approx$ 10$^{12}$ cm$^{-3}$), and coronal abundances that are subsolar in quiescence and appear to vary during flares. Centimeter-wavelength radio emission is generally attributed to nonthermal gyrosynchrotron radiation outside of flares \\citep{gudelbenz1996}. The phenomena observed on HR~1099 at 20 cm reported here, and on other active binary systems reported by others, has parallels with behavior on dMe stars: in both cases, there is evidence for highly polarized and time variable emission (dMe star behavior shows evidence for fast fluctuations [$\\Delta t \\leq$ 20 milliseconds] and extremely high brightness temperatues, [$T_{b} \\geq$ 10$^{16}$K]). Constraining the observational properties of these low-frequency bursts on active binary systems will reveal the extent of the apparent similarity in dynamics of these two kinds of coronal environments.  Radio observations show several large flares at 3.6 cm. The increase of flux and decrease of polarization during 3.6 cm flares is consistent with a constant quiescent level in intensity and polarization and a flare which increases in flux but remains unpolarized.  The 6--3.6 cm spectral indices during the flares show an increase during the flare rise, attaining the largest value as the flux reaches its maximum value \\citep[see also][]{richardsetal}.  This behavior is consistent with an increase in optical depth during the flare rise, and also consistent with little or no circular polarization during the flare; it is difficult to obtain large values of circular polarization under optically thick conditions. The behavior of the radio flux spectra illustrate the general trends of active binary gyrosynchrotron emission:  Spectral indices are flat or slightly negative during quiescence, and show a peak between 2 and 5 GHz during flares. The observed polarizations always increase with frequency, regardless of the behavior of the flux spectra with frequency \\citep[as also discussed in][]{whitef1995}.  There were no definite detections of HXR emission at energies larger than $>$ 12 keV during or outside of flares.  Yet the radio fluxes indicate the existence (and persistence) of accelerated particles.  To date, there have been only a few detections of hard X-ray (20--50 kev) emission by {\\it BeppoSAX} at the peak of very strong flares \\citep{favataschmitt1999,maggioetal2000,pallavic2001,francioetal2001}. This HXR emission however has been interpreted as thermal emission due to the very high temperatures (T $\\sim 10^8$ K) reached in these flares. Our {\\it BeppoSAX} PDS observation of HR 1099 gives only an upper limit to hard X-ray emission which is two orders of magnitude lower than the soft X-ray emission. The fact that no HXR emission has been detected by either {\\it RXTE} or {\\it BeppoSAX} is due to the fact that the temperatures reached by the observed X-ray flares on HR~1099 are much lower than those needed (about 100 MK) for thermal hard X-ray emission; and secondly, to the fact that the nonthermal electrons indicated by the radio observations are apparently insufficient to produce detectable non-thermal hard X-ray emission (as is also true for the Sun where the ratio of non-thermal to thermal hard X-ray emission in impulsive solar flares is $\\geq$ 10$^{-5}$).  Flares lasting several days appear to be a common feature on active stars.  Such long-duration flares have now been detected not just on active binary systems \\citep{kursterschmitt1996}, but on active evolved single giant stars \\citep{ayresetal1999,ayresetal2001} as well as higher gravity dMe flare stars \\citep{cullyetal1994}.  A feature common to many of these flares is the presence of a change in the light curve decay phase, from an initially fast decay to a slower decay.  Several flares on HR~1099 discussed in this paper also show evidence for such decay morphologies.  Unfortunately, the sensitivities of the EUV and SXR telescopes which observed these flares are not sufficient to determine the change of plasma parameters (n$_{e}$, T$_{e}$) with time during the decay phase; usually only a gross estimate of the total flare changes with respect to quiescent intervals is possible.  This renders an interpretation of the underlying cause of the observed decay change difficult:  a change to low-density structures during the late stages of the decay could explain the long decay timescales.  Another viable explanation is the presence of continued heating during the decay \\citep[e.g.][]{favataschmitt1999}. A possible clue may come from studying solar flares; a few long-duration solar events \\citep[e.g.][]{feldman1995} appear to possess a break during the decay, similar to the behavior seen in stellar long-duration events.   {\\it EUVE} observations reveal many flares; HR~1099 was in a high state of activity for three out of the four years it was observed during these campaigns.  There is suggestive evidence of a flare precursor  in four large EUV flares, consisting of a slow increase in flux before the flare rise.  The dynamic range of events exhibited by HR~1099 was less than a factor of 7 in the EUV.  Most flares in the EUV show a remarkable amount of symmetry between the flare rise and decay, an observation that is at odds with the fast rise and slow decay typical of chromospheric evaporation in solar flares.  An investigation of the time-resolved spectral variations confirms the creation of hot plasma during flare intervals. The existence of many weak lines blended with continuum radiation in the EUV hinders spectral interpretation, particularly with regard to elemental abundances and activity-related abundance changes.  This also limits the hottest temperatures to which EUVE spectra are sensitive. Densities determined from {\\it EUVE} spectra indicate high values (n$_{e}$ $\\sim$ few 10$^{12}$ cm$^{-3}$) during quiescence, but no evidence for a statistically significant enhancement during flares.   {\\it ASCA} observations of a small flare indicate an enhancement of the iron abundances by about a factor of three between quiescence and the peak of the flare.  The emission measure distribution shows the same general structure during quiescence and flares: A peak at 6--10 MK and another between 20 and 30 MK.  During flares, the high temperature peak becomes more prominent and moves to hotter temperatures. The temperatures derived from {\\it BeppoSAX} and {\\it ASCA} spectra are generally consistent with each other; however, the small flares observed with {\\it BeppoSAX} did not show any evidence for abundance changes. The behavior in different SXR bands as investigated with {\\it ASCA} and {\\it RXTE} reveal faster changes at the harder energies, indicative of hotter temperature plasma.  The multi-wavelength flare data illustrate fair correlation between EUV$/$SXR flares and radio/UV bursts. This can be interpreted in the light of solar flare mechanisms using the Neupert effect, where the time integral of the nonthermal radiation is proportional to the rise phase of the EUV$/$SXR light curve. The correspondence of radio and UV light curves suggests that the UV line fluxes also can be used as a proxy for the the nonthermal radiation impinging on the chromosphere. Three EUV$/$SXR flares showed a radio burst during the flare rise, with time delays between peak radio and EUV$/$SXR of 2.5--30 hours. These values are much larger than typical values for solar flares, where the delay is on the order of minutes. The highly polarized emission at 20 cm shows a low degree of correlation with gyrosynchrotron flaring activity at 3.6 cm. The ratios of flare luminosities in the EUV and 3.6 cm bandpasses indicate a remarkable amount of similarity with ratios derived from simultaneous observations of quiescence in the two bandpasses.  This suggests that a similar mechanism which forms the quiescent thermal/nonthermal radiation is also present, at some level, in producing the time-varying flare emission. The radio and EUV flares considered here last orders of magnitude longer than typical solar flares, where nonthermal radio radiation lasts $\\sim$ minutes, and EUV/SXR thermal radiation lasts $\\sim$ hours. Yet, the ratios of EUV-to-radio flare duration for the three flares discussed are in agreement with the solar and stellar flares discussed by \\citet{gudeletal1996}.  There is a noticeably greater enhancement in the flare luminosity compared with the quiescent luminosity for higher energy bandpasses.  For times when EUV and X-ray observations are simultaneous, often flares are obvious in the X-ray light curve but only appear as ``quiescent'' modulations in the EUV. More luminous flares tend to be hotter, and thus the flare spectral energy distribution will be shifted to higher energies, and involve predominantly continuum emission; flares will then favor the X-ray spectral regions over the EUV. This can lead to a bias that lowers the observed flare frequency, and affects the distribution of flares with energy.  This work was supported by NASA grants NGT5-50241, NAG5-7020, NAG5-3226, NAG5-2259, NSG5-4589, NAG5-2530, NAG5-7398, and NSF grant AST-0206367 to the University of Colorado. RAO is grateful for the support of a GSRP fellowship. This represents the results of VLA projects AB691, AB719, AB793, and AB874, and ATCA projects C302, C370, and C546. EF, RP and GT acknowledge partial support from the Italian Space Agency (ASI).  We gratefully acknowledge Keith Jones (University of Queensland, Australia; retired) for his effort in obtaining the early ATCA observations, and Bryan Deeney for his contribution in the IUE data analysis.  We thank the referee for a careful reading of this lengthy paper."
    },
    "0402/astro-ph0402049_arXiv.txt": {
        "abstract": "% We reinvestigate the UV spectrum of NGC\\,1535 by means of recently developed fully line-blanketed non-LTE models. These new models account for the wind in spherical geometry while handling the atomic data in a very similar way to the {\\sc Tlusty} code. This approach ensures at the same time realistic predictions of the photospheric absorption lines and of the emission lines formed in the wind. Our analysis confirms the results of previous studies. We derive $T_*=70\\,{\\rm kK}$, {\\raisebox{0.24ex}{$\\stackrel{{\\textstyle\\raisebox{-0.2ex}{.}}}{M}$}} $=10^{-7.8}\\,{\\rm M_\\odot/yr}$, and $v_\\infty$=2000\\,km/s. ",
        "introduction": "NGC\\,1535 is a planetary nebula with a hot central star which belongs to the group of Hydrogen-rich objects with weak winds. The stellar parameters and abundances of these stars can be determined from plane-parallel, static models by fitting their stellar absorption lines (c.f. Werner et al. 2003). A small subset of this group, to which NGC\\,1535 belongs, shows P-Cygni profiles in the UV which allows for an investigation of their weak stellar winds by means of spherical wind models.  The stellar parameters and the wind properties of NGC\\,1535 have been well established by Mend\\'{e}z et al. (MKH, 1992), Perinotto (P, 1993), Tinkler \\& Lamers (TL, 2002) and Bauer \\& Husfeld (BH, 1995). In the present paper we confirm these results by applying (for the first time) a new non-LTE code for expanding stellar atmospheres to recently observed HST and FUSE spectra. ",
        "conclusions": ""
    },
    "0402/astro-ph0402080_arXiv.txt": {
        "abstract": "The time dependence of the dark energy density can be an important clue to the nature of dark energy in the universe.  We show that future supernova data from dedicated telescopes (such as SNAP), when combined with data of nearby supernovae, can be used to determine how the dark energy density $\\rho_X(z)$ depends on redshift, if $\\rho_X(z)$ is not too close to a constant.  For quantitative comparison, we have done an extensive study of a number of dark energy models. Based on these models we have simulated data sets in order to show that we can indeed reconstruct the correct sign of the time dependence of the dark energy density, outside of a degeneracy region centered on $1+w_0 = -w_1 z_{max}/3$ (where $z_{max}$ is the maximum redshift of the survey, e.g., $z_{max}=1.7$ for SNAP).  We emphasize that, given the same data, one can obtain much more information about the dark energy density directly (and its time dependence) than about its equation of state. ",
        "introduction": "Most of the energy in our universe is of unknown nature to us.  The amount of this dark energy has been determined by recent experiments, including the Wilkinson Microwave Anisotropy Probe (WMAP) satellite observations \\cite{wmap} of the anisotropy in the cosmic microwave background radiation.  Our universe is spatially flat (the three-dimensional equivalent of a two-dimensional plane), with roughly 27\\% matter and 73\\% dark energy.  Determining the nature of this dark energy is one of the major fundamental challenges in astronomy and physics today.  There are many plausible candidates for dark energy.  For example, (1) a cosmological constant, i.e., constant vacuum energy originally proposed by Einstein in his equations of general relativity, (2) a time dependent vacuum energy, or scalar field known as ``quintessence'', that evolves dynamically with time\\cite{yaya}, or (3) modified Friedmann equation, e.g. the Cardassian models \\cite{card,mpcard,card03}, that could result as a consequence of our observable universe living as a 3-dimensional brane in a higher dimensional universe.  Other proposed modifications to the Friedmann equation include \\cite{others}.  The time-dependence of the density of dark energy can reveal the nature of dark energy at a fundamental level.  A powerful probe of dark energy is type Ia supernovae (SNe Ia), which can be used as cosmological standard candles to measure how distance depends on redshift in our universe. Observations of SNe Ia have revealed the existence of dark energy in the universe \\cite{SN1,SN2}. Current SN Ia data are not yet very constraining on the nature of the dark energy \\cite{Wang04}.  The distance-redshift relation of observed supernovae depends on the nature of dark energy.  Most researchers have chosen to parametrize dark energy by its equation of state parameter. However, it has been shown \\cite{mbs,barger} that it is extremely difficult to constrain the time-dependence of the dark energy equation of state using supernova searches (or any other technique relying on the luminosity distance); hence one might worry that one cannot differentiate between different dark energy models.  Fortunately, it has been shown that one can do much better if one parametrizes the dark energy by its density directly, instead of its equation of state \\cite{Wang01a,Wang01b,Tegmark02,Daly03}. The dark energy density is a more fundamental parameter than the dark energy equation of state parameter.  Obtaining the equation of state parameter requires one to perform an additional integral (compared to obtaining the dark energy density); this integral smears out much of the information one could otherwise learn.  Hence, given the same data, the uncertainties of the constraints on the dark energy density should be {\\it smaller} than that of the constraints on the dark energy equation of state.  In this paper we focus on extracting information about the dark energy density directly.  We will show how well, using future supernova data, one can determine whether the dark energy density changes with time, and whether it increases or decreases with time.    We begin in Section II with the basic equations for using supernovae to study dark energy.  In Section III we present four theoretical models which we will study to see how well we can reconstruct the time dependence of the dark energy: Models 1, 2, and 3 have dark energy density that is constant, increasing, and decreasing in time respectively. As our fourth set of models we consider those parametrized by an equation of state $w_X(z)=w_0+w_1 z$.  In Section IV, we simulated SNIa data for these models.  In Section V, we use the adaptive iteration method to see how well we can reconstruct the time dependence of the dark energy density for these models: we use three test functions with different time dependences to see which one best matches the data.  For each model we then run 1000 Monte Carlo samples to obtain error bars for our fit.  The results are presented in Section VI, followed by the conclusions. ",
        "conclusions": "\\label{sec:7}  We have investigated how well future supernova data from dedicated telescopes (such as SNAP), when combined with data of nearby supernovae, can be used to determine the time dependence of the dark energy density. For quantitative comparison, we have done an extensive study of a number of dark energy models, with dark energy density that is constant, increasing, and decreasing in time. Based on these models we have simulated data sets in order to show that we can indeed reconstruct the correct sign of the time dependence of the dark energy density.  Among the dark energy models we studied are those parametrized by an equation of state $w_X(z)=w_0+w_1 z$.  Here, $w \\equiv p/\\rho$.  We studied a grid of 147 models, for $ -1.2 \\leq w_0 \\leq -0.5$, and $-1.5 \\leq w_1 \\leq 0.5$.  We emphasize that it is the dark energy density that we reconstructed, {\\it not} the equation of state.  We find that there is a degeneracy region in the ($w_0, w_1$) parameter space centered near $1+w_0 = -w_1 z_{max}/3$ (where $z_{max}$ is the maximum redshift of the survey, e.g., $z_{max}=1.7$ for SNAP); the models that lie within this region cannot be differentiated from a $\\Lambda$ model even if $\\Omega_m$ is known independently to 1\\% accuracy (we compute the size of the region for $\\Omega_m$ known to varying degrees of accuracy).  Outside of this degeneracy region, we can detect the time variation of the dark energy density at 1$\\sigma$ or higher significance levels.  We emphasize that, given the same data, one can learn much more by reconstructing the dark energy density directly (and its time dependence) than by attempting to reconstruct its equation of state."
    },
    "0402/astro-ph0402563_arXiv.txt": {
        "abstract": "We present high-spatial resolution Plateau de Bure Interferometer CO(2--1) and SiO(2--1) observations of one intermediate-mass and one high-mass star-forming region. The intermediate-mass region IRAS\\,20293+3952 exhibits four molecular outflows, one being as collimated as the highly collimated jet-like outflows observed in low-mass star formation sources. Furthermore, comparing the data with additional infrared H$_2$ and cm observations we see indications that the nearby ultracompact H{\\sc ii} region triggers a shock wave interacting with the outflow. The high-mass region IRAS\\,19217+1651 exhibits a bipolar outflow as well and the region is dominated by the central driving source. Adding two more sources from the literature, we compare position-velocity diagrams of the intermediate- to high-mass sources with previous studies in the low-mass regime. We find similar kinematic signatures, some sources can be explained by jet-driven outflows whereas other are better constrained by wind-driven models. The data also allow to estimate accretion rates varying from a few times $10^{-5}$\\,M$_{\\odot}$yr$^{-1}$ for the intermediate-mass sources to a few times $10^{-4}$\\,M$_{\\odot}$yr$^{-1}$ for the high-mass source, consistent with models explaining star formation of all masses via accretion processes. ",
        "introduction": "Studies of massive molecular outflows have revealed many important insights in the formation of massive stars over recent years.  Based on the morphologies and energetics we can deduce physical processes taking place at the inner center of the regions. Several single-dish studies agree on the results that massive molecular outflows are ubiquitous phenomena in massive star formation and that they are far more massive and energetic than their low-mass counterparts (e.g., \\citealt{shepherd1996a,ridge2001,zhang2001,beuther2002b}).  A point these studies disagree on is the degree of collimation of massive outflows. Based on early studies by \\citet{shepherd1996b} and its follow-ups, it was believed that high-mass outflows tend to be less collimated than low-mass flows. As the outflow collimation is theoretically tightly connected with the accretion process, these studies favored the idea that massive stars might form via different physical processes, e.g., the coalescence of intermediate-mass protostars at the very center of dense evolving cluster \\citep{bonnell1998,stahler2000,bally2002}.  However, recent observations by \\citet{beuther2002b} show that the previously claimed lower collimation of massive outflows is mostly an observational artifact caused by the larger distances of the target sources (on the average a few kpc) and too low spatial resolution of most studies. Their data taken with the IRAM 30\\,m telescope at a spatial resolution of $11''$ are consistent with massive, bipolar outflows as collimated as their low-mass counterparts. This implies that massive stars can form in a qualitatively similar manner as low-mass stars, just with accretion rates increased by orders of magnitude.  These latter observations are still based on single-dish observations, and to substantiate the scenario a statistically significant number of high-spatial-resolution interferometer studies of massive molecular outflows is necessary. As a first step in that direction \\citet{beuther2002d} have observed the massive star-forming region IRAS\\,05358+3543 with the Plateau de Bure Interferometer (PdBI) in CO(1--0), SiO(2--1) and H$^{13}$CO$^+$(1--0). They observed a massive outflow from the central object of the evolving cluster which is jet-like and highly collimated with a collimation degree of 10. This is the upper end of collimation degrees observed for low-mass outflows as well \\citep{richer2000}. In addition to that collimated jet-like structure, they observed at least two more outflows within the same region.  In another source, IRAS\\,19410+2336, the two outflows observed at single-dish resolution split up at least into 7 separate outflows when observed with interferometers \\citep{beuther2003a}. Similar results were observed toward G35.2 by \\citet{gibb2003}. One of the main conclusions of these studies is that massive star-forming regions can appear confusing with single-dish instruments, but that it is possible to disentangle the structures with high enough spatial resolution into features well known from low-mass star formation.  Contrary, other high-spatial-resolution studies of high-mass star-forming regions indicate that massive outflows can also appear morphologically and energetically different to their low-mass counterparts (e.g., \\citealt{shepherd1998,shepherd2003}). As the high-spatial-resolution results are still based on poor statistical grounds, we pursued massive outflow studies with the PdBI\\footnote{IRAM is supported by INSU/CNRS (France), MPG (Germany), and IGN (Spain).}.  Here we present the results of two more regions~-- IRAS\\,19217+1651 and IRAS\\,20293+3952~-- observed at an angular resolution as high as $1.8''$ in CO(2--1) and SiO(2--1). The two sources are part of a large and well studied sample of 69 high-mass protostellar objects at early evolutionary stages prior to producing significant ultracompact H{\\sc ii} (UCH{\\sc ii}) regions \\citep{sridha,beuther2002a,beuther2002b,beuther2002c}. The two sources were chosen because they combine different features of massive star formation: IRAS\\,19217+1651 has a luminosity of $10^{4.9}$\\,L$_{\\odot}$ and shows a rather simple morphology with one mm continuum source associated with cm emission and H$_2$O and class II CH$_3$OH masers. Contrary, IRAS\\,20293+3952 contains a small UCH{\\sc ii} region, contributing most of the bolometric luminosity ($10^{3.8}$\\,L$_{\\odot}$, \\citealt{sridha}), and likely a cluster of younger intermediate-mass sources triggering the molecular outflows. While the single-dish outflow map of IRAS\\,19217+1651 shows a well-defined bipolar morphology, already the single-dish data of IRAS\\,20293+3952 show that we are dealing with multiple outflows in that region \\citep{beuther2002b}. Both sources cover a wide range of characteristics from intermediate- to high-mass star formation, the main source parameters are listed in Table \\ref{sources}.  After describing the observations in \\S \\ref{obs}, we present the observational results for both sources separately in \\S \\ref{obs_res}. Then we discuss the results in the framework of massive star formation and include literature data with special regard to the position-velocity structure of massive outflows in \\S \\ref{discussion}. Finally, \\S \\ref{conclusion} draws the conclusions, summarizes the current stage of massive molecular outflow studies, and outlines main topics to be tackled in the coming years. ",
        "conclusions": "\\label{conclusion}  The presented analysis of high-spatial-resolution observations of intermediate- to high-mass molecular outflows indicates that the outflow morphologies/kinematics and thus their driving mechanisms do not vary significantly compared to their low-mass counterparts. The higher the source luminosity the more energetic the outflows, but the qualitative signatures are similar. These observations indicate that similar driving mechanisms can be responsible for outflows of all masses.  We find an extremely collimated jet-like outflow emanating from the intermediate-mass source mm1 in IRAS\\,20293+3952. This outflow shows the highest degree of collimation at highest velocities and slightly lower collimation at lower velocities, similar to low-mass sources \\citep{bachiller1996}. The whole region IRAS\\,20293+3952 shows many signs of dynamical interactions, not only between the different outflows (at least four) but there are also indications of a shock wave from the nearby UCH{\\sc ii} region interacting with the collimated outflow. While the UCH{\\sc ii} region is the main source of luminosity in IRAS\\,20293, most of the mechanical force stems from the outflows of the intermediate-mass sources.  The high-mass source IRAS\\,19217+1651 shows a nice bipolar outflow, slightly less collimated than the outflow (A) in IRAS\\,20293, but still comparable to many low-mass flows. This region is strongly dominated by the central core which exhibits mm/cm continuum emission as well as H$_2$O and CH$_3$OH maser emission.  Position-velocity diagrams of molecular outflows from intermediate to high masses show similar signatures as known for low-mass outflows. Some sources are better explained by jet-driven outflows whereas others seem to be due rather to wind-driven outflows. IRAS\\,19217 exhibits signatures of both. The proposal from \\citet{lee2002} that a combination of both driving mechanisms can explain all outflows consistently also holds for our sample.  Estimated accretion rates are of the order a few times $10^{-5}$\\,M$_{\\odot}$yr$^{-1}$ for the intermediate-mass sources in IRAS\\,20293 and a few times $10^{-5}$\\,M$_{\\odot}$yr$^{-1}$ for the high-mass source IRAS\\,19217, consistent with models forming stars of all masses via accretion (e.g, \\citealt{norberg2000,mckee2003}).  The data presented in this paper further support the idea that massive stars form via similar accretion-based processes as their low-mass counterparts. The main difference appears to be their clustered mode of formation and increasing accretion rates and energetics with increasing stellar mass and luminosity. However, investigations of the most massive stars is just beginning, and studies like this one so far rarely exceeded sources with luminosities $>10^5$\\,L$_{\\odot}$. This is to a large degree due to the fact that there simply do not exist many sources with far higher luminosity which are in a state of evolution prior or at the very beginning to form a significant UCH{\\sc ii} region. Therefore, on the one hand we have to extend massive star formation research significantly to even higher luminosities to confirm the present results or to identify possible differences in that regime. On the other hand, the constrains set on the massive star-forming processes are yet mostly indirect, e.g., observing molecular outflows on large scales and inferring the processes likely taking place at the cluster centers. As the spatial resolution and sensitivity of (sub-)mm interferometers increase steadily, it is now necessary to really study the cluster centers and try to resolve the relevant processes in more direct ways. For example, massive disk which are crucial to explain the observed outflows need to be properly identified, resolved and studied to manifest its physical conditions. Furthermore, the strong radiation of the massive protostars significantly changes the chemistry of those central regions. The broad bandwidth and high spatial resolution of current and future (sub-)mm interferometers (SMA, PdBI, CARMA, and further on ALMA) will shed light on many such processes."
    },
    "0402/nucl-th0402057_arXiv.txt": {
        "abstract": " ",
        "introduction": "The main purpose of this article is to show how a mean field treatment of neutron star crust matter can be used  to address the previously unsolved problem of evaluating a quantity, namely the relevant neutron mobility coefficient {\\bb \\calK\\fb}, that is essential for the astrophysical applications that will be described in a separate article~\\cite{CCHII}. In terms of the relevant number density {\\bb \\nn\\fb} of effectively unbound neutrons, this coefficient determines a corresponding effective mass {\\bb \\mm_\\star=\\nn/\\calK\\fb} that characterises their average motion on a macroscopic scale (meaning one that is large compared with the spacing between nuclei) and that we therefore refer to as the macro mass, to distinguish it from the microscopic effective mass, {\\bb\\microm\\fb} say, characterising the dynamics of the neutrons on subnuclear scales. Whereas {\\bb\\microm\\fb} is well known to be typically rather smaller than the ordinary neutron mass {\\bb\\mm\\fb}, we reach the previously unexpected conclusion that there is likely to be a strong ``entrainment'' effect whereby the macro mass {\\bb \\mm_\\star\\fb} will typically become large, and in some layers extremely large, compared with {\\bb\\mm\\fb}.   A secondary purpose of this article is to draw the attention of nuclear theorists to the potentialities of the almost entirely unexplored branch of theoretical astrophysical nuclear physics that needs to be developed for this and many other purposes. The only relevant work of which we are aware so far is that of Oyamatsu and Yamada ~\\cite{Oyamatsu94}, who appear to be the only ones to have taken proper account of the neutron scattering by the nuclei inside the inner crust by the use of appropriate Bloch type periodicity conditions of the kind commonly employed for the treatment of electrons in ordinary terrestrial solid state physics. Their treatment however was restricted to a simple one dimensional model.  Under the conditions of ordinary terrestrial solid state physics, and even at the much higher densities characterising the matter in a white dwarf star, long range electric forces keep the nuclei so far apart that, in so far as the much stronger but short range nuclear interactions are concerned, each individual nucleus can be treated separately as if it were isolated. Until quite recently~\\cite{Oyamatsu94}, such a separate treatment of individual nuclei (considered as if isolated each in its own cell - with Wigner Seitz type boundary conditions) has been used in nearly all quantum mechanical calculations on neutron star crust matter since the pioneer work of Negele and Vautherin~\\cite{Vautherin73}. That kind of approximation is fully justifiable in the outer crust , where the densities are not too much greater than those found in a white dwarf. However such a treatment can no longer be considered entirely satisfactory in the inner crust, meaning the part with density above the ``neutron drip'' threshold at about 10$^{11}$ g/cm$^3$, where there are unconfined neutrons that travel between neighbouring nuclei, which thereby cease to be effectively isolated from one another.  While desirable for accuracy throughout the inner crust, a proper collective rather than individual treatment of the nuclei becomes not just desirable but absolutely essential for treating the problem with which Oyamatsu and Yamada ~\\cite{Oyamatsu94}~\\cite{Oyamatsu93} were concerned, namely that of the nuclear matter inside neutron star crust. Such a treatment is also essential for treating the problem with which the present work is concerned, namely that of stationary but non static configurations in which a neutron current flows relative to the lattice formed by the nuclei, something that obviously can not be discussed in the usual approach that treats the nuclei as if they were isolated in individual (e.g. Wigner Seitz type) boxes.  The flow of neutrons is treated here as a perturbation of a zero temperature ground state characterised just by the location of the relevant Fermi surface in momentum space. We thereby obtain provisional rough estimates of the relevant mobility coefficient which suggest that (unlike what occurs in the fluid core and on a microscopic scale) the macro mass {\\bb\\mm_\\star\\fb} that effectively characterises the neutron motion on a macroscopic scale can become very large compared with the ordinary neutron mass {\\bb\\mm\\fb}, particularly in the middle part of the conducting layer, for which three dimensional numerical results will be presented in a follow up article~\\cite{Chamel04}. The present article deals more specifically with simplified rod and plate type models that are relevant near the crust core interface, where the mass enhancement will be less extreme.  Bulgac, Magierski and Heenen have recently pointed out the importance of shell effects induced by unbound neutrons in neutron star crust by evaluating the Casimir energy for neutron matter in the presence of inhomogeneities from a semiclassical approach and more recently by performing a Skyrme Hartree-Fock calculation with ordinary periodic boundary conditions (see \\cite{BulMagHeen02, MagHeen02} and references therein). However this kind of boundary conditions does not properly account for Bragg scattering of dripped neutrons and is only a particular case of the more general Bloch type boundary conditions.  In the absence of any previous quantum mechanical calculation whereby nuclei on a crystal lattice are treated collectively (apart from the 1D calculation previously mentionned \\cite{Oyamatsu94}), even at the simplest level of approximation, we shall adopt the simple model suggested by Oyamatsu and Yamada, supplemented with Bloch type boundary conditions, in order to estimate the effective neutron mass {\\bb\\mm_\\star\\, .\\fb} This model treats the neutrons as independent fermions subject to an effective potential. We wish to draw the attention of nuclear theorists to the problem of including  Bloch boundary conditions in more sophisticated approximation schemes as a challenge for future work.   In the mean time, experience with the analogous problem of electron transport in ordinary solid state physics suggests that the results obtained from the Oyamatsu Yamada type treatment used here should not be too bad as a first approximation. Further encouragement comes from our own recent attempt to take up the challenge of allowing for coupling by an appropriate adaptation of the standard BCS pairing theory on which the prediction of neutron superfluidity (in the relevant low to moderate temperature range) is based: the upshot~\\cite{CCHIII} is that (although it is essential for the inhibition of resistivity) as far as the ``entrainment'' phenomenon is concerned the effect of the ensuing pairing ``gap'' will not be very large nor very difficult to calculate. ",
        "conclusions": "The scattering of dripped neutrons by the nuclei in the inner crust leads on a macroscopic scale to a modification of the neutron mass {\\bb  \\mm_ \\star \\fb}, which can be expressed via a well defined mobility scalar {\\bb  \\calK \\fb} by {\\bb  \\mm_\\star =\\nn/\\calK \\fb} in which {\\bb  \\nn \\fb} is the (arbitrary) density of such unbound neutrons, and {\\bb  \\calK \\fb} is found to be expressible as an integral of the group velocity {\\bb  \\vv^i \\fb} over the corresponding Fermi surface. This effective macro mass should not be confused with the effective micro mass, relevant for subnuclear scales, which is usually found to be smaller than the ordinary neutron mass.  Bragg scattering of dripped neutrons is taken into account here by applying Bloch type boundary conditions, which are well known in solid state physics but have been barely used in this nuclear context. We have computed numerical values of this mobility scalar in the bottom layers of the inner crust near the crust-core interface, for simple models in the ``pasta'' layers: equally spaced slab shaped nuclei (``lasagna'') and rod like nuclei on either a square or an hexagonal 2D lattice (``spaghetti''). The (anisotropic) entrainment effect is small at such densities since the system is nearly homogeneous. It appears that the mobility scalar tends to be systematically reduced compared to the homogeneous expression and the farther from homogeneity, the smaller is the mobility scalar. The resulting effective mass {\\bb  \\mm_\\star \\fb} is found to be larger than the bare neutron mass. This results can be interpreted as a macroscopic manifestation of the modifications in the shape of the Fermi surface. Notably as one goes from the homogeneous outer core to the crust, the spherical neutron Fermi surface gets distorted and even torn and pierced.    The main reason is that the neutron Bragg scattering by crustal nuclei leads to the opening of energy band gaps. A specific feature of those exotic phases is that the mobility scalar is bounded, {\\bb \\calK \\geq 2\\nn_{\\rmn} / 3 \\mm \\fb} for the ``lasagna'' phase and {\\bb \\calK \\geq \\nn_{\\rmn} / 3 \\mm \\fb} for the ``spaghetti'' phase due to the fact that the neutrons are still free to move in one or two dimensions, respectively. There is no such bound for three dimensional crystals, for which  smaller mobility scalars (larger effective masses) may be expected.  From these considerations, we can infer that the mobility scalar {\\bb  \\calK \\fb} will increase with the density, starting from zero in the outer crust below the neutron drip threshold where all neutrons are confined, since {\\bb  \\vv^i=0 \\fb} for all states,  to its largest possible value in the homogeneous neutron star mantle. At low densities near neutron drip, where the conduction neutron density is negligible compared to the total neutron density, {\\bb \\nn \\ll \\nn_{\\rmn}\\, ,\\fb}  the crystal potential {\\bb  \\VV\\fb} is quite large but the fraction of the cell occupied by the nuclei is small so that dripped neutrons propagate essentially freely. Consequently we may expect to get {\\bb  \\calK \\simeq \\nn/\\mm\\fb} and {\\bb  \\mm_\\star\\simeq \\mm\\fb} (on the understanding that, as discussed in Subsection \\ref{cutoff}, a conduction state is defined so as to have an associated group velocity that differs significantly from zero). On the other hand,  at very high densities where nuclei nearly merge, the potential is much weaker and smoothly varying and we shall have {\\bb \\nn\\simeq \\nn_{\\rmn}\\, ,\\fb} so we expect to find {\\bb  \\mm_\\star\\simeq \\mm\\, .\\fb}  In other words the mean effective mass {\\bb  \\mm_\\star\\fb} is expected to be close to the ordinary mass {\\bb  \\mm\\fb} at the top (above neutron drip density where neutrons start to leak out of nuclei) and bottom (where nuclei merge into a uniform mixture of neutrons, protons and electrons at density about $1/3$ or $2/3$ the nuclear saturation density {\\bb  \\nn_{\\rm sat}\\simeq 0.16\\, {\\rm fm}^{-3}\\fb} \\cite{Haensel01}) of the inner crust, but may reach a much larger value (with astrophysically interesting consequences) in the intermediate layers. This indicates that the unbound neutron band effects (that are effectively neglected in the  commonly used W-S approximation) could have important consequences as concerns the equilibrium neutron star crust structure and composition.  The exploratory character of the present work, justifies the simplicity of the single particle model we used. In a more realistic, Hartree-Fock calculation, the single particle potential and effective micro mass should have to be determined self-consistently. The potential may contain momentum dependent terms, as with effective nucleon-nucleon interactions of the Skyrme type, which leads to space varying effective neutron masses {\\bb \\mm^{_\\oplus}\\{ {\\bf r} \\} \\fb} which are typically smaller than the bare neutron mass. However the general arguments we developed, relying on the shape of the Fermi surface, suggest that the qualitative results of our calculations, namely the enhancement of the macroscopic effective mass {\\bb \\mm_\\star \\fb} will remain valid in more elaborate single particle schemes. Finally let us mention that as we have shown in a recent paper \\cite{CCHIII}, the neutron pairing is not expected to qualitatively alter our present conclusions.  \\vskip 0.5cm {\\bf Acknowledgements}  One of the authors (PH) was partly supported by the KBN grant no. 1-P03D-008-27 and by the LEA Astro-PF PAN/CNRS program. We wish to thank the referee for questions that have helped us to improve the clarity of our presentation.  \\vfill\\eject {\\bf APPENDIX}   \\appendix"
    },
    "0402/astro-ph0402094_arXiv.txt": {
        "abstract": "{Several solar analogs have been identified in the library of high resolution stellar spectra taken with the echelle spectrograph ELODIE. A purely differential method has been used, based on the $\\chi^2$ comparison of a large number of G dwarf spectra to 8 spectra of the Sun, taken on the Moon and Ceres. HD 146233 keeps its status of closest ever solar twin (Porto de Mello \\& da Silva, \\cite{PMDS97}). Some other spectroscopic analogs have never been studied before, while the two planet-host stars HD095128 and HD186427 are also part of the selection. The fundamental parameters found in the literature for these stars show a surprising dispersion, partly due to the uncertainties which affect them. We discuss the advantages and drawbacks of photometric and spectroscopic methods to search for solar analogs and conclude that they have to be used jointly to find real solar twins. ",
        "introduction": "The Sun is the best-known star : its fundamental parameters (radius, mass, age, luminosity, effective temperature, chemical composition) are known with a good accuracy, as well as its internal structure, activity, velocity field and magnetic field. Consequently the Sun is used as the fundamental standard in many astronomical calibrations. One of the motivations to identify stars that replicate the solar astrophysical properties is the necessity to have other reference stars, observable during the night under the same conditions as any other target. The pioneers of the subject (Hardorp \\cite{H78}, Cayrel de Strobel et al \\cite{cay81}) were also involved in resolving the problem of the photometric indexes of the Sun, inherent to the impossibility to observe it as a point-like source. In the last decade the motivation of finding such stars has been increased by an exciting challenge : the search for planetary systems that could harbour life. Solar analogs are straightforward targets for this hunt.  The first searches of solar analogs were performed by photometric and spectrophotometric techniques. Hardorp (\\cite{H78}) compared UV spectral energy distributions of nearly 80 G dwarfs to that of the Sun and found 4 stars that had a UV spectrum indistinguishable from solar : HD028099 (Hy VB 64), HD044594, HD186427 (16 Cyg B), HD191854. Neckel (\\cite{neck86}) established a list of bright stars with UBV-colours close to those of the Sun and confirmed the photometric resemblance of Hy VB 64 and 16 Cyg B to the Sun. With the advance of techniques in high resolution spectroscopy and solid state detectors, and with the progress in modelling stellar atmospheres, measurements of (T$_{\\rm eff}$, log\\,g, [Fe/H]) became of higher precision allowing the search for solar analogs by comparing their atmospheric parameters to those of the Sun. G. Cayrel de Strobel made a huge contribution to the subject with the detailed analysis of many candidates (Cayrel de Strobel et al \\cite{cay81}, Cayrel de Strobel \\& Bentolila \\cite{cay89}, Friel et al \\cite{friel93}) and a review of the status of the art (Cayrel de Strobel \\cite{cay96}). She also introduced the concepts of solar twin, solar analog and solar-like star. Porto de Mello \\& da Silva (\\cite{PMDS97}) presented the star HD146233 (18 Sco) with physical properties extremely close to those of the Sun, as the \"closest ever solar twin\".  A workshop on Solar Analogs was held  in 1997 at the Lowell Observatory to provide a solid basis to the hunt of solar analogs. After many discussions on the performances of different methods, a list of the best candidates was established, in which 4 stars received the agreement of almost all participants : HD217014 (51 Peg), HD146233 (18 Sco), HD186408 (16 Cyg A), HD186427 (16 Cyg B).  In this paper, we take advantage of a large and homogeneous dataset of high resolution echelle spectra which are compared directly to solar spectra, independently of any model or photometric measurements. The eye is replaced by a more reliable criterion, approximatively the reduced $\\chi^2$, computed over $\\sim$ 32000 resolution elements. This purely differential method allowed us to identify several stars whose optical spectrum looks closely like the Sun's, the best one being HD146233.  We describe in Sect. \\ref{s:ELO} our observational material and differential method, and we give the list of our Top Ten solar analogs. We have searched the literature for their colour indexes and atmospheric parameters and calculated absolute magnitudes from Hipparcos parallaxes. We discuss the uncertainties which affect these data and compare them to that of the Sun (Sect. \\ref{s:param}). In Sect. \\ref{s:Li}, we examine qualitatively their Li content and give information on their activity and age. In Sect. \\ref{s:other} we discuss several stars, having similar colours and absolute magnitude or similar atmospheric parameters as the Sun but slightly different spectra. ",
        "conclusions": "We have presented the 10 stars of the ELODIE library which exhibit the closest optical spectrum to the Sun at a resolution of 42000. They have colours, absolute magnitudes, atmospheric parameters and Li content which span a range of values larger than expected. It is surprising that a star like HD089269, colder, more metal poor and less luminous than the Sun is at the 4th position, whereas HD076151, having similar colours, absolute magnitude and atmospheric parameters is only at the 14th position. Activity may play an important role in discrimination. We have shown for instance that the good photometric analog HD001835 was a very bad spectroscopic analog because of its high activity. One also has to take into account, when comparing colours, absolute magnitudes and atmospheric parameters to those of the Sun, that these quantities are affected by significant uncertainties. Effective temperatures are particularly in question, with determinations for the same star differing by nearly 200K in some cases.  Our method consisting in measuring distances between spectra is powerful but it is also affected by uncertainties due to observing conditions, especially the pollution by telluric lines, which may perturb the order of the classification.  Among our Top Ten, several stars have never been mentioned before as solar analogs and have been very little studied. They are good candidates for planet hunting, especially HD047309 which is slightly more metal rich than the Sun. Two of our solar analogs, HD095128 and HD186427, are already known to have planets. HD159222 and HD076151 are also good candidates because they are good spectroscopic analogs (in the Top 15) and good photometric analogs.  The conclusion of this work is that none of the methods to search for solar twins is satisfactory when used by itself. The methods that have been already used are the comparison of colour indexes, of absolute magnitudes, of UV spectral energy distributions, of atmospheric parameters and of high resolution optical spectra. All these methods are affected by uncertainties and none of them is able to describe sufficiently all the stellar properties. Combining them is the best way to minimize their drawbacks, uncertainties and insufficiencies. Finally HD146233 is the only star in the ELODIE library which merits the title of solar twin because it has passed the filter of all methods. It is not however a perfect twin and differs from the Sun by its higher Li content, slightly higher age (6 Gyr against 4.6 Gyr for the Sun) and higher luminosity ($M_{\\rm V}=4.77$ against $M_{\\rm V \\odot}=4.82$)."
    },
    "0402/astro-ph0402607_arXiv.txt": {
        "abstract": "The lightcurve of the Large Magellanic Cloud (LMC) variable star MACHO 81.8997.87 shows evidence for photometric variations due to both stellar pulsation, with a 2.035 day period, and eclipsing behavior, with an 800.4 day period.  The primary star of the system has been identified as a first-overtone Cepheid but the nature of the secondary star has not been determined.  Here we present multicolor BVI photometry of a primary eclipse of the system and fit a model to the complete lightcurve to produce an updated set of elements.  These results are combined with 2MASS JHK photometry to give further insight into the identity of the companion star.  We find that the companion is most consistent with a late-K or an early-M giant but also that there are a number of problems with this interpretation.  The prospects for future observations of this system are also discussed. ",
        "introduction": "There is a small but growing list of regularly pulsating stars that are known to be members of eclipsing binary systems.  \\cite{myfirst} present observations and analysis of three eclipsing Cepheid variables for which data exist in the MACHO Project Large Magellanic Cloud (LMC) database.  \\cite{RB01} list nine eclipsing binary systems that contain $\\delta$ Scuti variables and new candidates have been identified since that time (see \\cite{gkdra} and \\cite{ecl_d}).  Most recently \\cite{eclRR} list three objects whose lightcurves show evidence for eclipses and RR Lyr-type pulsations.  Although one or more of these may be the result of photometric contamination, clearly this is a burgeoning field for obtaining long-sought direct measurements of pulsating star properties.  The astrophysical returns from systems that combine eclipsing and pulsating behavior can be considerable.  An eclipsing Cepheid system, if also a double-lined spectroscopic binary, can give a determination of the mass and luminosity of the Cepheid that is not only more accurate than existing measurements but also independent of assumed distance estimates.  Such a system would offer an independent calibration of the period-luminosity and period-luminosity-color relations and the most direct measurement of the Cepheid's mass.  Here we present additional observations and an updated analysis of the eclipsing Cepheid system MACHO 81.8997.87.  In particular, we more strongly constrain the nature of the system's secondary star. ",
        "conclusions": " \\begin{enumerate} \\item{Based on the updated set of optical magnitudes, colors and relative radii we can classify the components.  They are most consistent with an intermediate-mass overtone Cepheid with a late K or M-type giant companion.} \\item{This result is inconsistent with the expectations from evolutionary theory.  The companion is too cool and dim for the system to match theoretical isochrones.} \\item{In the near-infrared, a companion with cooler colors than standard giant stars is needed to replicate the observed system color.} \\end{enumerate}  Clearly, more observations are needed to fully realize the considerable potential of this system.  In particular one of the principal sources of the uncertainty in the companion's properties is the lack of observations of a secondary eclipse.  To facilitate follow-up work Table \\ref{future-tab} presents a table of predicted future dates of primary and secondary eclipses.  Given the low temperature of the companion, observations taken at near-infrared wavelengths should put a stronger constraint on the companion's properties.  In particular, precise photometry taken during the primary and secondary eclipses would allow better estimates of the individual colors of each component so their location in Figure \\ref{jhk-fig1} would be better determined.  Observations taken on an 8m class telescope would have sufficient resolution to identify possible sources of contamination within the crowded field.  The companion and Cepheid would appear to have similar fluxes between $J~(1.22~\\mu m)$ and $H~(1.63~\\mu m)$ and therefore (given the estimated radial velocities given above) radial velocity work should be attempted with a high-resolution near-infrared spectrograph on an 8m-class telescope.  Near-infrared spectra could also provide a more definitive classification of the companion star."
    },
    "0402/astro-ph0402431_arXiv.txt": {
        "abstract": "{{As Be stars are restricted to luminosity classes III-V, but early B-type stars are believed to evolve into supergiants, it is to be expected that the Be phenomenon disappears at some point in the evolution of a moderately massive star, before it reaches the supergiant phase. As a first stage in an attempt to determine the physical reasons of this cessation, a search of the literature has provided a number of candidates to be Be stars with luminosity classes Ib or II. Spectroscopy has been obtained for candidates in a number of open clusters and associations, as well as several other bright stars in those clusters. Among the objects observed, HD~207329 is the best candidate to be a high-luminosity Be star, as it appears like a fast-rotating supergiant with double-peaked emission lines. The lines of HD~229059, in Berkeley~87, also appear morphologically similar to those of Be stars, but there are reasons to suspect that this object is an interacting binary. At slightly lower luminosities, LS~I~$+56\\degr$92 (B4\\,II) and HD~333452 (O9\\,II), also appear as intrinsically luminous Be stars. Two Be stars in NGC~6913, HD~229221 and HD~229239, appear to have rather higher intrinsic magnitudes than their spectral type (B0.2\\,III in both cases) would indicate, being as luminous as luminosity class II objects in the same cluster. HD~344863, in NGC~6823, is also a rather early Be star of moderately high luminosity. The search shows that, though high-luminosity Be stars do exist, they are scarce and, perhaps surprisingly, tend to have early spectral types.} ",
        "introduction": "\\label{sec:intro}  In recent years, it has become obvious that rotation plays a fundamental role in the evolution of massive stars (Maeder \\& Meynet 2000). The initial rotational velocity determines the main-sequence lifetime of a star, its post-main-sequence evolution and likely even the nature of its final remnant after supernova explosion (Meynet \\& Maeder 2000, 2003). However, one of the most obvious manifestations of stellar rotation, the Be phenomenon, remains unexplained after more than a century of study (see Porter \\& Rivinius 2003 for a recent review).  The accepted definition of a Be star is that given by Collins (1987), namely, ``a non-supergiant B star whose spectrum has, or had at some time, one or more Balmer lines in emission''. This definition may need some qualification, as the Be phenomenon appears to extend into late O and early A subtypes, and also because some objects (such as Herbig AeBe stars) are generally excluded from the class (see Porter \\& Rivinius 2003 for a detailed discussion). It is, however, widely accepted for the ``classical Be stars'', i.e., fast rotating moderately massive stars in which the emission lines are produced in a circumstellar disk of material expelled from the photosphere.  The ``non-supergiant'' part of the definition is generally understood to represent a warning against potential confusion. Many B-type supergiants show relatively strong H$\\alpha$ emission lines, which are not produced in a circumstellar disk. These lines arise in the strong radiatively-driven winds of the supergiants (e.g., Leitherer 1988) and are considered morphologically normal (i.e., not unusual for the spectral type). As a consequence of their different formation mechanisms, H$\\alpha$ emission lines in classical Be stars and supergiants usually have rather different shapes, with the lines in Be stars showing rotational symmetry, while the lines in supergiants typically display P-Cygni profiles (see Fig.~\\ref{fig:superalpha}).  There may, however, be a deeper interpretation to this qualification, as it clearly states that no supergiant star displays the Be phenomenon. As an important fraction of O9-B1 main sequence and giant stars display the Be phenomenon (e.g., Zorec \\& Briot 1997) and these stars are sufficiently massive to become blue supergiants at a later stage, one would (naively) conclude that the Be phenomenon has to {\\it disappear} at some point during the post-Main-Sequence (post-MS) evolution of the star.  There are two immediately obvious physical effects to which this cessation of the Be phenomenon could be linked. On the one hand, the Be phenomenon is known to be related to fast rotation and supergiants do not rotate fast. If the unknown cause of mass loss requires fast rotation, then mass loss could stop as the rotation slows. Because of angular momentum conservation, a star is expected to slow down as it expands (cf. Steele 1999). On the other hand, as the luminosity increases, the star develops a slow radiative wind which could exert a force on the circumstellar disk that would lead to its dissipation. The disk would be swept away by radiation pressure even if suitable conditions for its formation were still present.  In reality, things could not be so simple. The relevant factor in the Be phenomenon is not high rotational velocity itself but its ratio to some critical velocity, which measures the tendency of the star to be centrifugally disrupted (e.g., Porter 1996; Keller et al. 1999; 2001). More complicated evolutionary models can therefore be envisaged in which loss of angular momentum via wind breaking stops the Be phenomenon (cf. Meynet \\& Maeder 2000; Keller et al. 2001).   In any case, evidence from young open clusters in the Milky Way and the SMC (e.g., Keller et al. 1999; 2001) supports the idea that fast rotating B stars become Be stars late in their MS life, as the highest fraction of Be stars appears to occur at the MS turn-off and above it. Therefore, whatever the mechanisms involved, if the Be phenomenon occurs preferentially in stars that are close to the end of their MS lives or have already started to move away from the MS and, at the same time, we know that it has already stopped when the star reaches the supergiant stage, it is tempting to conclude that the Be phenomenon is related to some physical conditions occurring only during a fraction of the lifetime of the star. If this is so, some information on the mechanism producing the Be phenomenon can be gained by determining at what phase in the evolution of the star the Be characteristics disappear and what are the physical parameters of the star when this happens.  In this spirit, a programme has been started to identify the highest luminosity at which a star may display the Be phenomenon. If one can identify stars of relatively high luminosity which still display the Be phenomenon, important constraints could be gained on the physical conditions that lead to its development. Some care is necessary here, as Be characteristics are mainly observational, and the assumption that they obey to a single physical mechanism may be misleading. Therefore a careful determination of the stellar parameters of putative high luminosity Be stars via fitting of stellar atmosphere models will be necessary before any conclusion can be reached.  \\begin{figure}[ht] \\resizebox{\\hsize}{!} {\\includegraphics[bb=60 165 520 600]{an789f1.ps}} \\caption{H$\\alpha$ emission lines in four bright early-B supergiants. The P-Cygni profiles may be considered typical of this class of objects. Note that the absorption trough tends to disappear towards early spectral types (it is very weak already in HD~14143), making the profile look like a single peak, specially for stars B0 and earlier. Note also the increase in the strength of the \\ion{C}{ii}~$\\lambda\\lambda$6578, 6583\\AA\\ doublet towards a maximum at B3-4. Compare these profiles with those in Fig.~\\ref{fig:minialpha}.} \\label{fig:superalpha} \\end{figure}   As a first step for such work, in this paper I set out to investigate the existence of high-luminosity Be stars. As Be stars are known to be widespread among luminosity classes III to V, attention is centred on OB stars of luminosity class I or II that may display Be characteristics. ",
        "conclusions": "\\subsection{Consistency check}  Before attempting to interpret results, it may be sound to check whether the spectral types derived provide a consistent way of comparing stars in clusters. For this, I have calculated the corresponding distance moduli for the stars under study using the intrinsic colours of Wegner (1994) and the spectral type to intrinsic magnitude calibration used in Negueruela \\& Marco (2003), conveniently interpolated when necessary. The distances obtained are listed in Table~\\ref{tab:consistent}.   \\begin{table*}[ht] \\caption{Distance moduli calculated using the spectral types derived here and photometry from the literature. The primary reference is Hiltner (1956), who has been found to be in good accord with other references when more than one measurement exists. For stars not observed by Hiltner photometry has been taken from Hoag et al. (1965, BD $+35\\degr$3955) and Coyne et al. (1975, HD~227733).   } \\label{tab:consistent} \\begin{tabular}{llccccc}\\hline Assoc. & Object & Spectral & $V$ & $(B-V)$ & $E(B-V)$& DM\\\\ & & Type &&&&\\\\ \\hline Vul OB1 & HD 344863& O9\\,III& 8.83 & 0.68 & 0.94 &11.4\\\\ &HD 344873 & O9.7\\,II & 8.77 & 0.77 &1.00 & 11.4\\\\ \\hline Cyg OB3& HD 227634& B0.2\\,II&7.91&0.25&0.46 &12.1\\\\ &HD 227733& B1.5\\,V(e)&10.31 & 0.24&0.46 & 11.7\\\\ &BD $+35\\degr$3955&B0.7\\,Iab&7.38 &0.25& 0.45 &12.3\\\\ &BD $+35\\degr$3956&B0.5\\,Ve& 8.85 & 0.19&0.43 &11.3\\\\ \\hline Cyg OB1 &HD 229221 & B0.2\\,IIIe& 9.21 & 0.92&1.14 & 10.8\\\\ &HD 229227 & O9.7\\,III & 9.38 & 0.80 & 1.04& 11.6 \\\\ &HD 229234 & O9\\,II& 8.92 & 0.77 &1.04 & 11.6 \\\\ &HD 229238 & B0.2\\,II & 8.88 & 0.90 &1.11 & 11.2\\\\ &HD 229239 & B0.2\\,III & 8.92 & 0.87& 1.10& 10.6 \\\\ &HD 194280 & OC9.7\\,Iab& 8.39 & 0.76&  0.99 & 11.6\\\\ &HD 194334 & O7\\,II(f)& 8.77 & 0.84& 1.13 &11.2\\\\ \\hline \\end{tabular} \\end{table*}  \\subsubsection{NGC~6871}  The four stars observed have basically identical $E(B-V)$, fully compatible with the average $E(B-V)=0.46$ found by previous workers. Their spectroscopic distances show a large spread, but it would fall within what is expected from the method if  BD $+35\\degr$3956 is indeed a binary (for two stars of similar spectral type, its $DM$ would be $\\approx12$). Obviously, none of the Be stars in this cluster are of high luminosity.   In a recent spectroscopic survey, Balog \\& Kenyon (2002) find 6 Be stars in NGC~6871. All of them appear close to the main sequence. It appears that no B-type star has moved off the main sequence yet, which together with the presence of the O-type stars HD~190918 and HD~190864, favours an age $\\la 10$ Myr for NGC~6871. This is outside the age range generally quoted for ``Be-rich'' clusters, though the Be fraction appears to be rather high.  \\subsubsection{Cyg OB1}  The distance moduli have been calculated assuming that $R=3.1$, though there are strong reasons to believe that the reddening to NGC~6913 does not follow the standard law (Crawford et al. 1977; Wang \\& Hu 2000). It is hoped that they can still be used for internal comparison.  Surprisingly, all the luminosity I and II objects give very consistent distances, but the two objects proposed as Be stars, which have luminosity class III give rather shorter distances. In the case of HD~229221, this is partly due to the overestimate of the interstellar extinction, as the observed $E(B-V)$ must have a component of circumstellar origin, like in all Be stars. If we rather assume that all the stars are at the same distance, we find that HD~229221 and HD~229239 are (almost) as luminous as cluster members with luminosity class II, while HD~229227, a fast rotator but not a Be star, is rather fainter. We have to conclude that the two confirmed Be stars, HD~229221 and HD~229239 are rather luminous stars.    For Berkeley~87, no attempt has been made at deriving distance moduli, as none of the stars observed appears to be a single non-emission-line star.  \\subsection{Variability}  Variability is a well known characteristic of the Be phenomenon. We find that several stars that have been classified as emission-line stars in the past do not display any sign of emission now. In principle, we will ignore the possibility of a wrong identification and rather cavalierly assume that the lack of emission lines in modern spectra must be taken precisely as demonstration of variability, identifying our targets as Be star (rather than other types of emission line stars).  The only exceptions to this boldness would be HD~194280 and HD~194334. The former is well observed, has a very high luminosity and a single reference to emission lines. The latter does not look a strong case either. In the case of BD $+36\\degr$4032, we may assume that the presence of double lines may have induced confusion with emission components. Moreover, there appears to be no reliable source identifying the object as an emission-line star.  Among the initial candidates, the three objects in NGC~6871 and V439~Cyg have been found to be indeed Be stars, but close to the main sequence, rather than luminous stars. We are left with 4 candidates that displayed emission lines at the time of the observations (HD~229221, HD~207329, LS~I~$+56\\degr$92 and HD~229059, which may not be a Be star after all), two other candidates that did not display emission lines, but have been reported as Be stars by several {\\em independent} sources (HD~344863 and HD~229239) and one possible Be star which has only been reported once (HD~344873).  Another high-luminosity Be star reported in the literature is HD~333452. This object was classified B0\\,III?np by Morgan et al. (1955), the use of both an uncertainty marker and the ``p'' tag clearly indicating some peculiarity. The Luminous Star catalogue classifies it as OB,ce,le,r. However, HD~333452 was observed by Steele et al.~(1998) in 1998 and appeared as a normal absorption-line O9\\,II star.  Another potential high-luminosity Be star is \\object{BD $+56\\degr$511}, in \\object{$h$ Per}. This object has spectral type B1\\,III (Steele et a. 1999) and is widely reported to be a mild Be star (cf. Vrancken et al. 2000). Vrancken et al. (2000) find an effective gravity of only $\\log g = 3.1\\pm0.1$, based on model atmosphere analysis, which would suggest a rather high luminosity. However, using the generally accepted distance modulus to $h$ Per $(m-M)=11.7$ (e.g., Marco \\& Bernabeu 2001; Keller et al. 2001) and $V=9.11$, $(B-V)=0.38$ from Keller et al. (2001), the implied absolute magnitude (for standard reddening) is $M_{V}=-4.4$, which does not look particularly high for B1\\,III.  \\subsection{Outlook} \\label{sec:outlook}  As outlined in Section~\\ref{sec:intro}, one of the drivers to undertake this study was the suspicion that stars that appear as B0-1 Be stars could later appear as mid-to-late B-type stars of luminosity classes I or II, while still keeping their emission lines. It may therefore appear surprising that all the high-luminosity Be stars found, with the relative exception of LS~I~$+56\\degr$92, are {\\em early} B-type objects.  This might be (at least, in part) due to a selection effect. The criteria for telling luminosity class II from III are relatively straightforward for B0 stars, but become more subtle for mid-B stars and may be rather difficult to apply for emission line objects. It might then be that some objects classified as Be giants have slightly higher luminosities but have not been recognised as such.  The interesting point, however, is the very existence of these early-type luminous Be stars. Objects like HD~344863 or HD~333452 must have started their lives as rather massive stars ($M\\ga 25M_{\\sun}$), in a region of the HR diagram (spectral types O8 or earlier) where Be stars are extremely scarce. As a matter of fact, only one Be star with spectral type earlier than O8 is well documented in the literature, HD~155806 (O7.5\\,IIIe, Conti \\& Leep 1974), and even for this one observations suggest some sort of ``peculiarity'' with respect to other Be stars (Hanuschik et al. 1996). Understanding these very few mid-range O-type stars displaying emission lines that morphologically resemble Be stars is very important in order to constrain the physical causes leading to the Be phenomenon. Based on the few examples known, one might suspect that these objects develop emission lines only late in their lives, but the crucial point would be to determine whether they actually bear the same phenomenology as classical Be stars. Therefore this parameter range will be studied in further work.  Meanwhile, the present work shows that there are at least a few stars of high luminosity presenting emission lines that morphologically will classify them as Be stars. More detailed analysis of their astrophysical parameters is required in order to determine whether the resemblance implies a common physical mechanism. It would appear, however, rather tempting to conclude that a B4II emission-line star displaying the same morphology as a B4IIIe star is also a classical Be star.   The number of these luminous Be stars, however, seems to be rather small, as few convincing candidates were found in the literature and a substantial fraction do not qualify as such when looked at in detail. The existence of an upper limit in the luminosity that a Be star can have appears certain, and the implication that the Be phenomenon must cease at some point in the life of a Be star has a bearing on our quest to understand this phenomenon."
    },
    "0402/astro-ph0402357_arXiv.txt": {
        "abstract": "{ We present benchmark problems and solutions for the continuum radiative transfer (RT) in a 2D disk configuration. The reliability of three Monte-Carlo and two grid-based codes is tested by comparing their results for a set of well-defined cases which differ for optical depth and viewing angle. For all the configurations, the overall shape of the resulting temperature and spectral energy distribution is well reproduced. The solutions we provide can be used for the verification of other RT codes. We also point out the advantages and disadvantages of the various numerical techniques applied to solve the RT problem. ",
        "introduction": "Observations show that many astrophysical sources such as young stellar objects (YSOs), post-AGB stars and active galactic nuclei (AGN) are surrounded by dust. Dust grains scatter, absorb and re-emit radiation originating from the primary energy sources, thus modifying their spectral energy distributions (SEDs). Moreover many embedded objects cannot be directly studied in the visible, since dust may entirely obscure them at optical wavelengths. Their structure can be only inferred from the thermal dust emission. Therefore, modelling of their intensity and polarization maps as well as their SEDs is necessary. This can only be done by solving the radiative transfer (RT) equation (e.g. Yorke \\cite{yorke85}). Analytical solutions for this equation do exist only for the simplest cases, far from representing the complexity of dust-enshrouded objects. Hence, the development of sophisticated numerical RT codes is unavoidable.  Early attempts for spherically symmetric configurations were performed by Hummer \\& Rybicki (\\cite{hummer71}), Scoville \\& Kwan (\\cite{scoville76}) and Leung (\\cite{leung76}) including rough assumptions such as grey opacity and/or neglecting scattering. The first formal solution for the dust continuum in spherical geometry was obtained by Rowan-Robinson (\\cite{rowan80}), directly integrating the RT equation, an operation known as ray-tracing. Since then many other codes treating 1-D slab or spherical configurations (e.g. Yorke \\& Shustov \\cite{yorke81}; Lef\\`evre et al. \\cite{lefevre82}; Martin et al. \\cite{martin84}; Rogers \\& Martin \\cite{rogers84}; Henning \\cite{henning85}; Groenewegen \\cite{groenewegen93}; Winters et al. \\cite{winters94}) or the inverse RT problem (Steinacker et al. \\cite{stein02b}) have been developed. Nevertheless, it became soon clear that a 1-D geometry is often too restrictive. Distinct non-spherical features such as bipolar outflows (e.g. Bachiller \\cite{bachiller96}), bipolar reflection nebulae (e.g. Lenzen \\cite{lenzen87}) and disks (e.g. McCaughrean \\& O'Dell \\cite{mccau96}) are typical of many astronomical objects. Nowadays,  multidimensional codes implementing different methods and numerical schemes are being applied to treat the RT in such configurations.  In contrast to hydrodynamical simulations, benchmark tests for radiative transfer computations are rare. The only practical approach to test the reliability of RT calculations is to compare solutions of well-defined problems by several independent codes. This has been done for the 1D case by Ivezic et al. (\\cite{ivezic97}). A benchmark project for 1-D plane-parallel RT and vertical structure calculations for irradiated passive disks is available on the web \\footnote{http://www.mpa-garching.mpg.de/PUBLICATIONS/DATA/radtrans/ \\hspace*{1.1cm}benchmarks/}. As for 1-D RT in molecular lines, a comparison of results from different codes has been performed by van Zadelhoff et al. (\\cite{vanzadelhoff02}) \\footnote{see also: http://www.strw.leidenuniv.nl/$\\sim$radtrans/}. Going from spherical symmetry to 2D and 3D spatial configurations, we add two or three more variables to the RT problem. Numerically, this implies  10$^4$ or 10$^6$ more numbers to store when a decent resolution of 100 points in each variable is used. In addition, the geometry makes the solution of the integro-differential RT equation more complex. This explains why benchmark tests for 2D and 3D configurations are lacking. It also implies that reaching an agreement to the level of 1D RT computations using state-of-the-art computer equipment is unrealistic. A previous attempt to test 2D RT calculations has been made by Men'shchikov \\& Henning (\\cite{mensh97}). They compare results from their approximate method with those of a fully-2D program (Efstathiou \\& Rowan-Robinson \\cite{efstathiou90}) applying the same geometry.  Here, we propose to test the behaviour of five different RT codes in a well defined 2D configuration, point out advantages and disadvantages of the various techniques applied to solve the RT problem and provide benchmark solutions for the verification of continuum RT codes. As modelling sources with high optical depth and strong scattering is the challenge of multi-dimensional RT codes, we explicitly include a test case at the limit of the current computational capabilities. In Sect.~\\ref{sect:bprob} we briefly introduce the RT problem and we define our test case. Sect.~\\ref{sect:codes} is devoted to the description of the codes we used and to explain their differences. Solutions for the dust temperature and emerging SEDs are presented in Sect.~\\ref{sect:bres}. In the last section we discuss our results. ",
        "conclusions": "Before presenting our findings, we briefly discuss the general features of the computed SEDs for different viewing angles ($i$) and optical depths ($\\tau$). As mentioned in Sect. \\ref{defmodel}, optical depths are measured in the disk midplane from the inner to the outer boundary. Thus, the optical depths we refer to are the highest in the disk.  Both panels of Fig. \\ref{fig:sed} show clearly that the far-infrared region (longward 100~\\micron ) remains unaffected when varying the viewing angle. On the other hand, the short-wavelength part of the spectrum is strongly modified. When the disk is seen face-on the spectrum is dominated by unattenuated stellar radiation. As the disk inclination increases, more and more of the stellar flux is extincted by the dust in the disk. For the $\\tau_{\\mathrm v}=100$ case at $i=77.5\\degr$ this reduction amounts to a factor of  $e^{-100}$ in the visual. However, due to the high albedo of the dust grains, a large fraction of the stellar radiation is scattered above the disk into the line of sight. Dust scattering is also responsible for the excess of emission at stellar wavelengths seen at small disk inclinations (see left panel of Fig. \\ref{fig:sed} especially for high optical depths). The optical depth also affects the strong 10~\\micron{} feature produced by the SI -- O stretching. While in most cases the feature appears strongly in emission, for the $\\tau_{\\mathrm v}=100$ test case at $i=77.5\\degr$ the feature appears in absorption (see Fig. \\ref{fig:sed} right panel). The 20~\\micron{} feature is much weaker than the 10~\\micron , but it is already visible for the model with $\\tau_{\\mathrm v}$ = 1. All these features are in agreement with earlier radiative transfer computations of disks (e.g.~Efstathiou \\& Rowan-Robinson \\cite{efstathiou90},  Menshchikov \\& Henning \\cite{mensh97}).  Our aim in this paper is to provide benchmark solutions for the 2-D continuum radiative transfer problem in circumstellar disks. The problems we present have optical depths up to 100, which is actually the limit of current computational capabilities for most of the codes. The corresponding total dust mass in the disk of about 0.01 solar masses covers most of the observed disks around low mass stars. For more massive disks around intermediate and high mass stars as well as tori obscuring active galactic nuclei, the numerical strategies have to be modified, using e.g. the diffusion approximation for high optical depths. We used five independent radiative transfer codes that implement  different numerical schemes. We compare both the resulting temperature structure and the emerging SEDs. For the lowest optical depth case ($\\tau_{\\mathrm v}=0.1$) we also compared the results against a semi-analytic solution which treat scattering only as extinction term. The other three cases ($\\tau_{\\mathrm v}=1, 10, 100$) cannot be solved in a semi-analytic way, since multiple scattering and absorption-reemission events strongly affect the solution.  We find that the overall shape of the temperature distribution and of the emerging SEDs is well reproduced by all the codes. Differences in the temperature are smaller than 1~\\%  for all the codes in the most optically thin case. Even for the most optically thick model, differences in the temperature remain below $15$\\%. As for the SEDs, deviations among the codes are smaller than 3\\% at all wavelengths and disk inclinations for the most optically thin model. For the models with $\\tau_{\\mathrm v}=1$ and 10 at all disk inclinations and for the most optically thick case for disk inclinations of 12.5 and 42.5$\\degr$, differences do not exceed 10\\%. Only for the most optically thick case and an almost edge-on disk differences around 10~\\micron{} exceed 20\\% in the case of STEINRAY. We stress that this is the case for which the numerics is the most difficult: the codes have to treat both a very optically thin atmosphere and a thick disk midplane. Independent tests using two of the MC codes show that the frequency resolution cannot account for the infrared deviations among the codes in the almost edge-on disk and the most optically thick model. Grid resolution especially in radial direction together with cumulative numerical errors play a major rule. The presented results provide a robust way to test other continuum RT codes and demonstrate the possibilities of the current computational capabilities. Temperature distributions and SEDs for all the test cases are available at the web site: \\\\ http://www.mpia.de/PSF/PSFpages/RT/benchmark.html"
    },
    "0402/astro-ph0402161_arXiv.txt": {
        "abstract": "We constrain the mass profile and orbital structure of nearby groups and clusters of galaxies. Our method yields the joint probability distribution of the density slope $n$, the velocity anisotropy $\\beta$, and the turnover radius $r_0$ for these systems. The measurement technique does not use results from N-body simulations as priors. We incorporate 2419 new redshifts (included here) in the fields of \\ngroup\\ systems of galaxies with $z < 0.04$. The new groups have median velocity dispersion $\\sigma=360$ km s\\m. We also use 851 archived redshifts in the fields of 8 nearly relaxed clusters with $z < 0.1$. Within $R \\lesssim 2 \\rtwo$, the data are consistent with a single power law matter density distribution with slope $n = $\\grouprange\\ for systems with $\\sigma < 470$ km s\\m\\ and $n = $\\clusterrange\\ for those with $\\sigma > 470$ (95\\% confidence). We show that a simple, scale-free phase space distribution function $f(E,L^2) \\propto (-E)^{\\alpha-1/2} L^{-2 \\beta}$ is consistent with the data as long as the matter density has a cusp. Using this DF, matter density profiles with constant density cores ($n=0$) are ruled out with better than 99.7\\% confidence. ",
        "introduction": "The galaxies in a cluster account for less than 10\\% of its total mass. X-ray observations of the hot intracluster gas support the notion that most of the matter is distributed in a smooth, dark halo. It is of clear interest to use observations to deduce not only the total mass, but also the structure of this halo. For example, the logarithmic derivative of the density can be compared with N-body simulations of structure formation, allowing constraints on dark matter physics. These simulations show that the inner regions of halos composed of collisionless dark matter have power law density profiles $\\rho \\propto r^{-n}$ with $n \\sim 1-1.5$ \\citep{NFW,Fukushige97,Nakano99,Moore99}. However, if the dark matter particles interact through the weak or the strong force, constant density cores develop within a fraction of a Hubble time \\citep{Burkert00,Dave01, Balberg02}. Measurements of the slope $n$ in groups and clusters of galaxies are thus a potential a test of the collisionless nature of dark matter.   The methods for determining the masses of clusters fall into three broad categories: fluid dynamics, lensing, and stellar dynamics. The fluid dynamical methods, which rely on the X-ray properties of the intracluster medium, presuppose that the gas is either in hydrostatic equilibrium or in a time-independent cooling flow. Because measuring the gas temperature involves complex modeling of X-ray spectra, these methods offer at best $\\sim 50$ resolution elements across the brightest X-ray sources. Furthermore, recent X-ray observations indicate that the central regions of cooling flow clusters cannot be described self-consistently by the standard models, leaving the shape of the mass profile within $\\sim 200$ kpc of the cluster center uncertain\\footnote{Throughout this paper we use $H_0 = 100$ km s\\m\\ Mpc\\m.}  \\citep{Markevitch99,Soker01,Fabian01}.  Lensing models detect either strong or weak gravitational deflection of the light from distant sources. Strong lensing allows direct estimation of the surface mass density from the spectacular but rare giant arcs, e.g. CL0024+16 \\citep{Tyson98}. Weak lensing estimates depend on statistical reconstruction of the surface density from distortions in the shapes of field galaxies \\citep{Clowe00,Sheldon01}. These techniques are attractive because they make no assumptions about the equilibrium state of the cluster. However, they could suffer contamination from the large scale structure surrounding the cluster \\citep*{Metzler01}. This contamination may account for the disagreement between X-ray and lensing masses, and makes the determination of the true shape of the density profile difficult.  Stellar dynamical methods grow out of the century-old tradition, beginning with Jeans and Eddington, of modeling the structure of star clusters. Here, instead of addressing a self-gravitating system of stars, one regards galaxies as point masses adrift in a larger sea of dark matter. The fundamental dynamical problem is then to calculate the spherically symmetric gravitational potential that causes the observed motions of the galaxies. Obviously, the analogy to stellar systems is far from perfect; the ratio of the size of a galaxy to that of its host cluster is typically $\\approx$ 15 kpc / 1.5 Mpc $ = 0.01$, whereas for stars in a galaxy it is $\\approx 10^{-12}$. As a result, interaction cross sections are much larger in galaxy clusters. Furthermore, although stellar systems contain as many as $10^{12}$ members, clusters rarely have more than $\\approx 500$ luminous members, and the most abundant systems---groups of galaxies---have closer to $\\approx 30$ \\citep{Carlberg96, Mahdavi99}.  Still, if correctly applied, stellar dynamics can trace the gravitational potential of clusters at the largest and smallest scales as well as the other available techniques. Modeling of spherical infall patterns \\citep{Geller99,Rines00} can map the mass profile outside $\\approx 5$ Mpc. The equilibrium models we consider are sensitive to the shape of the dark matter halo in the innermost regions of clusters and groups.  The most popular equilibrium techniques make use of the moments of the data to constrain the depth or shape of the cluster potential. For example, the virial theorem yields an estimate of the total mass of the cluster \\citep{Heisler85,Biviano93,Oegerle95,Carlberg96,Girardi00}: \\begin{eqnarray} M & = & \\frac{3 \\pi}{2 G}\\sum_{i=1}^N v_{z,i} ^2 R_h, \\label{eq:virial} \\\\ R_h & = & \\sum_{i=1}^N R_i^{-1} \\\\ v_{z,i} & = & \\frac{c z_i - \\sum_{i=1}^N c z_i/N} {1+\\sum_{i=1}^N  z_i} \\label{eq:vzi}, \\end{eqnarray} where $v_z$ is the line-of-sight velocity in the center of mass frame \\citep{Danese}, $z_i$ is the redshift of the $i$th galaxy, $R_i$ is its projected distance from the cluster center, $c$ is the speed of light, and $G$ is Newton's constant.  The velocity moments may also be used to infer the structure of the potential in greater detail, using either analytic galaxy distribution functions \\citep{KentGunn} or the Jeans equation for a collisionless stellar system \\citep{Fabricant89,BinneyTremaine87,Hartog96,Carlberg97}, \\begin{equation} \\frac{1}{\\nu} \\dd{(\\nu \\sigma_r^2)}{r} + \\frac{2 (\\sigma_r^2-\\sigma_t^2)}{r} = - \\dd{\\Phi}{r}, \\label{eq:jeanstwo} \\end{equation} where $r$ is the true three-dimensional distance from the cluster center, $\\nu(r)$ is the number distribution of the galaxies, $\\sigma_t(r)$ and $\\sigma_r(r)$ are the tangential and the radial velocity dispersions, and $\\Phi$ is the gravitational potential.  To interpret the data, one usually chooses a form for $\\Phi(r)$ and the anisotropy parameter $\\beta(r) \\equiv 1-\\sigma_t^2/\\sigma_r^2$, and projects the solution to the Jeans equation to obtain $\\sigma_z(R)$, the theoretical line-of-sight velocity dispersion profile. This profile may be compared with a real cluster by splitting the galaxies into radial bins and calculating $\\sigma_z^2 \\propto \\sum v_z^2$ in each bin.  There are several disadvantages to using the velocity moments to constrain $\\Phi$. First, the quality of the observed velocity dispersion profile $\\sigma_z(R)$ is usually poor for clusters. Even with several hundred velocities, dividing the galaxies into radial bins produces a noisy profile that is not very informative about the radial variation of $\\sigma_z$. Second, even in the ideal limit of a perfectly observed $\\sigma_z(R)$, vastly different combinations of $\\Phi(r)$ and $\\beta$ can yield similar solutions to equation (\\ref{eq:jeanstwo}) \\citep{BinneyTremaine87}. Third, even if a unique solution is possible (e.g., by assuming a constant mass-to-galaxy ratio, $\\grad^2 \\Phi \\propto \\nu$), there is no guarantee that the solutions satisfy the requirement that the phase space density of the member galaxies be everywhere positive or zero \\citep{van00}. Finally, it is desirable to avoid binning the velocities altogether in order to make as powerful a use of the scarce data as possible. We therefore turn to maximum likelihood methods, which avoid some of the problems of the Jeans equations as applied to discrete systems \\citep{Merritt93,van00}.  Constructing and maximizing a suitable likelihood function guarantees a positive definite galaxy phase space distribution.  Here we apply the maximum likelihood method to nearby ($z \\lesssim 0.1$) systems of galaxies. Our goal is to derive joint constraints on $\\Phi$ and $\\beta$ for poor groups of galaxies as well as for rich clusters. To this end we conduct deep optical observations of the RASSCALS X-ray emitting galaxy groups \\citep{Mahdavi00}. We also assemble a catalog of published redshifts in 8 nearby relaxed clusters of galaxies (\\S \\ref{sec:data}). Using these samples, we construct an ensemble group and ensemble cluster which serve to constrain the galaxy phase space distribution in each type of system (\\S \\ref{sec:phase}). This distribution then serves as a maximum likelihood estimator of the gravitational potential and of the orbital structure of the galaxy population. In particular, we calculate joint five-dimensional confidence volumes in the central anisotropy, central matter density slope, total mass, matter core radius, and interloper fraction (\\S \\ref{sec:dynamics}). We discuss the implications of our work (\\S\\ref{sec:discuss}) and conclude (\\S \\ref{sec:conclude}). ",
        "conclusions": "Here we describe our constraints on the inner matter density slope $n$, the mass normalization $\\tilde{M}_{200}$, the transition radius $\\tilde{r}_{5/2}$, the velocity anisotropy $\\tilde{\\beta}$, and the interloper fraction $P_I$.  It bears repeating that our constraints result from maximization of the likelihood function described above, using the full unbinned data set. The binned profiles we show are for illustrative purposes and do not indicate ordinary $\\chi^2$ fitting.  \\subsection{Models with Constant-Density Cores} \\label{sec:freeg}  First we investigate the possibility that the ensembles have constant-density cores ($n=0$). In \\S\\ref{sec:df} we showed that such cores do not support radially anisotropic models; we are limited to $\\beta \\le 0$. Therefore the minimization occurs over five dimensions: ($\\tilde{\\beta}$,$\\gamma$,$\\tilde{r}_{5/2}$,$\\tilde{M}_{200}$,$P_I$).  The dotted line in figure \\ref{fig:results} shows the results of the minimization. We show the most likely grand total velocity histogram and galaxy surface density as given by equations (\\ref{eq:nvz})-(\\ref{eq:sigma}).  We find that matter profiles with constant-density cores do not produce a galaxy density profile as steep as our data.  Although the grand-total velocity histogram $N(\\tilde{v}_z)$ seems accurately reproduced, the model's galaxy density in the innermost regions is too small. The constant density models are disfavored with $q = 0.003$ for the low-$\\sigma$ ensemble, and $ 10^{-5}$ for the high-$\\sigma$ ensemble.  Note that the inconsistency with the galaxy surface density $\\Sigma$ is not the only discrepancy between the data and the $n=0$ model. Our $\\chi^2$ goodness-of-fit method described above shows that the $n=0$ models predict too many high-velocity galaxies at intermediate radii. This discrepancy is not easily visualized through the plots of $N(\\tilde{v}_z)$ or $\\Sigma(\\tilde{R})$.  Just to be sure that the $n = 0$ models are a poor fit, we also try relaxing the requirement that the orbits not be radially anisotropic. We fit the full range of negative to positive $\\beta$, and find that $\\beta \\approx 0.25, \\gamma \\approx 4$ maximize the likelihood for both ensembles. These values are formally unphysical, yielding a galaxy density that is steeper than the $n = 0$ total matter density near the central regions. However, even these fits are rejected with $q = 0.02$ for groups and $q = 0.003$ for clusters.  It is interesting to note (Figure \\ref{fig:results}) that the inconsistency occurs near $r = 0$, and not at larger radii. If our choice of mass model (equation \\ref{eq:mainpot}) were the cause for the poor fit, we would expect the inconsistency to occur at the larger radii, because that is where we impose a $\\rho \\propto r^{-4}$ behavior. We conclude that models with a constant density core ($n = 0$) are at best barely consistent with the data.   \\begin{figure*} \\resizebox{7in}{!}{\\includegraphics{result1.eps}} \\figcaption{Predictions (not fits) of the maximum-likelihood analysis of unbinned distance and velocity data. Shown are the surface number density (\\emph{top}) and the grand total velocity histogram (\\emph{bottom}).  The solid line represents the most likely model overall, whereas the dashed line shows the most likely model when the inner slope of the matter density is forced to equal 0. The latter model is rejected with 99.7\\% or better confidence. The dotted vertical line represent the completeness limit of the high-$\\sigma$ ensemble. \\label{fig:results}} \\end{figure*}  \\subsection{Models with Fixed $\\gamma$} \\label{sec:freen}  Now we allow the inner slope of the mass density $n$ to vary, maximizing $p(\\tilde{R},\\tilde{v}_z|\\mb{a})$ for several different values of $\\gamma$, the slope of the galaxy density at large radii. The jointly constrained parameter set is therefore ($n$,$\\tilde{\\beta}$,$\\tilde{r}_{5/2}$, $\\tilde{M}_{200}$,$P_I$).  As described in \\S \\ref{sec:df}, $\\gamma$ is also related to the slope of the energy term in the distribution function (equation \\ref{eq:df}). The larger the value of $\\gamma$, the more peaked the distribution of galaxies with small kinetic energies.  Our data can provide a lower, but not an upper limit on $\\gamma$. For both the low- and the high-$\\sigma$ ensembles, the quality of the fit increases with $\\gamma$, asymptotically approaching a fixed value as $\\gamma \\rightarrow \\infty$. Our simulations (\\S\\ref{sec:sims}) show a similar effect, and suggest that the smallest value of $\\gamma$ that gives an acceptable fit is likely to give the most accurate results. Therefore, instead of constraining $\\gamma$ continuously, we list the minimum value of $\\gamma$ that yields a quality of fit $q > 0.1$. This constraint corresponds to $\\gamma=8$ for the low-$\\sigma$ and $\\gamma=12$ for the high-$\\sigma$ ensemble.  Figures \\ref{fig:results} and \\ref{fig:conf} show the results.  Aside from difference in the value of $\\gamma$, the high- and low- velocity dispersion systems provide very similar fits: the galaxy orbits are consistent with isotropic ($\\tilde{\\beta} = 0$), and both ensembles have $n\\approx2$ and transition radius $r_{5/2}$ much greater than $\\rtwo$. In other words, the matter distributions for both samples are consistent with a singular isothermal sphere within the limits of the survey.  \\begin{figure*} \\begin{center} \\resizebox{7in}{!}{\\includegraphics{multcont.eps}} \\figcaption{Results of the likelihood maximization for free $n$. Solid contours represent the low-$\\sigma$ ensemble and dashed contours represent the high-$\\sigma$ ensemble. Shown are marginalized joint 68\\% and 95\\% probability contours in the matter density slope $n$, normalized anisotropy $\\tilde{\\beta}$, mass normalization $M_{200}$, transition radius $r_{5/2}$, and interloper fraction $P_I$. \\label{fig:conf}} \\end{center} \\end{figure*}  \\subsection{Discussion}  Neither the cluster nor the group data supports a total matter distribution with a constant-density core ($n=0$). This result is in keeping with recent studies of more distant clusters in the X-ray \\citep{Lewis03,Arabadjis02,Allen02} as well as weak lensing maps \\citep{Gavazzi03,Clowe02,Clowe00}. Optical observations of cluster velocity data have also yielded similar results. For example, \\cite{Biviano03} uses the Jeans equation and isotropic velocity dispersion profiles to rule out constant-density cores in the Two Degree Field Galaxy Redshift Survey \\citep{Colless01}.  Our results suggest that within a projected region $R \\lesssim 2\\rtwo$, the matter distribution in both low- and high-$\\sigma$ systems of galaxies is close to a single power law with $n=2$. The density declines rapidly only outside $2\\rtwo$. This result is consistent with N-body simulations. For example, in the original paper describing the simulations of NFW, the density profiles of dark matter halos are consistent with a single power law of slope 2 between $0.1 \\rtwo$ and $2 \\rtwo$ (see their Figure 3). The only discrepancy, or flattening, occurs at radii smaller than $0.1 \\rtwo$, where our datasets have 56 (low-$\\sigma$) and 131 (high-$\\sigma$) members, perhaps insufficient to provide a robust constraint.   A major limitation of our method is the fixed form of the galaxy distribution function $f(\\ene,L^2)$. By modeling the DF as a power law in energy and angular momentum, we neglect all other possible forms for the DF, some of which may indeed be consistent with an $n=0$ model. On the other hand, the $\\chi^2$ tests we perform reassure us that the data are at least statistically consistent with cuspy matter distributions.  The work of \\cite{van00}, whose methods we adapt and extend, involves a similar analysis, and it is instructive to compare the two different approaches. Instead of fixing the form of the DF, they fix the form of the surface number density $\\Sigma(R)$, and calculate the DF using an inversion of the integral in equation (\\ref{eq:dfint}). As a result, their DF is more general than ours, $f \\propto g(\\ene) L^{-2 \\beta}$, with $g$ being the result of the inversion. Their method has the advantage of generality, but the disadvantage that it requires a two-stage fit: first $\\Sigma(R)$ must be fit and fixed, and then $\\beta$ and $n$ measured separately, without an indication of how changes in the $\\Sigma(R)$ fit would affect the resulting constraints. Our method sacrifices generality by setting $g(\\ene) = \\ene^{\\alpha-1/2}$, but gains the advantage that all but one of the parameters are fit simultaneously ($\\gamma$ is varied discretely). As a result, the correlation among the parameters is easy to understand via marginalized confidence contours (Figure \\ref{fig:conf}).  Another difference between our work and \\cite{van00} is the sampling of the parameter space. We constrain the continuous five-dimensional region $(n,\\beta,M_0,r_0,P_I)$, whereas \\cite{van00} sample discrete points within that space. They find that an NFW profile with $n=1$ matches the CNOC data, but they do not consider $n=2$ models; we find that our clusters and groups have steeper matter distributions, with the best fit model closer to $n=2$.  The fixed functional form of our DF can also explain the large transition radii. Careful examination of equation (\\ref{eq:nu}) shows that increasing $\\gamma$, the slope of the galaxy density as $r \\rightarrow \\infty$, also steepens the slope of the galaxy density $n_g$ near $r = 0$. It is still the case that $n_g = 2 \\beta$ at $r = 0$, but the derivative of the slope $d n_g / d r$ becomes larger and larger as $\\gamma$ is increased. Our data have a steep galaxy density slope near $r = 0$, and there are only two ways to fit this: increasing $\\beta$ (and thus generating more radially anisotropic orbits) or increasing $\\gamma$ (and thus steepening the overall galaxy energy distribution and hence the galaxy density). This is where the velocity distribution enters. As Figure \\ref{fig:results} shows, the grand total velocity histogram is not strongly peaked at $\\tilde{v}_z = 0$ (in fact, it is consistent with a Gaussian distribution). Dramatically increasing $\\beta$ would yield a velocity distribution that is much more sharply peaked than the data. Thus the only choice left the model is to increase $\\gamma$ to larger values, $\\gtrsim 8$. However, the slope of the galaxy density is not as steep as 8 at the limits of the survey.  Hence, a large transition radius $r_0$ is required to fit the outer regions of the galaxy number density profile. At the same time, the velocity data allow the matter profile to be nearly a power law within the region constrained by the data. Note that the upper limits on $\\tilde{r}_{5/2}$ are possible because if $\\tilde{r}_{5/2}$ becomes too large, then the galaxy density profile begins to resemble a single power law as well, and the position data do not favor this limit (Figure \\ref{fig:results}).  Another shortcoming is that our method constrains the total matter density, rather than the dark matter density alone. While the dark component is thought to dominate the mass, it does not do so overwhelmingly---in typical clusters, $\\sim 25\\%$ of the mass is in baryons, chiefly in the form of the X-ray emitting medium. In the innermost regions of clusters with a dominant central galaxy, it is actually the stellar mass density and not the dark matter density that dominates the total distribution \\citep{Koopmans03}. Lensing and stellar dynamical models which take into account the starlight separately from the dark component show conflicting results---while the central regions of some clusters exhibit flat $n=0$ \\citep{Sand02} dark matter densities, others show evidence for steep $n \\sim 2$ dark matter profiles \\citep{Davis03}.  It would be of great value to constrain the inner density slope $n$ independently of the slope at infinity. Unfortunately, because the data are sparse, we cannot derive true constraints on the outer galaxy density slope $\\gamma$, and measuring the outer total matter density would be even more difficult. With $10^5$ redshifts a true, simultaneous constraint on the outer and inner slopes would be possible; such large data sets could allow a direct measurement of the phase space distribution function \\citep{Merritt93b}, without requiring us to parameterize the shape of the distribution function, as in equation (\\ref{eq:df}). Data sets that are only factors of two larger are unlikely to make that difference. Also, leveraging independent constraints on the inner and outer slopes by extending the survey beyond $\\rtwo$ is unlikely to provide better constraints. In these outer regions, most of the galaxies are experiencing spherical infall, which makes equilibrium models clearly unapplicable, and caustic modeling techniques preferable \\citep{vanHaarlem93,Diaferio99,Geller99,Rines00}. However, one can envision a technique that combines equilibrium dynamics at small scales with the study of the infall caustics outside the virialized regions of the cluster. In this way one could obtain a total mass profile over several decades in $r$.  The most promising method of increasing the accuracy of this analysis is to extend the known membership of groups and clusters to dwarf galaxies with absolute magnitudes as faint as $M_R = -11$. A typical group with velocity dispersion $\\sigma \\approx 300$ km s\\m\\ is likely to have $\\approx 60$ faint dwarf members within 1 Mpc \\citep{Trentham02}. A large spectroscopic survey of well-selected galaxies brighter than $m_R = 22$ in our group sample would likely yield triple the current membership count and allow for much more robust constraints."
    },
    "0402/astro-ph0402482_arXiv.txt": {
        "abstract": "s{We present the results from the analysis of 26 extragalactic radio sources of type FRII which were observed with the VLA at 5 GHz and around the 1.4 GHz band. The sources were selected to have redshifts in the range $ 0.3<z<1.3$, radio powers between $6.9 \\times 10^{26} {\\rm WHz^{-1}}<P_{151 {\\rm MHz}}<1.3 \\times 10^{28} {\\rm WHz^{-1}}$ and angular size $\\theta \\ge 10''$. We found that the depolarisation and the rms variations in the rotation measure increased with redshift. The flux values obtained from the observations were used to derive by means of analytical modelling the jet--power, density of the central environment, age of the source and its lobe pressure and the results were then compared with the observations. We find no significant correlations with the density parameter suggesting that the depolarisation and the rms variations in the rotation measure are indicative of the environment becoming more disordered rather than denser.  The age and size of a source are correlated and both were found to be independent of redshift and radio--power. Jet--power strongly correlated with the radio--power. The lobe pressure was found to be anti--correlated with size which could explain why there are no sources beyond a few Mpc in size. We found no significant correlation between size and density which demonstrates that the sample is a fair representation of the population.} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0402/gr-qc0402126_arXiv.txt": {
        "abstract": "It was recently shown by Whinnett \\& Torres \\cite{WT} that the phenomenon of spontaneous scalarization (SC) in compact objects (polytropes) was accompanied also by a {\\it spontaneous violation of the weak energy condition} (WEC). Notably, by the encounter of negative-energy densities at several star places as measured by a static observer. Here we argue that the negativeness of such densities is not generic of scalar tensor theories of gravity (STT). We support this conclusion by numerical results within a class of STT and by using three realistic models of dense matter. However, we show that the ``angular parts'' of the additional conditions for the WEC to hold $\\rho_{\\rm eff} + T^{i\\,({\\rm eff})}_{\\,\\,\\,i} \\geq 0$ tend to be ``slightly violated'' at the outskirts of the star. ",
        "introduction": "Scalar-tensor theories of gravity (STT) are one of the simplest alternatives to Einstein's general relativity (GR) \\cite{damour0}. A particular example of these, is the Brans-Dicke theory. The general feature of STT is the introduction of a new fundamental scalar field which couples to gravity in a non-minimal fashion (NMC). Nonetheless, the NMC is such that it preserves the Equivalence Principle, because the scalar-field couples to the ordinary matter only through the spacetime metric.  The NMC in STT unlike simple Einstein-Higgs systems, gives rise to field equations that can differ substantially from the usual Einstein's equations. Therefore, when applied to the astrophysical and cosmological settings these theories can produce severe deviations from the general relativity predictions. The fact that STT introduce in general only a few new parameters make them in principle easily testable. Perhaps the most notable probes in this direction are those connected with the binary pulsar, the primordial nucleosynthesis, the Cosmic Microwave Background, and the luminosity distance with Type-II Supernova. The success of GR lies in the fact that with suitable matter models, the theory is compatible with the available information in the above systems. Modifications of the gravitational sector in the way of STT lead thus to a serious change in the geometry-matter relation that can upset the compatibility of the theory with the observations. The stringent bounds on STT make difficult to accommodate these theories to all the observations while at the same time predicting new observable phenomena. One of the potentially observable new phenomenon is precisely the phenomenon of spontaneous scalarization (SC) arising in compact objects. As it was first shown by Damour \\& Esposito-Far\\`ese \\cite{damour1}, STT can induce non-perturbative effects in neutron stars (NS) which, roughly speaking, consists in a phase-transition that endows a compact star with a new global quantity, the {\\it scalar charge}. This ``charge'' is the analogous of the {\\it magnetization} in ferromagnets at low temperatures and the central-energy density (or equivalently the baryon mass) of the object plays in turn the role of the temperature in spontaneous magnetization. Damour \\& Esposito-Far\\`ese \\cite{damour1,damour2} also shown that the consequences of SC can be dramatic even if the boundary-conditions (asymptotic conditions) for the scalar-field are chosen so as to satisfy the solar-system bounds. In particular, they showed that the binary-pulsar dynamics is very sensitive to the new phenomenon, and therefore, that STT can be further constrained even if they successfully pass the local and the cosmological tests.  In this paper we will not be concerned with the confrontation of STT with the experiments, but rather with a result obtained by Whinnett \\& Torres \\cite{WT} who presented evidence that the the phenomenon of SC in NS is linked to a violation of the weak energy condition (WEC). Notably, Whinnett \\& Torres \\cite{WT} found that the effective energy density of STT as measured by a static observer and defined as $\\rho_{\\rm eff}:=n^a n^b G_{ab}/(8\\pi)$ (where $G_{ab}$ is the Einstein tensor and $n^a$ a time-like vector parallel to the static Killing vector $\\xi^a_t$) is negative and suggested that such a finding might be independent of the class of STT used (although depending of the chosen value of the NMC constant of the theory). A negative $\\rho_{\\rm eff}$ is a sufficient condition for the violation of the WEC. The authors argue that the violation of the WEC in NS might led to their instability (even for masses quite below to their corresponding Landau-Tolman-Volkoff-Oppenheimer limit), a very undesirable feature that would presumably rule out the STT as an alternative for a viable spacetime theory.  While it is true that the effective energy density $\\rho_{\\rm eff}$ of STT as defined above is not automatically positive definite due to the presence of second order derivatives of the scalar field, the main goal of this paper is to show that there are classes of STT that do not violate the condition $\\rho_{\\rm eff} \\geq 0$ for the problem at hand. As we shall argue in Sec. III, the violation of the condition $\\rho_{\\rm eff} \\geq 0$ as found by Whinnett \\& Torres \\cite{WT} is linked to the use of a parametrization {\\it \\`a la} Brans-Dicke which prejudices the choice of the specific theory which can potentially violate such condition. In our case, we use a parametrization of the STT where the effective gravitational constant can be chosen to be positive-definite. In such a case, as it is shown in Sec. III, the theories parametrized in both ways may be very different (the transformation that would be necessary to change from one parametrizaction to the other might not be well defined).  Altough the class of STT analyzed by us satisfy the condition $\\rho_{\\rm eff} \\geq 0$ for the NS of interest, they are prone to violate ``slightly'' the angular parts of the additional conditions $\\rho_{\\rm eff} + T^{i\\,({\\rm eff})}_{\\,\\,\\,i} \\geq 0$ for the WEC to hold. Since this violation is rather weak and occurs when the energy-density and pressure of the star fluid is very low (near the star surface), we argue that the stability of the neutron star is not in jeopardy, and therefore that a large class of STT cannot be considered as unviable spacetime theories only by this fact. ",
        "conclusions": "Our main conclusion is that the large violations of the WEC deep inside the neutron stars as an additional feature to the phenomenon of spontaneous scalarization is not generic of all the classes of scalar tensor theories of gravity as the work of Torres \\& Whinnett work might be thought to suggest. We have argued that there are classes of STT (at least a subset with $F(\\phi)>0$ and $\\partial^2_{\\phi\\phi} F> 0$) where $\\rho_{\\rm eff} \\geq 0$ and where the violations of the WEC in neutron stars occur only near and beyond the surface of the star via the angular parts of $\\rho_{\\rm eff} + T^{i\\,({\\rm eff})}_{\\,\\,\\,i}$. Strong numerical evidence supports this conclusion.   \\medskip"
    },
    "0402/astro-ph0402211_arXiv.txt": {
        "abstract": "{ We present a new tool for the photometric estimate of stellar masses in distant galaxies. The observed source spectral energy distributions are fitted by combining sets of various single stellar populations, with different normalizations and different amounts of dust extinction, for a given (Salpeter) IMF. This treatment gives us the best flexibility and robustness when dealing with the widest variety of physical situation for the target galaxies, including inactive spheroidal and active starburst systems. We tested the code on three classes of sources: complete samples of dusty ISO-selected starbursts and of K-band selected ellipticals and S0s in the HDF South, and a representative sample of $z\\sim 2$ to 3 Lyman-break galaxies in the HDF North. We pay particular attention in evaluating the uncertainties in the stellar mass estimate, due to degeneracies in the physical parameters, different star formation histories and different metallicities. Based on optical-NIR photometric data, the stellar masses are found to have overall uncertainties of a factor of $\\sim 2$ for E/S0s,  while for the starburst population these rise to factors 2$-$5 (even including ISO/15$\\mu$m photometric data), and up to $\\ge 10$ for Ly-break galaxies. Our analysis reveals in any case the latter to correspond to a galaxy population significantly less massive ($M<$ a few $10^{10}M_\\odot$) than those observed at lower redshifts (for which typically $M>$ several $10^{10}M_\\odot$), possibly indicating substantial stellar build-up happening at $z\\sim 1$ to 2 in the field galaxy population. Using simulated deep SIRTF/IRAC observations of starbursts and Lyman-break galaxies, we investigate how an extension of the wavelength dynamic range will decrease the uncertainties in the stellar mass estimate, and find that they will reduce for both classes to factors of 2$-$3, comparable to what found for E/S0s and good enough for statistically reliable determinations of the galaxy evolutionary mass functions. ",
        "introduction": "One of the still open critical questions of modern cosmology is to understand the epoch at which galaxies assembled the bulk of their stellar content. In the so-called monolithic scenario, the assembly of galaxies took place on rapid timescales at high redshifts, then galaxies evolved passively to present days. On the contrary, in the hierarchical scenario (Kaufmann \\& Charlot, 1998) galaxy formation is predicted to be a more continuous process and elliptical galaxies to assemble through merging of lower mass disc galaxies at moderate redshifts.  Many authors have analysed the evolution of the global star formation rate (SFR) through cosmic history, back to z$\\ge$3, using different observational tracers of the SFR in distant galaxies. It is generally accepted that galaxies produced stars more actively in the past than today, but the true rates of star formation are affected by a variety of uncertainties and biases, expecially related to the amount of dust in galaxies and its effect on the SFR tracers.  A complementary approach consists in measuring the dynamical or stellar masses of distant galaxies, instead of the instantaneous SFR. Dynamical masses are directly connected to theoretical predictions, but very difficult to measure, requiring high spatial and spectral resolution spectroscopy of selected samples of faint high-redshift galaxies. Stellar masses, on the contrary, are less well determined by theory, but have the advantage that can be derived using multiband optical and near-IR photometry as a powerful alternative to time-expensive spectroscopic investigations. Indeed galaxy near-IR SEDs show a moderate dependence on the age of the contributing stars (e.g. Franceschini \\& Lonsdale, 2003), because in a typical galaxy the stellar mass is dominated by low-mass stars, with evolutionary timescales of the order of the Hubble time. As discussed by several authors (e.g. Lancon et al. 1999, Origlia \\& Oliva 2000), these moderate-mass stars emit predominantly in the near infrared (NIR, J-to-K restframe bands), especially during their cool giant phase, and are only slightly affected by dust extinction.  In this work we present a new spectro-photometric synthesis tool, aimed at the estimate of stellar masses in galaxies from a thorough analysis of their spectral energy distributions (SEDs). The tool combines a set of single stellar populations of different ages, assuming different star formation rates and different amounts of extinction for each. The code is tested by fitting the optical-NIR broadband SEDs of two samples of intermediate ($z=0.5-1.5$) redshift galaxies, luminous IR-selected starbursts in the HDFS (Franceschini et al., 2003) and mostly passively-evolving K-band selected ellipticals from Rodighiero et al. (2001). A benchmark higher-redshift population, also targeted by our analysis, was extracted from a sample of $z\\sim 2-3$ Lyman-break galaxies (Papovich et al., 2001).  We pay particular attention to the uncertainties in the mass estimates, due to degeneracies in extinction, age and star formation history, and discuss how forthcoming near-to-mid infrared data from SIRTF will further constrain the photometric mass estimate for these sources.  The paper is structured as follows. Section 2 describes the three samples of selected sources. Section 3 presents the model tools used to fit the observed optical-NIR spectral energy distributions of the sources and the different star formation histories assumed for the two classes of galaxies analyzed. Section 4 discusses our results, with a particular care for degeneracies. Section 5 shows perspectives for future SIRTF/IRAC observations and Section 6 summarizes our conclusions. We assume a $H_0=65$ $[$km s$^{-1}$ Mpc$^{-1}]$, $\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$ cosmology. ",
        "conclusions": "We have presented a new spectrophotometric synthesis code for the integrated light of distant and high-redshift galaxies. This tool is aimed, in perspective, at investigating various fundamental physical parameters, like the age of the dominant stellar populations, extinction, SF history, ongoing star-formation rate, etc. We have concentrated in this paper to study in particular the ability of the code to constrain the stellar mass from SEDs fitting. The tool has been tested on a sample of luminous infrared starbursts detected by ISO in the Hubble Deep Field South at 15 $\\mu$m (Franceschini et al., 2003), on a HDFS K-band selected sample of ellipticals and S0s (Rodighiero et al. 2001), and on a set of the Lyman-break galaxies in the HDF South from Papovich et al. (2001).  We have fitted the optical-NIR SEDs (from U to K bands) of the mid-IR starbursts by assuming a discrete star formation history between 1.2\\,10$^{10}$ years ago and today, by means of the combination of up to 10 single stellar populations. Each SSP was weighted by different SFR's and absorbed by different amounts of dust. For a best exploitation of the available constraints we also included in the analysis the 15-$\\mu$m ISO flux, to assess the contribution of extinguished young stars to the emitted spectrum. Based on these data, our estimated mass uncertainties range from a factor $\\sim$2 to occasionally a factor of 5. In such dust-obscured starburst galaxies, substantial uncertainty in the mass estimate comes from the still possible existence of strongly extinguished young stars with low M/L ratios, contributing to the far-IR flux but undetectable in the rest-frame optical.  We have also analysed a sample of morphologically-selected E/S0 galaxies, at z$<$1.5.  Their SEDs have been reproduced as the combination of stellar populations with solar metallicities, or by combining SSPs of different metallicities (Z=0.02 and Z=0.008). Additional solutions have been sought with a parameterized continuous sequence of stellar ages. The SEDs of spheroidal galaxies with red colors are well fitted by combining intermediate and old populations only, while for ellipticals with relatively blue colors (the majority in our complete sample) acceptable solutions are found only by assuming some some recent or even ongoing star formation. The three different methods provide results well consistent with each other.  For Lyman-break galaxies the available rest-frame optical-UV data are much less constraining. The consequences of assuming relatively old ($\\sim 10^9$ yrs) stellar populations in these systems have been discussed and it has been found that the lack of long-wavelength data imply uncertainties in the mass estimate of up to a factor of $\\simeq 10$. In spite of these large uncertainties, there are indications that typical Lyman-break galaxies correspond to a galaxy population significantly less massive then those observed at lower redshifts, possibly indicating substantial stellar build-up to ocurring at $z\\sim 1$ to 2 in the field galaxy population.  The situation is expected to significantly improve, particularly for higher-z galaxies, with the forthcoming SIRTF/IRAC observations. Our simulations show that such an extension of the wavelengths dynamic range will reduce the mass uncertainties to factors of 2$-$3 for various classes of galaxies up to z$\\sim$3. This promises to be good enough for statistically reliable determinations of the galaxy evolutionary mass functions."
    },
    "0402/astro-ph0402027_arXiv.txt": {
        "abstract": "{ During the period Aug.23-Sept.24 2003, the INTEGRAL observatory performed an ultra deep survey of the Galactic Center region with a record sensitivity at energies higher than 20 keV. We have analized images of the Galactic Center region obtained with the ISGRI detector of the IBIS telescope (15-200 keV) and present here a catalog of detected sources. In total, 60 sources with a flux higher than 1.5 mCrab have been detected. 44 of them were earlier identified as Galactic binary systems, 3 are extragalactic objects. 2 new sources are discovered. } ",
        "introduction": " ",
        "conclusions": ""
    },
    "0402/astro-ph0402301_arXiv.txt": {
        "abstract": "Numerical turbulence with hyperviscosity is studied and compared with direct simulations using ordinary viscosity and data from wind tunnel experiments. It is shown that the inertial range scaling is similar in all three cases. Furthermore, the bottleneck effect is approximately equally broad (about one order of magnitude) in these cases and only its height is increased in the hyperviscous case--presumably as a consequence of the steeper decent of the spectrum in the hyperviscous subrange. The mean normalized dissipation rate is found to be in agreement with both wind tunnel experiments and direct simulations. The structure function exponents agree with the She-Leveque model. Decaying turbulence with hyperviscosity still gives the usual $t^{-1.25}$ decay law for the kinetic energy, and also the bottleneck effect is still present and about equally strong. ",
        "introduction": "In recent years there has been growing awareness of the detailed structure of the kinetic energy spectrum of hydrodynamic turbulence. In addition to the basic Kolmogorov $k^{-5/3}$ spectrum with an exponential dissipation range there are strong indications of intermittency corrections (possibly throughout the entire inertial range) and there is also the so-called bottleneck effect \\cite{bottleneck,LohseMuellerG}, i.e.\\ a shallower spectrum near the beginning of the dissipative subrange; see also Ref.~\\cite{SJ93}. These features can be seen both in high resolution simulations \\cite{Kan03} and in measurements of wind tunnel turbulence \\cite{PKJ03}.  Over the past few years it has become evident that in numerical turbulence the bottleneck effect is rather pronounced \\cite{PorterWoodward1998,GotohFukayama2001,Kan03}. However, some of the simulations used hyperviscosity or other kinds of subgrid scale modeling. Hyperviscosity has frequently been used in turbulence studies in order to shorten the dissipative subrange \\cite{BLSB81,MFP81,McW84,PP88,BO95}. However, hyperviscosity has also been suggested as a possible source of an artificially enhanced bottleneck effect \\cite{BSC98,BM00}. Meanwhile, the apparent discrepancy in the strength of the bottleneck effect between simulations and experiments has been identified as being due to the difference in the diagnostics: in wind tunnel experiments one is only able to measure one-dimensional (longitudinal or transversal) energy spectra, while in simulations one generally considers shell integrated three-dimensional spectra. The two are related by a simple integral transformation \\cite{Batchelor,Hinze,MoninYaglom}. It turns out that, while the bottleneck effect can be much weaker or even completely absent in the one-dimensional spectrum, it is generally much stronger in the three-dimensional spectrum \\cite{DHYB03}.  In order to see the bottleneck effect in simulations, it is important to have sufficiently large resolution of around $1024^3$ meshpoints. This raises the question to which extent the bottleneck effect seen in simulations with hyperviscosity is an artifact or a real feature that becomes noticeable only above a certain resolution. It is thus possible that the reason for an exaggerated bottleneck effect in the hyperviscous simulation is related to the fact that hyperviscosity increases the effective resolution beyond the threshold above which the bottleneck effect can be seen.  In this paper we consider forced hydrodynamic turbulence using hyperviscosity proportional to $\\nabla^6$ (instead of the usual $\\nabla^2$ viscosity operator). We find that the bottle\\-neck effect is enhanced in amplitude--but not in width, compared with direct simulations at the currently largest resolution of $4096^3$ on the Earth Simulator \\cite{Kan03}. One of the important results of these very high resolution simulations is that an inertial range begins to emerge that is clearly distinct from the bottleneck effect. Furthermore, the (negative) slope in the inertial range is steeper than the standard Kolmogorov power law exponent of $5/3$ by about $0.1$, so it is approximately $1.77$.  As in earlier papers \\cite{DHYB03}, we consider weakly compressible turbulence using an isothermal equation of state. The root mean square Mach number is between $0.12$ and $0.13$; for this type of weakly compressible simulations, we find that the energies of solenoidal and potential components of the flow have a ratio $E_{\\rm pot}/E_{\\rm sol} \\approx 10^{-4}\\mbox{--}10^{-2}$ for most scales; only towards the Nyquist frequency the ratio increases to about $0.1$. Compressibility is therefore not expected to play an important role. ",
        "conclusions": "The present investigations have shown that turbulence simulations with hyperviscosity are able to reproduce virtually the same inertial range scalings as simulations with ordinary viscosity. Specifically, the structure function exponents show scaling behavior that is consistent with the She-Leveque \\cite{SL94} model. However, the transversal structure functions show a slightly higher degree of intermittency than the longitudinal ones. This, in turn, is quite consistent with a number of turbulence simulations by other groups \\cite{gotoh02,PPW02}. A possible explanation for the difference between longitudinal and transversal structure functions has been offered by Siefert \\& Peinke \\cite{SP03}, who find different cascade times for longitudinal and transversal spectra. The spectra show inertial range scaling similar to that found both in wind tunnel experiments \\cite{PKJ03} and in very high resolution direct simulations \\cite{Kan03}. In all three cases (hyperviscous and direct simulations as well as wind tunnel experiments) the inertial range spectrum is found to be compatible with the $k^{-1.77}$ behavior found by Kaneda et al.\\ \\cite{Kan03}. As discussed above, this result is not compatible with the results from the structure function scalings and the She-Leveque relation. However, we believe that the presently resolved inertial range is still too short to distinguish conclusively between 1.77 and the She-Leveque value of 1.70. Also, the simulation data of decaying turbulence suggest a weaker correction of 0.03, giving a slope of 1.70 that is compatible with the She-Leveque scaling.  Another important result is that the width of the bottleneck seems to be independent of the use of hyperviscosity, and that only its height increases with the order of the hyperviscosity. This result is also confirmed in the case of decaying turbulence. Finally, we note that the normalized dissipation rate is independent of the Reynolds number, and that the asymptotic value of $C_{\\epsilon}\\approx 0.5$ is in agreement with both experimental and numerical results \\cite{PKW02,Kan03}.  One should of course always be concerned about the possible side effects of using hyperviscosity. One worry is that hyperviscosity may actually affect almost all of the inertial subrange \\cite{BSC98,BM00}. The current simulations confirm that the bottleneck requires at least an order of magnitude in $k$-space, and so does the dissipative subrange, leaving almost no inertial range at all--even in a simulation with $1024^3$ meshpoints. Thus, using hyperviscosity appears to be a reasonable procedure for gaining information about the inertial range at moderate cost, even though one should still use a reasonably high resolution to isolate true inertial range features from those in the bottleneck subrange. On the other hand, hyperviscosity is not a universally valid approximation. An example is in magnetohydrodynamics when magnetic helicity is finite and a large scale magnetic field builds up in a closed or fully periodic box \\cite{BS02}. As long as it is possible to understand the origin of peculiar features arising from hyperviscosity or hyper-resistivity (as is the case in helical hydromagnetic turbulence) there may well be circumstances where turbulence with hyperviscosity can provide a useful model for certain studies. One should bear in mind, however, that the height of the bottleneck depends on the order of the hyperviscosity. For example, if we choose $n > 3$ in \\Eq{22prelim}, the height of the bottleneck will be even more exaggerated \\cite{BO95}."
    },
    "0402/astro-ph0402137_arXiv.txt": {
        "abstract": "Studies of high-redshift galaxies usually interpret offsets from the Tully-Fisher (TF) relation as luminosity evolution.  However, apparent luminosity offsets may actually reflect anomalous velocity widths. Rotation curve anomalies such as strong asymmetries or radial truncation are probably common in high-$z$ samples, due to frequent galaxy interactions and in some cases low S/N data, although low physical resolution may mask these anomalies.  In this paper we analyze well-resolved, one-dimensional optical emission-line rotation curves from two low-$z$ samples: the Close Pairs Survey, which contains a high frequency of interacting galaxies, and the Nearby Field Galaxy Survey (NFGS), which represents the general galaxy population.  Unlike most low-$z$ TF samples, but in the spirit of many high-$z$ samples, these surveys reflect the natural diversity of emission-line galaxy morphologies, including peculiar, interacting, and early-type galaxies.  We adopt objective, quantitative criteria to reject galaxies with severe kinematic anomalies, and we use a statistical velocity width measure that is insensitive to minor kinematic distortions.  Severely anomalous galaxies are roughly twice as frequent in the Close Pairs Survey as in the NFGS, and these galaxies' TF offsets collectively resemble the ``differential luminosity evolution'' claimed in some high-$z$ studies, with larger offsets at lower luminosities.  With the anomalous galaxies rejected, however, the TF relations for the Close Pairs Survey and the NFGS are quite similar.  Furthermore, the two surveys follow very similar relations between color and TF residuals.  The Close Pairs Survey color--TF residual relation extends to bluer colors and brighter TF residuals.  Strong outliers from this relation are virtually always kinematically anomalous.  As a result, the color--TF residual relation can serve as a powerful tool for separating reliable luminosity offsets from offsets associated with kinematic anomalies.  This tool may prove especially useful at high $z$, where direct detection of kinematic distortions is not always feasible.  Although we cannot reliably measure luminosity evolution for galaxies with kinematic anomalies, the TF offsets associated with these anomalies may offer a sensitive probe of evolution in the frequency and intensity of mergers and interactions on different mass scales.  We perform a preliminary reanalysis of high-$z$ TF data from the FORS Deep Field and find: (1) overall luminosity evolution of $\\sim$0.3 mag; (2) strong slope evolution driven by kinematically anomalous galaxies, which show TF offsets of up to $\\sim$2 mag at low luminosities; and (3) an additional zero-point offset of $\\sim$0.2 mag linked to kinematically anomalous galaxies. ",
        "introduction": "\\label{sc:intro}  As the fundamental scaling relation between luminosity and rotation velocity for disk galaxies, the Tully-Fisher relation \\citep[TF relation,][]{tully.fisher:new} evolves along with the galaxies that define it, reflecting general trends in mass assembly and star formation. Numerous studies have sought to trace the star formation history of disk galaxies via the redshift evolution of the zero point of the TF relation, under the assumption that TF zero point offsets represent luminosity evolution.  Awkwardly, some studies find minimal luminosity evolution to redshifts as high as $z\\sim1$ \\citep[e.g.,][]{forbes.phillips.ea:keck,vogt.phillips.ea:optical,bershady.haynes.ea:rotation}, while others report substantial 1.5--2 mag offsets at lower redshifts \\citep[e.g.,][]{rix.guhathakurta.ea:internal,simard.pritchet:internal}. Efforts to reconcile these results have generally invoked differential evolution, in which only low-luminosity galaxies evolve significantly \\citep[e.g.,][]{simard.pritchet:internal,ziegler.b-ohm.ea:evolution}.  However, some ``luminosity offsets'' may actually be velocity offsets.  In the TF relation, underestimated rotation velocities look exactly like enhanced luminosities.  The following three examples are particularly relevant to high $z$ studies.  (1) Optical emission-line data for high-$z$ ($z$ $\\sim$ 0.25--1 for this paper) blue compact galaxies may not extend to large enough radii to sample peak rotation velocities, based on 21 cm HI studies of analogous galaxies at low $z$ (\\citeauthor{barton.:possible} 2001; \\citeauthor{pisano.kobulnicky.ea:gas} 2001; see also \\citeauthor{kobulnicky.gebhardt:obtaining} 2000; \\citeauthor{courteau.sohn:galaxy} 2003).  Of course, high-$z$ studies must employ optical lines rather than HI to measure rotation velocities.  Also, most high-$z$ studies have selection biases favoring the bright galaxy cores and strong emission lines typical of blue compact galaxies.  One might hope that high-$z$ studies would be insensitive to radially truncated emission-line data, because unlike low-$z$ analyses, high-$z$ analyses usually derive rotation velocities by analyzing kinematic and photometric profiles together \\citep[e.g.,][]{vogt.forbes.ea:optical,simard.pritchet:analysis,ziegler.b-ohm.ea:evolution}. However, such modeling techniques typically rely on simplifying assumptions that blue compact galaxies probably routinely violate, such as the assumption of a close correspondence between emission-line and underlying disk-continuum fluxes, or the assumption that rotation curves can be simply parametrised based on exponential-disk fits to the spatial flux distribution regardless of disturbances or central mass concentrations.  (2) Low S/N can also cause radially truncated emission and underestimated rotation velocities, especially in samples already biased toward galaxies with centrally concentrated emission.  Accounting for S/N-induced rotation curve truncation could significantly reduce discrepancies between high-$z$ TF studies \\citep[as discussed for the Simard \\& Pritchet and Vogt et al.\\ studies by][]{kannappan:kinematic}.  (3) Distorted rotation curves may also yield unreliable rotation estimates and systematic velocity offsets.  In low-$z$ TF samples that contain interacting or morphologically peculiar galaxies, disturbances in longslit optical emission-line rotation curves clearly correlate with apparent luminosity boosts from the TF relation \\citep[][]{barton.geller.ea:tully-fisher,kannappan.fabricant.ea:physical}. We suspect that these apparent boosts may not be pure luminosity offsets, especially when they are larger than would be expected based on colors or H$\\alpha$ equivalent widths \\citep{kannappan.fabricant.ea:physical}.  At low $z$, large TF offsets associated with distorted or truncated rotation curves are most common for low-luminosity galaxies, and the affected galaxies often display emission-line S0 or irregular morphologies, sometimes with independent evidence of interactions (\\citeauthor{kannappan.fabricant.ea:physical} 2002; see also \\citeauthor{kobulnicky.gebhardt:obtaining} 2000; \\citeauthor{barton.geller.ea:tully-fisher} 2001).  Although most low-$z$ TF studies would reject such galaxies \\citep[e.g.,][]{courteau:optical,haynes.giovanelli.ea:i-band,tully.pierce:distances}, all of the high-$z$ studies that report substantial faint-end luminosity evolution employ selection criteria that would admit them \\citep[e.g.,][]{rix.guhathakurta.ea:internal,simard.pritchet:internal,mall-en-ornelas.lilly.ea:internal,ziegler.b-ohm.ea:evolution}. Furthermore, the frequency of such galaxies may be enhanced in high-$z$ samples to the extent that the interaction rate increases with $z$ \\citep{patton.pritchet.ea:dynamically,murali.katz.ea:growth}. These points raise the obvious concern that high-$z$ TF samples may contain a population of galaxies whose velocity offsets mimic differential luminosity evolution.  Another key consideration in interpreting apparent luminosity offsets is the possibility of third-parameter dependence in TF residuals. Numerous studies have examined possible physical drivers of TF offsets, including morphology, surface brightness, gas content, environment, and color, for TF samples chosen by a variety of criteria \\citep[e.g.,][ and additional references therein]{roberts:twenty-one,rubin.burstein.ea:rotation,giraud:two-color,pierce.tully:distances,mould.han.ea:nonlinearity,pierce.tully:luminosity-line,sprayberry.bernstein.ea:mass-to-light,courteau.rix:maximal,mcgaugh.schombert.ea:baryonic,verheijen:ursa*1,barton.geller.ea:tully-fisher,kannappan.fabricant.ea:physical}. In a recent analysis, \\citet{kannappan.fabricant.ea:physical} demonstrate that TF residuals correlate strongly with star formation indicators --- color and EW(H$\\alpha$) --- in the Nearby Field Galaxy Survey \\citep[NFGS,][]{jansen.franx.ea:surface,kannappan.fabricant.ea:physical}, a statistically representative survey of all galaxy types with no bias against interacting, peculiar, or early-type galaxies.  The inclusion of such galaxies distinguishes the NFGS TF sample (and most high-$z$ TF samples) from the majority of low-$z$ TF samples, which restrict analysis to a limited range of morphologies that may show only weak correlations between color and TF residuals \\citep[e.g.,][ see Kannappan et al.\\ 2002 for further discussion]{courteau.rix:maximal}. However, the Ursa Major cluster sample of \\citet{verheijen.sancisi:ursa}, which approximates a volume-limited sample, shows a stronger color--TF residual correlation \\citep{verheijen:ursa*1,kannappan.fabricant.ea:physical}, and \\citet{bershady.haynes.ea:rotation} also find initial evidence for a color--TF residual correlation at high $z$.  The existence of this correlation implies that high-$z$ samples that differ in average color because of different selection criteria will also differ in average TF zero-point offset.  If high-$z$ galaxies follow the same color--TF residual relation the NFGS follows, then we can use this relation to correct high-$z$ TF offsets for any bias toward blue colors (or we can use the EW(H$\\alpha$)--TF residual relation to correct for any bias toward strong emission lines).  Moreover, once such biases are removed, we can compare the remaining zero-point offset with the luminosity evolution predicted by the color--TF residual relation (based on true differences in mean color between high and low $z$), in order to determine whether TF zero-point evolution includes not only luminosity evolution, but also additional evolution reflecting the growth of stellar-to-total mass fractions over cosmic time \\citep{kannappan.gillespie.ea:interpreting}.  Obtaining a well-defined color--TF residual (CTFR) relation and measuring evolutionary offsets reliably may require special attention to galaxies with distorted or radially truncated rotation curves.  In this paper, we demonstrate such an analysis at low $z$ using the Close Pairs Survey of \\citet{barton.geller.ea:tully-fisher}.  The interacting galaxies in this survey display luminosity enhancements and misleading velocity offsets much like high-$z$ galaxies, as previously shown by \\citet{barton.geller.ea:tully-fisher}.  However, at low $z$ we can use high-resolution kinematic data to identify problem rotation curves objectively, using quantitative tests of radial truncation and asymmetry of shape based on those introduced by \\citet{kannappan.fabricant.ea:physical}.  Without explicitly accounting for kinematic anomalies, \\citeauthor{barton.geller.ea:tully-fisher} could not decouple luminosity offsets from velocity offsets and found no statistically significant CTFR relation for the Close Pairs Survey.  We recover the CTFR relation for the Close Pairs Survey by eliminating galaxies with severely truncated or asymmetric rotation curves based on quantitative criteria, and by analyzing modestly asymmetric rotation curves with a robust velocity width measure that does not assume a functional form. Using these procedures, we find that the TF and CTFR relations for the Close Pairs Survey look very similar to the corresponding relations for the NFGS.  Furthermore, the tightness of the Close Pairs Survey CTFR relation suggests that if a similar relation holds at higher $z$, determining whether galaxies lie on or off its locus may serve as a way to distinguish reliable luminosity evolution from TF offsets associated with kinematic anomalies.  Below, we describe the Close Pairs Survey and the NFGS (\\S~2), as well as our analysis methods (\\S~3), including quantitative criteria for identifying strongly asymmetric or truncated rotation curves.  We then analyze the TF and CTFR relations for the Close Pairs Survey, with attention to kinematic anomalies, and compare the Close Pairs Survey relations to the corresponding NFGS relations (\\S~4).  In \\S~5 we examine the possible drivers of kinematic anomalies. We go on to consider the implications of our results for high-$z$ TF studies in \\S~6.  Finally, we summarize our conclusions in \\S~7. ",
        "conclusions": "We have demonstrated robust methods for measuring luminosity evolution in TF samples with a high frequency of rotation curve anomalies, such as might be expected at high redshift.  The Close Pairs Survey of \\citet{barton.geller.ea:tully-fisher} is ideal for this analysis, as a low-$z$ TF sample with high-quality data and many similarities to high-$z$ TF samples: optical emission-line rotation curves, morphology-blind selection, and a large number of interacting galaxies.  The Nearby Field Galaxy Survey \\citep[NFGS,][]{jansen.franx.ea:surface,kannappan.fabricant.ea:physical} offers a low-$z$ reference sample with similar features, but with a more typical number of interacting galaxies.  We have extended \\citeauthor{barton.geller.ea:tully-fisher}'s previous TF analysis of the Close Pairs Survey, which showed that both starbursts and kinematic disturbances can create apparent ``luminosity evolution'' for galaxies in interacting pairs, by our demonstration of methods for isolating potentially spurious luminosity offsets associated with severe kinematic anomalies from reliable luminosity offsets clearly linked to star formation.  The largest apparent luminosity offsets in the Close Pairs Survey TF relation correspond to galaxies with severe kinematic anomalies \\citep[asymmetric rotation curve shapes and/or radially truncated rotation curve extents, using objective measures adapted from][]{kannappan.fabricant.ea:physical}.  The pattern of these galaxies' TF offsets looks much like the differential luminosity evolution claimed in many high-$z$ studies, with the largest TF offsets at luminosities fainter than M$_{\\rm B}\\sim-21$.  Excluding the galaxies with asymmetric or truncated rotation curves, however, and adopting a robust velocity width measure insensitive to minor kinematic distortions, we find that the TF relations for the Close Pairs Survey and the NFGS are very similar, with no significant evidence for overall luminosity enhancement in paired galaxies relative to the general population.  Nonetheless, we do find evidence for luminosity enhancement when we compare the color--TF residual (CTFR) relations for the two surveys.  Two galaxies that are not objectively flagged as kinematically anomalous extend the CTFR relation to very blue colors and large luminosity offsets, apparently reflecting interaction-induced star formation.  Of course, kinematically anomalous galaxies probably experience luminosity boosts as well, and in fact two such galaxies are also exceptionally blue.  However, anomalous galaxies are typically outliers from the CTFR relation, and their TF offsets may include velocity offsets.  Unfortunately, it is presently impossible to separate luminosity and velocity offsets for these galaxies.  If, as expected from hierarchical merging scenarios, the galaxy interaction rate was higher in the past, kinematic anomalies may pose a serious problem for high-$z$ TF studies.  Severe anomalies are roughly twice as common in the Close Pairs Survey as in the NFGS ($\\sim$20\\% vs.\\ $\\sim$10\\% of galaxies brighter than $M_B=-18$). Galaxy interactions probably explain the 3--4 times higher rate of strong rotation curve asymmetries in the Close Pairs Survey compared to the NFGS.  The externally triggered gas inflow associated with interactions can also lead to centrally concentrated line emission and thereby rotation curve truncation, a problem that may be compounded by low S/N data.  However, our data are inconclusive as to the primary source of rotation curve truncation in the Close Pairs Survey. To the extent that gas inflow processes play a role at high $z$, some of the assumptions inherent in high-$z$ rotation-curve fitting techniques may break down, leading to artificially low velocity widths, as observed at low $z$.  The frequency of kinematic anomalies at high $z$ that would meet our criteria is presently unknown.  We have shown that TF outliers associated with kinematic anomalies in the Close Pairs Survey occupy the same part of TF parameter space as the galaxies responsible for TF slope evolution in some high-$z$ studies \\citep[e.g.,][]{simard.pritchet:internal,ziegler.b-ohm.ea:evolution}. These studies tend to have morphology-blind selection criteria that would include kinematically anomalous galaxies, which often have blue compact, emission-line S0, or peculiar morphologies.  In contrast, studies that show less slope evolution \\citep[e.g.,][]{vogt.phillips.ea:optical} tend to favor large disks and probably contain fewer kinematic anomalies.  We have also shown that the slope evolution in the high-$z$ FORS Deep Field sample \\citep{b-ohm.ziegler.ea:tully-fisher} is largely driven by ``low-quality'' data points, labeled as such by B\\\"ohm et al.\\ based on perturbations or limited radial extent in the rotation curves.  Using the FORS Deep Field data, we find that most or all of the TF evolution measured at high $z$ can be modeled as an overall $\\sim$0.3 mag luminosity offset at fixed slope, consistent with evolution along the CTFR relation, plus a differential evolution component associated with kinematically anomalous galaxies, which show offsets as large as $\\sim$2 mag at low luminosities but add only a small $\\sim$0.2 mag enhancement to the total TF offset for the survey.  We note that the use of the outdated \\citet{pierce.tully:luminosity-line} TF calibration \\citep[superseded by][]{tully.pierce:distances} as a low-$z$ reference relation contributes 0.3--0.4 mag of spurious luminosity evolution to many high-$z$ TF studies.  At present, only the $\\sim$0.3 mag offset consistent with the CTFR relation can be reliably interpreted as luminosity evolution.  TF slope evolution associated with kinematic anomalies may be interesting for its own sake, as a source of data on mass-dependent evolution in the frequency of mergers and interactions (or the frequency of gas-dynamical disturbances caused by these events).  Consistent with mass-dependent evolutionary trends in star formation histories and the luminosity function \\citep[e.g.,][]{cowie.songaila.ea:new}, kinematically anomalous galaxies in the NFGS tend to be dwarf galaxies, fainter than M$_{\\rm B}=-18$, while analogous galaxies at high $z$ can be as bright as M$_{\\rm B}\\sim-21$.  We have considered four strategies for isolating reliable luminosity offsets from offsets possibly associated with kinematic anomalies at high $z$: identification of anomalies based on rotation curve properties, identification based on morphology, identification based on the color--TF residual relation, and inclusion of anomalies in optimally matched low-$z$ calibration samples that reproduce the distribution of anomalies expected at high $z$ (as well as high-$z$ selection criteria, data quality, and analysis techniques).  {\\em The color--TF residual relation may offer the simplest and most powerful tool currently available for measuring luminosity evolution independent of kinematic anomalies at high $z$,} especially when combined with optimal low-$z$ calibration samples.  Unreliable TF offsets associated with kinematic anomalies are typically CTFR outliers.  Conversely, reliable luminosity enhancements lie on the CTFR relation and extend it toward bluer colors.  Preliminary evidence for a CTFR relation at high $z$ has already been reported \\citep{bershady.haynes.ea:rotation}.  If the high-$z$ CTFR relation proves as tight as the Close Pairs Survey CTFR relation, then identifying CTFR outliers will serve as the preferred method for isolating kinematic anomalies in studies of luminosity evolution.  Once established, the high-$z$ CTFR relation may be applied to measuring not only luminosity evolution but also the evolution of stellar populations and stellar-to-total mass fractions \\citep[][]{kannappan.gillespie.ea:interpreting}.  The CTFR relation and the analogous relations for EW(H$\\alpha$) and EW([OII]) can also be used to reconcile discrepancies between high-$z$ TF studies with different selection biases in color or emission-line strength \\citep{kannappan:kinematic,kannappan.gillespie.ea:interpreting,kannappan.fabricant.ea:calibrating}. Matching selection criteria at low and high $z$ is only the first step, however, because of the potential for luminosity-dependent evolution in the frequency of kinematic anomalies.  By combining well-matched low-$z$ calibration samples with careful modeling of kinematic anomalies and the CTFR relation, future high-$z$ TF studies should be able to properly account for the major uncertainties of existing studies and reach consensus on how galaxy luminosities have evolved over cosmic time."
    },
    "0402/astro-ph0402247_arXiv.txt": {
        "abstract": "We have analyzed data from two sets of calibration observations of the Moon made by the {\\it Chandra X-Ray Observatory}. In addition to obtaining a spectrum of the bright side that shows several distinct fluorescence lines, we also clearly detect time-variable soft X-ray emission, primarily \\ioka\\ and \\iolya, when viewing the optically dark side. The apparent dark-side brightness varied by at least an order of magnitude, up to $\\sim 2 \\times 10^{-6}$ phot s$^{-1}$ arcmin$^{-2}$ cm$^{-2}$ between 500 and 900 eV, which is comparable to the typical 3/4-keV-band background emission measured in the \\rosat\\ All-Sky Survey. The spectrum is also very similar to background spectra recorded by \\chandra\\ in low or moderate-brightness regions of the sky. Over a decade ago, \\rosat\\ also detected soft X-rays from the dark side of the Moon, which were tentatively ascribed to continuum emission from energetic solar wind electrons impacting the lunar surface. The \\chandra\\ observations, however, with their better spectral resolution, combined with contemporaneous measurements of solar-wind parameters, strongly favor charge transfer between highly charged solar-wind ions and neutral hydrogen in the Earth's geocorona as the mechanism for this emission.  We present a theoretical model of geocoronal emission and show that predicted spectra and intensities match the \\chandra\\ observations very well. We also model the closely related process of heliospheric charge transfer and estimate that the total charge transfer flux observed from Earth amounts to a significant fraction of the soft X-ray background, particularly in the \\rosat\\ 3/4-keV band. ",
        "introduction": "\\label{sec:intro}  As reported by \\citet{cit:schmitt1991}, an image of the Moon in soft X-rays (0.1--2 keV) was obtained by the {\\it R\\\"{o}ntgen Satellite} (\\rosat) using its Position-Sensitive Proportional Counter (PSPC) on 1990 June 29. This striking image showed an X-ray-bright sunlit half-circle on one side, and a much dimmer but not completely dark side outlined by a brighter surrounding diffuse X-ray background.  Several origins for the dark-side emission were considered, but the authors' favored explanation was continuum emission arising from solar wind electrons sweeping around to the unlit side and impacting on the lunar surface, producing thick-target bremsstrahlung. Given the very limited energy resolution of the PSPC, however, emission from multiple lines could not be ruled out.  A significant problem with the bremsstrahlung model was explaining how electrons from the general direction of the Sun could produce events on the opposite side of the Moon, with a spatial distribution which was ``consistent with the telescope-vignetted signal of a constant extended source.'' An elegant alternative explanation would be a source of X-ray emission {\\em between} the Earth and the Moon, but at the time, no such source could be envisioned. If this source were also time-variable, it would account for the Long Term Enhancements (LTEs) seen by \\rosat. These occasional increases in the counting rate of the PSPC are vignetted in the same way as sky-background X-rays, indicating an external origin \\citep{cit:snowden1995}. LTEs are distinct from the particle-induced background, and are uncorrelated with the spacecraft's orientation or position (geomagnetic latitude, etc.), although \\citet{cit:freyberg1994} noted that LTEs appeared to be related, by a then unknown mechanism, to geomagnetic storms and solar wind variations. The final \\rosat\\ All-Sky Survey (RASS) diffuse background maps \\citep{cit:snowden1995,cit:snowden1997} removed the LTEs, so far as possible, by comparing multiple observations of the same part of the sky, but any constant or slowly varying ($\\tau \\ga 1$ week) emission arising from whatever was causing the LTEs would remain.  A conceptual breakthrough came with the \\rosat\\ observation of comet Hyakutake \\citep{cit:lisse1996} and the suggestion by \\citet{cit:cravens1997} that charge transfer (CT) between the solar wind and neutral gas from the comet gave rise to the observed X-ray emission. In solar-wind charge transfer, a highly charged ion in the wind (usually oxygen or carbon) collides with neutral gas (mostly water vapor in the case of comets) and an electron is transferred from the neutral species into an excited energy level of the wind ion, which then decays and emits an X-ray. This hypothesis has been proven by subsequent observations of comets such as C/LINEAR 1999 S4 by \\chandra\\ \\citep{cit:lisse2001} and Hyakutake by {\\it EUVE} \\citep {cit:krasno2001} (see also the review by \\citet{cit:cravens2002}), and is supported by increasingly detailed spectral models \\citep{cit:khar2003,cit:khar2000}. A more extensive history of the evolution of the solar-wind CT concept can be found in \\citet{cit:cravens2001} and \\citet{cit:robertson2003}.  Citing the cometary emission model, \\citet{cit:cox1998} pointed out that CT must occur throughout the heliosphere as the solar wind interacts with atomic H and He within the solar system. \\citet{cit:freyberg1998} likewise presented \\rosat\\ High-Resolution Imager data that provided some evidence for a correlation between increases in the apparent intensity of comet Hyakutake and in the detector background; he further suggested that this could be caused by charge transfer of the solar wind with the Earth's atmosphere. A rough broad-band quantitative analysis by \\citet{cit:cravens2000} predicted that heliospheric emission, along with CT between the solar wind and neutral H in the Earth's tenuous outer atmosphere (geocorona), accounts for up to half of the observed soft X-ray background (SXRB). Intriguingly, results from the Wisconsin Soft X-Ray Background sky survey \\citep{cit:mccammon1990} and RASS observations \\citep{cit:snowden1995} indicate that roughly half of the 1/4-keV background comes from a ``local hot plasma.'' \\citet{cit:cravens2001} also modeled how variations in solar-wind density and speed should affect heliospheric and geocoronal CT emission observed at Earth, and found strong correlations between the measured solar-wind proton flux and temporal variations in the \\rosat\\ counting rate.  In this paper we present definitive spectral evidence for geocoronal CT X-ray emission, obtained in \\chandra\\ observations of the Moon. Data analysis is discussed in \\S\\ref{sec:data}, and results are presented in \\S\\ref{sec:results}. As we show in \\S\\ref{sec:interp}, model predictions of geocoronal CT agree very well with the observed \\chandra\\ spectra. In \\S\\ref{sec:sxrb} we estimate the level of heliospheric CT emission, discuss the overall contribution of CT emission to the SXRB, and assess the observational prospects for improving our understanding of this subtle but ubiquitous souce of X-rays. ",
        "conclusions": "\\label{sec:conclusions}  As described in this paper, we have detected significant time-variable soft X-ray emission in \\chandra\\ observations of the dark side of the Moon which is well explained by our model of geocoronal charge transfer. The observed brightness ranged from a maximum of $\\sim 2 \\times 10^{-6}$ phot s$^{-1}$ arcmin$^{-2}$ cm$^{-2}$, with most of the emission between 500 and 900 eV, to a minimum at least an order of magnitude lower. Predicted intensities, which are based in part on detailed solar wind data, match observation to within a factor of two, which is within the model uncertainty. Emission from \\ioka, \\iolya, and a blend of high-$n$ \\ion{O}{8} Lyman lines is detected with high confidence, as well as probably \\imgka\\ and perhaps high-$n$ emission from \\ion{C}{6}. We also include estimates of heliospheric emission and find that the total charge transfer emission amounts to a substantial fraction of the soft X-ray background, roughly one-third of the rate measured in the \\rosat\\ R45 band. Future observations with microcalorimeter detectors should allow much more accurate assessments of the contribution of charge transfer emission to the SXRB because of its unique spectral signatures."
    },
    "0402/astro-ph0402592_arXiv.txt": {
        "abstract": "We discuss the main properties of the Galactic globular cluster (GC) blue straggler stars (BSS), as inferred from our new catalog containing nearly 3000 BSS. The catalog has been extracted from the photometrically homogeneous $V$ vs. ($B-V$) color-magnitude diagrams (CMD) of 56 GCs, based on WFPC2 images of their central cores.  In our analysis we used consistent relative distances based on the same photometry and calibration.  The number of BSS has been normalized to obtain relative frequencies ($F_{\\rm BSS}$) and specific densities ($N_{\\rm S}$) using different stellar populations extracted from the CMD. The cluster $F_{\\rm BSS}$ is significantly smaller than the relative frequency of field BSS.  We find a significant anti-correlation between the BSS relative frequency in a cluster and its total absolute luminosity (mass).  There is no statistically significant trend between the BSS frequency and the expected collision rate. $F_{\\rm BSS}$ does not depend on other cluster parameters, apart from a mild dependence on the central density.  PCC clusters act like normal clusters as far as the BSS frequency is concerned.  We also show that the BSS luminosity function for the most luminous clusters is significantly different, with a brighter peak and extending to brighter luminosities than in the less luminous clusters.  These results imply that the efficiency of BSS production mechanisms and their relative importance vary with the cluster mass. ",
        "introduction": "Globular Clusters (GCs) are important astrophysical laboratories for investigating the stellar dynamics and stellar evolution of low-mass stars (e.g., Meylan \\& Heggie 1997).  In recent years, it became clear that we can not study these two astrophysical processes independently if we want to understand GCs and properly address several long-standing problems concerning their stellar content.  Among the most puzzling products of the interplay between stellar evolution and dynamics are the blue straggler stars (BSS). This group of stars was originally identified by Sandage (1953) in the cluster M3 as a bluer and brighter extension of the main sequence (MS) turn-off (TO) stars.  At present, the most popular mechanisms suggested to account for their origin are {\\em primordial binary evolution} (McCrea 1964), i.e., mass transfer and/or coalescence in primordial binary systems (Carney et al.\\ 2001), and {\\em collisional merging}, i.e., the collision of single and/or binary systems (Bailyn 1995).  Unfortunately, current photometric investigations do not allow us to figure out the mechanism that triggers the formation of BSS in GCs, and indeed it has been suggested that both primordial binary evolution and collisions are probably at work in different clusters (Ferraro, Fusi Pecci, \\& Bellazzini 1995; Piotto et al.\\ 1999, Ferraro et al. 2003), or even within the same cluster (Ferraro et al.\\ 1997).  The observational scenario concerning BSS formation has been recently enriched by the results of a spectroscopic survey by Preston \\& Sneden (2000, hereafter PS00). On the basis of multi-epoch radial velocity data of field blue metal-poor (BMP) stars, PS00 found that more than 60\\% of the stars in their sample are binaries.  On the basis of empirical evidence, PS00 concluded that at least 50\\% of BMP stars are BSS. Moreover, PS00 suggested that the BSS in their sample must have formed via mass transfer in binaries. Finally, PS00 found that the specific frequency of BSS in the local halo is an order of magnitude larger than in GCs.  This discrepancy opens several new questions concerning the origin of field and cluster BSS.  In an attempt to better understand the properties of BSS stars in GCs, we took advantage of our homogeneous database of color-magnitude diagrams (CMD) from WFPC2 images (Piotto et al.\\ 2002) to select a sample of nearly 3000 BSS in 56 GCs characterized by different morphological and dynamical properties. In this paper we exploit the new BSS catalog to investigate empirically whether the BSS population is related to any of the properties of the parent GC. Here we present the results we believe to be the most relevant and original.  The entire BSS photometric catalog and further details on the BSS extraction will be published in a forthcoming paper (De Angeli et al.\\ 2004), and it will become available at the Padova Globular Cluster Group web site. ",
        "conclusions": "\\label{conc}  The statistically significant anticorrelation of the BSS relative frequency with the integrated luminosity and the independence of the expected collision rate discussed in the previous Section are noteworthy, and we will concentrate on them.  These observational facts are complemented by the finding by PS00 that field BSS have a frequency $F_{\\rm BSS}=4.0$, an order of magnitude larger than the BSS frequency of the bulk of the GCs.  \\begin{figure} \\plotone{f2.ps} \\figcaption{Number of BSS (circles), HB (triangles), and RGB (squares) stars per absolute visual flux unit as a function of the integrated cluster magnitude ({\\it top panel}) and of the collision rate ({\\it lower panel}). Open circles represent PCC clusters.} \\end{figure}  Figure 2 shows the same results of Fig.\\ 1, in a different way that may be more enlightening.  Here we look at the number of BSS, HB, and RGB stars relative to the total flux in the same region.  We call this quantity $N_{\\rm S}$; it is defined by $$\\log N_{\\rm S}=\\log\\left[{N \\over (F_{\\rm HST}/10^{-0.4 V_{\\rm tot}})} \\cdot {1 \\over 10^{-0.4 M_V}}\\right],$$ where $N$ is the total number of BSS (or HB or RGB) stars, corrected for completeness, $F_{\\rm HST}$ is the total flux from all the stars that we measured in the region, $V_{\\rm tot}$ is the integrated apparent magnitude of the cluster, and $M_V$ is its integrated absolute magnitude.  (Note that our CMDs typically extend well below the turnoff, so that the contribution from the fainter stars is negligible.)  The first factor in the brackets can be understood as follows: the quantity $F_{\\rm HST}/10^{-0.4 V_{\\rm tot}}$ is the fraction of the total cluster flux that is sampled by the HST field.  If the BSS were distributed like the flux, then $N / (F_{\\rm HST}/10^{-0.4 V_{\\rm tot}})$ would be the number of BSS we would expect if we could observe the whole cluster. In fact, as the BSS are more concentrated to the center than is the flux, the quantity above is still a reasonable approximation to the total number of BSS to be expected.  The defect in the approximation increases the scatter in $N_{\\rm s}$, but it does not introduce any systematic effects, since (as we have verified) there is no correlation between the fraction of flux included and $M_V$. The second factor in the brackets is just our previous normalization to the size of the cluster, but now in luminosity units.  Interestingly enough, Fig.\\ 2 confirms that the HB and RGB stars are very good tracers of the cluster population, as their absolute density remains constant over more than 4 magnitudes in cluster total luminosity. This fact removes the risk that the results of Fig.\\ 1 might be due to some anomalous gradient in the distribution of HB and RGB stars (cf.\\ Djorgovski, Piotto \\& Capaccioli 1993). Figure\\ 2 confirms that the density of BSS decreases with increasing total cluster mass and that there is no correlation between the density of BSS and the collisional parameter. However, we note that, given the small size of the error bars, the dispersion of the the BSS density is much larger than the dispersion of the HB and RGB star densities, and that, as noticed in Fig.~1, clusters with $\\Gamma_\\star<10^{-15}$ have a 2--3 times larger BSS density than clusters with higher collision rates. The lack of an overall dependence of $F_{\\rm BSS}$ and $N_s$ on the collisional parameter seems to suggest that direct collisions of single or binary stars are not the main formation mechanism of BSS. At first glance, the evolution of primordial binaries also does not seem to be the dominant formation mechanism for BSS in all GCs. In the simple hypothesis that the binary fraction is the same in all clusters, we would expect the BSS density to show a behavior similar to that of the HB and RGB stars, in Fig.\\ 2.  On the other hand, the evolution of primordial binaries is affected by the cluster environment, and, in particular, it is accelerated in clusters where the encounter probability is higher.  Indeed, in a paper parallel to this one, using the mechanism proposed by Davies \\& Hansen (1998) to explain the production of millisecond pulsars in GCs, Davies, Piotto, \\& De Angeli (2004) demonstrate that in clusters with high encounter probability the formation of BSS from primordial binaries has been favored in the past.  Now these binaries cannot form BSS anymore (they have already evolved), and this explains the observed relative absence of BSS in many high mass, high collision rate clusters.  It also explains the relatively larger fraction of BSS among the field stars, where the even lower-density environment makes the evolution of binaries via encounters slower than in any GC, allowing them to produce BSS for a more extended time interval (till the present).  Davies et al.\\ (2004) show also that only in the most luminous GCs (specifically, clusters with $M_V<-8.8$) do the BSS start to be produced predominantly by stellar collisions.  A better way to characterize the physical properties of the BSS is to look at their luminosity function (LF). In order to overcome possible dependencies of the LF on the cluster metallicity, distance, and reddening, we have divided the luminosity of each BSS by the turn-off luminosity of the parent cluster. Figure 3 shows the LFs for GCs with different total luminosity. The cut in $M_V$ has been set at $M_V=-8.8$, where the theory (Davies et al.\\ 2004) predicts that the BSS should become predominantly collisional.  \\begin{figure} \\plotone{f3.ps} \\figcaption{BSS LFs for clusters with different integrated magnitude.} \\end{figure}  Interestingly enough, clusters with $M_V<-8,8$ have a BSS LF which is significantly different from the BSS LF of less luminous clusters (Fig.~3), in that the LFs for the most luminous clusters have a brighter peak and are significantly shifted toward brighter magnitudes.  If the relative importance of the BSS production mechanisms depends on the cluster mass, we would then expect to see a dependence of the BSS LF on $M_V$, as is observed in Fig.\\ 3. In general, a BSS produced by collision is expected to have a different luminosity with respect to a BSS from mass transfer or merger of binaries, due to the resulting interior chemical profile. How much different is still controversial.  Indeed, recent detailed smoothed particle hydrodynamic simulations performed by Sills et al.\\ (2002) have shown that collision products are not chemically homogeneous. This has the effect of producing a BSS structure less blue and less bright than expexted on the basis of the \\lq{fully mixed}\\rq\\ models (e.g., Bailyn \\& Pinsonneault 1995). Nevertheless, Sills et al.\\ (2001) have also shown that collision products emerge as rapidly rotating blue stragglers, and so far we lack a full understanding of the changes in the evolutionary properties due to rotationally induced mixing.  It is also worth noting that PCC clusters seem to have normal BSS population. This might be due to the fact that the core collapse phase is very short, and confined to the very central part of the clusters, and therefore does not affect the BSS production over the last few Gyr."
    },
    "0402/astro-ph0402071_arXiv.txt": {
        "abstract": "{Analysis of the low resolution UV(IUE) spectra of 15 hot post-AGB candidates is presented. The UV(IUE) spectra of 10 stars suggest partial obscuration of the hot stars due to circumstellar dust. The reddened continua of these 10 stars were used to model and estimate the circumstellar extinction. The circumstellar extinction law was found to be linear in $\\lambda^{-1}$ in the case of IRAS13266-5551 (CPD-55 5588), IRAS14331-6435 (Hen3-1013), IRAS16206-5956 (SAO 243756), IRAS17074-1845 (Hen3-1347), IRAS17311-4924 (Hen3-1428), IRAS18023-3409 (LSS 4634), IRAS18062+2410 (SAO 85766), IRAS18371-3159 (LSE 63), IRAS22023+5249 (LSIII +5224) and IRAS22495+5134 (LSIII +5142). There seems to be no significant circumstellar extinction in the case of IRAS17203-1534, IRAS17460-3114 (SAO 209306) and IRAS18379-1707 (LSS 5112). The UV(IUE) spectrum of IRAS12584-4837 (Hen3-847) shows several emission lines including that of HeII. It may be a massive young OB-supergiant or a low mass star in the post-AGB phase of evolution.  IRAS16206-5956 (SAO 243756) and IRAS 18062+2410 (SAO 85766) show variability in the UV which in addition to stellar pulsations may be attributed to a dusty torus in motion around the hot central stars. The UV spectrum of the bipolar PPN, IRAS17423-1755 (Hen3-1475) indicates that the central B-type star is obscured by a dusty disk. The stars were placed on the log g$-$log T$_{\\rm eff}$ diagram showing the post-AGB evolutionary tracks of Sch\\\"onberner. Terminal wind velocities of the stars were estimated from the CIV and NV stellar wind features. The presence of stellar wind in some of these stars indicates ongoing mass-loss. ",
        "introduction": "Low and intermediate mass stars (M $\\simeq$ 0.8 - 8 M$_{\\odot}$) pass through the post-asymptotic giant branch (post-AGB) phase of evolution on their way to becoming planetary nebulae (PNe). From an analysis of the Infrared Astronomical Satellite Point Source Catalog (IRAS PSC) cooler post-AGB stars having G,F,A supergiant like character were first identified (Parthasarathy \\& Pottasch 1986, Lamers et al. 1986, Pottasch \\& Parthasarathy 1988a, Hrivnak et al. 1989). These stars were  found to have circumstellar dust shells with far-IR colors and flux distributions similar to the dust shells of PNe. Later, from an analysis of IRAS data, Parthasarathy \\& Pottasch (1989) found a few hot (OB spectral types) post-AGB candidates. Their supergiant like character, the presence of cold detached dust shells, far-IR colors similar to PNe and high galactic latitudes suggested that they may be in a post-AGB phase of evolution. Thus, there seems to be an evolutionary sequence ranging from the cooler G,F,A supergiant-like stars to hotter O-B types, evolving from the tip of the AGB towards young PN stage (Parthasarathy, 1993c).  Pottasch et al. (1988b) and van der Veen \\& Habing (1988) identified a region of the IRAS color-color diagram (F(12$\\mu$)/F(25$\\mu$) $<$ 0.35 and F(25$\\mu$)/F(60$\\mu$) $>$ 0.3) which was mainly populated by stars in transition from the AGB to the PN phase. Based on their far-IR colors and low resolution optical spectra, several hot post-AGB candidates were identified (Parthasarathy \\& Pottasch 1989, Parthasarathy 1993a, 1993c, Parthasarathy et al., 2000a). The optical spectra of these objects show strong Balmer emission lines and in some cases low excitation nebular emission lines such as [NII] and [SII] superposed on the OB stellar continuum. The absence of [OIII] 5007\\AA~ line and the presence of low excitation nebular emission lines indicate that photoionisation has just started. It is important to study these stars in the UV to obtain better estimates of their temperatures and to look for signatures of circumstellar reddening, mass-loss and stellar winds. The UV(IUE) spectra of some hot post-AGB stars (eg. Hen3-1357, Parthasarathy et al. 1993b, 1995, Feibelman, 1995) have revealed violet shifted stellar wind P-Cygni profiles of CIV, SiIV and NV, indicating hot and fast stellar wind, post-AGB mass-loss and rapid evolution. In this paper we have analysed the UV(IUE) spectra of 15 hot post-AGB candidates. ",
        "conclusions": "We analysed the UV(IUE) spectra of 15 hot post-AGB candidates. In 11 cases (IRAS13266-5551 (CPD-55 5588), IRAS14331-6435 (Hen3-1013), IRAS16206-5956 (SAO 243756), IRAS17074-1845 (Hen3-1347), IRAS17311-4924 (Hen3-1428), IRAS17423-1755 (Hen3-1475), IRAS18023-3409 (LSS 4634), IRAS18062+2410 (SAO 85766), IRAS18371-3159 (LSE 63), IRAS22023+5249 (LSIII +5224) and IRAS22495+5134 (LSIII +5142)), the UV spectra revealed obscuration of the hot central stars due to circumstellar dust. While IRAS17423-1755 (Hen3-1475) was not detected at all in a 35 minute exposure, the UV continua of the remaining 10 stars were found to be considerably reddened. We found that the circumstellar extinction in these 10 stars varies linearly as $\\lambda^{-1}$. A $\\lambda^{-1}$ law for the circumstellar extinction was also found in the case of the post-AGB star, HR4049 (Waters et al., 1989, Monier \\& Parthasarathy, 1999). In the context of Mie scattering (Spitzer, 1978), linear extinction arises from dust grains small compared to the wavelength of light. The shortest wavelength of light at which the extinction is linear can give an estimate of the size of the smallest grains in the circumstellar environment of these stars ($\\lambda$ = 2$\\pi$a, where, a is the radius of the dust grain). However our IUE SWP observations are limited to 1150\\AA~ ($\\lambda^{-1}$ $\\approx$ 8.7$\\mu^{-1}$). Shortward of 1300\\AA~ the spectra are noisy and often contaminated by Lyman $\\alpha$. Taking 1300\\AA~ as the shortest observed wavelength at which the extinction is linear in $\\lambda^{-1}$, we may infer an upper limit of ~a $\\approx$ 200\\AA~ for the radii of the small grains. Waters et al. (1989) speculate that the destruction of these grains in the vicinity of the hot central stars of PPNe and PNe may give rise to smaller grains and polyaromatic hydrocarbons (PAHs). PAH features at 8.2, 8.6 and 11.3 $\\mu$ have been detected in the circumstellar environment of several post-AGB stars, PPNe and PNe (see eg. Beintema et al., 1996). It would be interesting to study the infrared spectra of our hot post-AGB candidates to know more about the chemical compositions (carbon-rich or oxygen-rich nature) of the dust grains and the evolution of these grains in the circumstellar environment of these stars. Variation of IRAS16206-5956 (SAO 243756) and IRAS18062+2410 (SAO 85766) in the UV may be due to stellar pulsations and/or due to variable circumstellar extinction similar to that observed in the case of HR4049 (Waters et al., 1989, Monier \\& Parthasarathy, 1999). Significant circumstellar extinction was not observed in the case of IRAS17203-1534, IRAS17460-3114 (SAO 209306) and IRAS18379-1707 (LSS 5112). The effective temperatures and gravities of these three stars were estimated using Kurucz model atmospheres.  F$_{\\rm fir}$/F$_{\\rm star}$ $>>$ 1.0 in the case of IRAS14331-6435 (Hen3-1013), IRAS17311-4924 (Hen3-1428), IRAS17423-1755 (Hen3-1475), IRAS18062+2410 (SAO 85766) and IRAS22023+5249 (LSIII +5224) indicates the presence of dusty disks around these stars. From the UV(IUE) spectra we found that 7 (IRAS12584-4837, IRAS13266-5551 (CPD-55 5588), IRAS17203-1534, IRAS17311-4924 (Hen3-1428), IRAS17460-3114 (SAO 209306), IRAS18023-3409 (LSS 4634) and IRAS22023+5249 (LSIII +5224)) of the 15 hot post-AGB candidates have stellar wind velocities in excess of 1000kms$^{-1}$ indicating post-AGB mass-loss."
    },
    "0402/hep-ph0402128_arXiv.txt": {
        "abstract": "After recent results from solar neutrino experiments and KamLAND we can definitely say that neutrinos from SN1987A underwent flavor conversion, and the conversion effects must be taken into account in the analysis of the data.  Assuming the normal mass hierarchy of neutrinos we calculate the permutation factors $p$ for the Kamiokande-2, IMB and Baksan detectors.  The conversion inside the star leads to $p = 0.28 - 0.32$; the oscillations in the matter of the Earth give partial (and different for different detectors) regeneration of the original $\\bar{\\nu}_e$ signal, reducing this factor down to 0.15 - 0.20 (at $E = 40$ MeV). We study in details the influence of conversion on the observed signal depending on the parameters of the original neutrino spectra.  For a given set of these parameters, the conversion could lead to an increase of the average energy of the observed events up to 50\\% and of the number of events by a factor of 2 at Kamiokande-2 and by a factor of 3 - 5 at IMB. Inversely, we find that neglecting the conversion effects can lead up to 50\\% error in the determination of the average energy of the original $\\bar{\\nu}_e$ spectrum and about 50\\% error in the original luminosity.  Comparing our calculations with experimental data we conclude that the Kamiokande-2 data alone do not favor strong conversion effect, which testifies for small difference of the original $\\bar{\\nu}_e$ and $\\bar{\\nu}_\\mu$ spectra. In contrast, the combined analysis of the Kamiokande and IMB results slightly favors strong conversion effects (that is, large difference of the original spectra). In comparison with the no oscillation case, the latter requires lower average energy and higher luminosity of the original $\\bar{\\nu}_e$ flux. % ",
        "introduction": "After recent results from solar neutrino experiments, and first of all SNO~\\cite{Ahmad:2001an,Ahmad:2002jz,Ahmad:2002ka}, as well as from the reactor experiment KamLAND~\\cite{Eguchi:2002dm}, which have selected the Large Mixing Angle (LMA) MSW solution of the solar neutrino problem, we can definitely say that neutrinos from SN1987A got converted. The neutrino flavor transformations influenced the signals observed in 1987 by Kamiokande-2 (K2)\\cite{Hirata:1987hu,Hirata:1988ad}, IMB \\cite{Bionta:1987qt,Bratton:1988ww} and Baksan \\cite{Alekseev:1987ej}. Conversion effects must be taken into account in the analysis of the data and in the determination of the properties of the original neutrino fluxes.  Results obtained without conversion are to some extent incorrect.    The conversion of neutrinos associated to  SN1987A has been extensively studied  before (see \\cite{Wolfenstein:1987pj}-\\cite{Barger:2002px} for an incomplete list of references). Here we comment on some papers which are relevant for our present discussion.    It was suggested~\\cite{talkalexei} that the difference of signals detected by the K2 and IMB detectors can be explained partially by the oscillations of neutrinos in the matter of the Earth since the distances crossed by neutrinos on the way to these two detectors were different. The suggestion implied, however, a large lepton mixing, which was  not  a favored idea at that time. The detailed calculations have been done 13 years later~\\cite{Lunardini:2000sw}, when certain hints appeared that LMA could be the correct solution of the solar neutrino problem.  In connection to SN1987A, Wolfenstein considered antineutrino conversion in the star in the case of large mixing~\\cite{Wolfenstein:1987pj}. He concluded that conversion leads to a harder energy spectrum of the observed events and, possibly, to a larger number of events.  In the attempt to restrict the large lepton mixing, the conversion of antineutrinos in the non-resonance channel has been considered in details~\\cite{Smirnov:1994ku}. From the analysis of the SN1987A data a bound on the permutation factor ($p < 0.35$), and consequently,on the mixing angle has been obtained. It was found that the bound is weaker in the LMA range, and the Earth matter effect further relaxes it.  A detailed analysis of the SN1987A data based on the Poisson statistics has been performed by Jegerlehner et al. \\cite{Jegerlehner:1996kx}, who found that the data do not allow a definite conclusion on the oscillations hypothesis.  In the event that a large neutrino mixing is confirmed (as it has been recently), the data analysis would point toward average neutrino energies (at production in the star) lower than what theoretically predicted.  By combining solar and SN1987A data, the authors of refs. \\cite{Kachelriess:2001sg,Kachelriess:2000fe} concluded that the LMA region was the most suitable, among the large mixing solutions of the solar neutrino problem, to reconcile SN1987A data and predictions from numerical supernova codes, in agreement with \\cite{Smirnov:1994ku}.  In ref. \\cite{Lunardini:2000sw}, following the early suggestion~\\cite{talkalexei}, we have considered the possibility that certain features of the energy spectra of the events detected by K2 and IMB can be explained by neutrino oscillations in the matter of the Earth. This fixes several bands in the $\\Delta m^2 - \\cos 2\\theta$ plane. It was concluded that the data favor the parameters of LMA solution and the normal mass hierarchy.  The inverted mass hierarchy is disfavored, unless the $1-3$ mixing angle is very small \\cite{Minakata:2000rx} (see however \\cite{Barger:2002px}).  The combined analysis of results from solar neutrino experiments and KamLAND lead to the values of oscillation parameters \\be \\Delta m^2 = 7.1^{+3.2}_{-2.2} \\cdot 10^{-5} {\\rm eV}^2 , ~~~ \\tan^2 \\theta = 0.40 \\pm 0.10~, \\label{bfglob} \\ee which coincides with the third band (from the bottom in $\\Delta m^2$ scale) found in~\\cite{Lunardini:2000sw}.   In this paper we revisit the conversion of neutrinos from SN1987A using the latest information on neutrino mass spectrum and mixing.  We address the questions of how neutrinos were converted, how  conversion modified  the observed signals and what could be the error in the determination of the original neutrino fluxes if conversion is not taken into account.\\\\  The analysis of the SN1987A data and a determination of the original spectra as precise as possible are needed not only to understand what happened in 1987 but also to compare with the results of future detections of neutrino bursts from supernovae. Detections of supernova neutrinos are rare events and furthermore each supernova is unique. Indeed, the mass of the progenitor, luminosity, rotation, magnetic fields, chemical composition can be substantially different, and, as a consequence, the properties of the neutrino fluxes can vary. Thus, future high statistics detections are not expected to reproduce the same features as those of SN1987A, but will give somehow complementary information. The comparison of neutrino signals from different supernovae would be extremely important for understanding the latest stage of star evolution, the dynamics of core collapse and explosion.\\\\  The paper is organized as follows. In sec. 2 we consider conversion of antineutrinos and calculate the $\\bar{\\nu}_e$ survival probabilities and permutation factors for Kamiokande-2, IMB and Baksan. In sec. 3 the effects of conversion on the observed signals are studied depending on the parameters of the original spectra. In sec. 4 we compare the predictions with the experimental results and make some indicative conclusions on the properties of the original fluxes. The results are summarized in sec. 5. ",
        "conclusions": "1). After the identification of the solution of the solar neutrino problem and KamLAND results, we can definitely say that neutrinos from SN1987A underwent  flavor conversion inside the star and oscillations in the matter of the Earth.  In the assumption of the normal mass hierarchy, the conversion probabilities can be calculated with good precision. We find that the permutation factor is about $p= 0.28 - 0.32$ due to conversion inside the star and oscillations in the matter of the  Earth suppress the permutation. The Earth effect increases with energy, and   at $E \\sim 40$ MeV $p$ decreases down to 0.15- 0.20.\\\\  \\noindent 2). The conversion effects on observables depend strongly on the properties of the original neutrino fluxes, in particular, on the average energies and widths of the original spectra of $\\bar{\\nu}_e$ and $\\bar{\\nu}_{\\mu/\\tau}$. For a given set of these parameters, conversion leads to an increase of the number and of the average energy of the observed events as well as of the widths of the observed spectra. The conversion effects are different for K2 and IMB. By varying parameters in the ranges allowed by astrophysics, we have found that the average energy of events can increase by 30 - 50 \\% and the number of events by 50 \\% for K2 and by up to a factor of 3 for IMB.  {\\it Vice versa}, conversion changes the average energies and luminosities of the original neutrino fluxes extracted from the observations. In particular, it leads to a decrease of the original energy and (for fixed $E_{0e}$) of the luminosity of $\\bar{\\nu}_e$. We find that the decrease can be up to $\\sim 50-70\\%$. This leads to uncertainty in the determination of parameters of the original spectra.\\\\  \\noindent 3). Comparing calculations with the real signals from SN1987A we find that the K2 data alone do not show significant conversion effects, and can be well described by original fluxes with small conversion effects. This would testify for small difference of the original $\\bar{\\nu}_e$ and $\\bar{\\nu}_{\\mu}$ spectra. At the same time the K2 data do not exclude strong conversion. In this case, however, the original spectra should have lower energies and higher luminosities in comparison with the no oscillation case.  The characteristics of the original spectra extracted from the K2 and IMB data exhibit substantial differences.  These can be partially reduced by conversion effects: that is, the conversion improves the combined fit of the K2 and IMB data. The improvement is not dramatic, though, and does not lead to substantially more coherent overall picture.  It requires strongly different original spectra, $r_E = 1.5 - 2$, low average energy of the $\\bar{\\nu}_e$-spectrum, $E_{0e}=6.5 - 8.5$ MeV, and high luminosity: $L_e = (8 - 12) \\cdot 10^{52}$ ergs.    The oscillations in the matter of the Earth substantially modify the high energy parts of the spectra of events at K2 and IMB.   Our conclusions are valid for normal mass hierarchy or inverted hierarchy provided that $\\theta_{13}$ is negligibly small.  If the hierarchy is inverted and $\\sin^2 \\theta_{13}\\gta few \\cdot 10^{-4} $, the $\\barnue \\leftrightarrow \\barnux$ permutation in the star is complete \\cite{Dighe:1999bi,Lunardini:2003eh}, so that the $\\barnue$ flux at Earth is entirely due to the original $\\numu/\\nutau$ flux. It follows that the parameters extracted from the data analysis refer to the non-electron flavors produced inside the star and therefore would lead to a completely different test of supernova theory, with respect to the non-oscillation case.  In this same scenario of mixing and hierarchy, the amount of $\\barnue \\leftrightarrow \\barnux$ permutation could change at late times, as the supernova shock-wave reaches the external layers of the star, thus modifying the adiabaticity character of the $\\theta_{13}$-induced MSW resonance \\cite{Schirato:2002tg,Takahashi:2002yj,Lunardini:2003eh,Fogli:2003dw}. The effect on the time integrated neutrino signal is however small \\cite{Takahashi:2002yj,Fogli:2003dw} and negligible with respect to the large uncertainties of statistical and astrophysical nature on the SN1987A data.\\\\   \\noindent 4). In general, one needs to perform an analysis which employs the whole information contained in the data: both integral and differential (energy spectrum, arrival time) as well as errors in the determination of the energies of events, background and angular information. In view of the small number of detected events, the optimal type of analysis would be along the lines of the work of Loredo and Lamb \\cite{Loredo:2001rx}, with the conversion effects taken into account. The present study will allow to better understand the results of such a global fit. Our conclusions concerning the interpretation of the observed signals have qualitative character only.\\\\  \\noindent 5). It would be important to compare these results on SN1987A with those of future SN neutrino detections.  The latter will have high statistics and therefore will provide the possibility to disentangle the oscillation effects and the properties of the original fluxes.  We mark that many of the results presented here have general character and therefore will apply to future data as well. In particular, the effects of conversion in the star will be valid.  At the level of probabilities (permutation factor), the effects of oscillations in the Earth will probably be different, due to different trajectory of the neutrinos in the Earth. However the difference with respect to SN1987A may be negligible in the observed energy spectra if the energy resolution of the detector is larger than the size of the spectral modulations due to oscillations. The analysis of future supernova data will result in a better understanding of the generation of neutrino fluxes and therefore in more reliable predictions of fluxes also for SN1987A. In this way a more precise interpretation of the SN1987A signals can be done.   \\subsection*"
    },
    "0402/astro-ph0402651_arXiv.txt": {
        "abstract": "We have conducted a study of optical and \\HI properties of spiral galaxies (size, luminosity, \\Halpha flux distribution, circular velocity, \\HI gas mass) to investigate causes (\\eg nature versus nurture) for variation within the cluster environment. We find \\HI deficient cluster galaxies to be offset in Fundamental Plane space, with disk scale lengths decreased by a factor of $25$\\%.  This may be a relic of early galaxy formation, caused by the disk coalescing out of a smaller, denser halo (\\eg higher concentration index) or by truncation of the hot gas envelope due to the enhanced local density of neighbors, though we cannot completely rule out the effect of the gas stripping process.  The spatial extent of \\Halpha flux and the \\Bb radius also decreases, but only in early type spirals, suggesting that gas removal is less efficient within steeper potential wells (or that stripped late type spirals are quickly rendered unrecognizable).  We find no significant trend in stellar mass-to-light ratios or circular velocities with \\HI gas content, morphological type, or clustercentric radius, for star forming spiral galaxies throughout the clusters. These data support the findings of a companion paper that gas stripping promotes a rapid truncation of star formation across the disk, and could be interpreted as weak support for dark matter domination over baryons in the inner regions of spiral galaxies. ",
        "introduction": "Recent years have seen the emergence of a standard model for the growth of structure -- the hierarchical clustering model -- in which the gravitational effects of dark matter drive the evolution of structure from the near-uniform recombination epoch until the present day.  Simple models for galaxy formation in the context of these CDM cosmogonies have been remarkably successful in reproducing the properties of the local galaxy population (Kauffmann, White, \\& Guiderdoni 1993; Cole \\etal 1994) and have been extended to predict the sizes, surface densities and circular velocities of spiral galaxy disks (Dalcanton, Spergel, \\& Summers 1997; Mo, Mao, \\& White 1998).  The result has been a testable scenario that predicts the basic structural properties of the disk galaxy population in any specific cosmogony of CDM type, and how the ensemble of disk galaxies should evolve with redshift.  The models assume that most spirals formed as the central galaxies of isolated halos, an assumption supported by the fact that they must have undergone minimal dynamical disturbance since the formation of the bulk of the disk stars (T\\'oth \\& Ostriker 1992).  The size of the disk is thus expected to scale with that of its halo, and at high redshift the predicted distribution of halo sizes is shifted to smaller values.  By $z$ = 1 the predicted change in disk size is almost a factor of 3 for an Einstein-de Sitter Universe, though only about 1.5 for a flat universe with $\\Omega_0$ = 0.3.  High redshift field spirals show some evidence for this effect (Vogt 2000), though surface brightness selection biases can mimic this trend and it is difficult to separate the two factors (Vogt \\etal 2004c). Alternatively, one can search for evidence within the fossil record of local clusters.  Specifically, spiral galaxies which formed early in the vicinity of rich clusters may preserve signatures of early coalescence (Kauffmann 1995), even as we observe them today.  The difficulty, of course, is to distinguish between such cosmological variations in disk structure and in the direct effects of the cluster environment.  Some fundamental disk properties (\\eg scale length) may prove to be more a function of the cluster-wide environment than of the galaxy-specific environment, while the properties of the gas reservoir are strongly dependent upon the individual merger history and on interactions with the intracluster medium.  There have been a number of theoretical approaches to modeling the Tully-Fisher (Tully \\& Fisher 1977) relation and its implications for disk formation and evolution (Rhee 1996; McGaugh 2000; van den Bosch 2000 and references therein).  Its successful use to probe the peculiar velocity field about clusters is dependent on the assumption that the relevant galaxy properties do not vary significantly in different environments.  Ideally Tully-Fisher studies utilize many galaxies within a cluster to reduce the distance error estimate; by measuring velocity widths from \\Halpha rotation curves, \\HI selected samples can be augmented with gas-poor galaxies for which the \\HI line profile cannot furnish a velocity width (\\cf Giovanelli \\etal 1997a).  This increases the sampling of a cluster, and in addition \\HI stripped spirals are the most likely to be true cluster members rather than members of infalling groups, offset both on the sky and in projection.  Evidence for systematic changes in the Tully-Fisher relation as a function of environment could invalidate such studies.  Variations in the mass and light distribution of galaxies with environment could thus have a significant effect, as the Tully-Fisher relation might vary not only between the field and clusters, but, if environmental effects are important, from cluster to cluster depending upon its evolutionary state.  Furthermore, Salucci, Frenck, \\& Persic (1993) argue that the scatter in the Tully-Fisher relation can be decreased by decomposing the mass distribution within galaxies into a dark and a luminous component to isolate the contribution of the disk alone to the circular velocity of the system (see however Jablonka \\& Arimoto 1992), implicitly emphasizing the need for detailed knowledge of the environmental variations of the dark and luminous mass distributions.  Though most of the Tully-Fisher studies performed to date have not contained large or sufficiently well--defined samples to properly assess environmental impacts, there is some evidence (Biviano \\etal 1990) against a strong dependence of the Tully-Fisher relation calibration upon mean cluster densities, radial positions of galaxies within clusters, or galaxy morphologies.  There is also evidence that the Tully-Fisher relation and the $D - \\sigma$ relation, and other fundamental plane relations, do not necessarily produce the same results when applied in parallel.  The peculiar velocities derived by Aaronson \\etal (1986) for spirals, and by Lucey \\etal (1991) for ellipticals and S0s, within the cluster A2634, for example, differ significantly. The effects of the cluster environment upon the two populations may be a very relevant factor, as is the evidence for an environmental correlation in the zeropoint of the $D - \\sigma$ relation (Lucey \\etal 1991; Guzm\\'an 1993). However, the effects of infalling group velocities and, in this case, the possible inclusion of galaxies from the very nearby cluster A2666 (Scodeggio \\etal 1995) must also be examined.  Our program integrates both optical and \\HI observations; we seek to form a consistent picture of infalling spirals which is sensitive to both gas depletion and star formation suppression.  Our sample is made up of 329 galaxies, 296 selected from 18 nearby clusters and 33 isolated field galaxies observed for comparative purposes.  It extends over a wide range of environments, covering three orders of magnitude in cluster X-ray luminosity and containing galaxies located throughout the clusters from rich cores out to sparsely populated outer envelopes.  We have obtained \\Halpha rotation curves to trace the stellar disk kinematics within the potential at high resolution and to explore the strength of current star formation, \\HI line profiles to map the overall distribution and strength of \\HI gas, and \\Ib imaging to study the distribution of light in the underlying, older stellar population.  The sample contains spirals of all types, and is unbiased by the strength of flux from \\HII regions or by \\HI gas detection.  This paper is a companion to Vogt, Haynes, Herter \\& Giovanelli (2004a, \\pone), which details the observations and reduction of the data set, and to Vogt, Haynes, Herter \\& Giovanelli (2004b, \\ptwo) which explores the evidence for spiral galaxy infall. ",
        "conclusions": "We have conducted a study of optical and \\HI properties of spiral galaxies to explore the role of gas stripping as a driver of morphological evolution in clusters.  For \\HI deficient local cluster spirals (with less than 40\\% of the predicted atomic gas mass), we observe a decrease of a factor of $25$\\% in \\Ib disk scale lengths relative to \\HI normal spirals in spirals of all morphological types.  This may be a relic of early galaxy formation, caused by the disk coalescing out of a smaller, denser halo (\\eg higher concentration index) or reflect an environmental effect (\\eg truncation of the hot gas reservoir at large radius due to the high local density of neighboring galaxies), or it may be be the product of post disk-formation effects (\\eg gas stripping, tidal interactions).  The few \\HI normal spirals observed in the cores of rich clusters are similarly decremented, in support of the former.  \\Bb radii R$_b$ are decreased by 25\\% relative to \\HI normal field galaxies in \\HI deficient early type spirals (Sa through Sb).  They show an additional decrease of 55\\% in the extent of the \\Halpha flux along the disk, on or within the stripping radius predicted for the \\HI gas loss from models. Gas-rich {\\it asymmetric} spirals reflect this trend on the truncated side of the disk, implying that it is caused by the gas stripping process.  Late type spirals (Sbc through Scd) exhibit neither trend, suggesting that either gas removal is less effective within the potential or that once significant star formation suppression occurs these galaxies are no longer identified as late type spirals.  For star forming spiral galaxies associated with all of the clusters, we find no significant trend in stellar mass-to-light ratios or circular velocities with \\HI gas content, morphological type, or clustercentric radius.  This could be interpreted as support for dark matter domination over baryons well within the optical radius of disks.  In summary, we have explored the formation and evolution of spiral galaxies in local clusters through a combination of optical and \\HI properties.  We find evidence that the spirals within rich cluster cores, believed to coalesce from their halos at an early epoch, formed with intrinsically smaller disks than the local field population.  We have explored the relationship between \\HI gas stripping and the consequential suppression of young star formation in spiral galaxies.  We find that the ram pressure stripping and the suppression of star formation both occur quickly within spirals infalling into an intracluster medium of hot gas, but find little evidence for substantial mass stripping beyond that of the \\HI gas."
    },
    "0402/astro-ph0402184_arXiv.txt": {
        "abstract": "Gravitational waves from oscillating neutron stars in axial symmetry are studied performing numerical simulations in full general relativity. Neutron stars are modeled by a polytropic equation of state for simplicity. A gauge-invariant wave extraction method as well as a quadrupole formula are adopted for computation of gravitational waves. It is found that the gauge-invariant variables systematically contain numerical errors generated near the outer boundaries in the present axisymmetric computation. We clarify their origin, and illustrate it possible to eliminate the dominant part of the systematic errors. The best corrected waveforms for oscillating and rotating stars currently contain errors of magnitude $\\sim 10^{-3}$ in the local wave zone. Comparing the waveforms obtained by the gauge-invariant technique with those by the quadrupole formula, it is shown that the quadrupole formula yields approximate gravitational waveforms besides a systematic underestimation of the amplitude of $O(M/R)$ where $M$ and $R$ denote the mass and the radius of neutron stars. However, the wave phase and modulation of the amplitude can be computed accurately. This indicates that the quadrupole formula is a useful tool for studying gravitational waves from rotating stellar core collapse to a neutron star in fully general relativistic simulations. Properties of the gravitational waveforms from the oscillating and rigidly rotating neutron stars are also addressed paying attention to the oscillation associated with fundamental modes. ",
        "introduction": "One of the most important roles of numerical simulations in general relativity is to predict gravitational waveforms emitted by general relativistic and dynamical astrophysical phenomena. Rotating stellar core collapse and nonspherical oscillation of neutron stars are among the possible sources of gravitational waves. Therefore, fully general relativistic numerical simulation for them is an important subject in this field \\cite{HD}.  To date, there has been no systematic work for computation of gravitational waves from rotating stellar core collapse to a neutron star in fully general relativistic simulation (but see \\cite{Siebel}). The gravitational waveforms have been computed only in the Newtonian gravity \\cite{Newton,Newton1,Newton2,Newton3,Newton4,Newton5,Newton6} or in an approximate general relativistic gravity \\cite{HD} using the so-called conformal flatness approximation (or Isenberg-Wilson-Mathews approximation). As demonstrated in \\cite{HD}, general relativistic effects modify the evolution of the collapse and emitted gravitational waveforms significantly. Thus, the simulation in full general relativity appears to be the best approach for accurate computation of gravitational waves.  In the case that the progenitor of the core collapse is not very rapidly rotating, nonaxisymmetric instabilities do not set in and, hence, the collapse will proceed in an axisymmetric manner. In such a collapse, the amplitude of gravitational waves measured in a local wave zone at $r \\approx \\lambda$ where $\\lambda$ denotes the gravitational wave length will be smaller, by two or three orders of magnitude, than that in highly nonaxisymmetric phenomena such as mergers of binary neutron stars and black holes. The amplitude of gravitational waves from an oscillating neutron star is also likely to be small due to its small nonspherical deformation. Technically, it is not easy to extract gravitational waves of small amplitude from metric computed in numerical simulations, in which a numerical noise is in general contained. The numerical noise is generated due to the following reasons:  Gravitational waves are usually extracted from the metric in the wave zone in general relativistic simulations. Although they should be extracted at the null infinity, the outer boundaries of computational domain are located at a finite radius whenever the 3+1 formalisms are adopted. Thus, the outer boundary conditions are imposed at finite radii and in general they are approximate conditions. As a result, a small numerical error may be excited around the outer boundaries. Here, the possible candidates of the numerical error are unphysical nonwave modes, spurious gauge modes, back reflections at the outer boundaries, and roundoff errors.   In this paper, we study gravitational waves from oscillating neutron stars in axial symmetry. Neutron stars in equilibrium are simply modeled by $n=1$ polytropes. Oscillations of neutron stars are followed by axisymmetric numerical simulations in full general relativity. Gravitational waves are extracted using a gauge-invariant wave extraction technique. The gauge-invariant variables are not contaminated by gauge modes and, hence, we can focus on other error sources using this variables. We also adopt a quadrupole formula for approximately computing gravitational waveforms to clarify its validity.  This work was planned from the following four motivations. The first one is to specify the error sources contained in the gauge-invariant variables extracted in the local wave zone. As mentioned above, they could be contaminated by nonwave components and numerical errors. In particular, it is important to specify systematic error components contained in the gauge-invariant variables since as indicated in Sec. IV, the systematic errors may be eliminated at least partly if their origin is clarified.  The second motivation is to understand how large computational domains are needed to extract gravitational waveforms within $\\sim 10$\\% error. Since the gauge-invariant variables are extracted at finite radii, gravitational waveforms (in particular the amplitude) are slightly different from the asymptotic ones. It is important to clarify how magnitude of the error depends on the radius at which we impose the outer boundary conditions and on the radius at which we extract gravitational waves. A similar study was carried out about 15 years ago by Abrahams and Evans \\cite{AE}. However, they were interested only in specific gauge conditions which were often used in axisymmetric numerical simulations in general relativity at that time. Moreover, the simulations were carried out only for non-rotating stars. In this paper, we adopt a different gauge condition often used nowadays in three-dimensional simulations, and report numerical results both for nonrotating and rotating stars.  The third motivation, in which we are most interested in the present study, is to investigate validity of a quadrupole formula in fully general relativistic simulations. For computation of gravitational waves generated by oscillations of gravitational field such as quasinormal mode ringings of black holes, quadrupole formulas cannot work. However, in rotating stellar core collapse to a neutron star and in oscillating neutron stars in which gravitational waves are generated mainly by matter motions, quadrupole formulas may be able to yield an accurate waveform. This method can be applied much more easily than geometrical methods in which gravitational waves are extracted from metric in the wave zone. Thus, a quadrupole formula which can yield high-quality approximate waveforms will be a robust method for computing gravitational waves of small amplitude from a noisy numerical data set. Note that a similar work has been already done by Siebel et al. \\cite{Siebel0,Siebel} in a null-cone formulation. We here carry out the similar study for a 3+1 approach.  The last motivation is to understand the properties of oscillations of rotating neutron stars. During rotating stellar core collapse, gravitational waves associated with oscillations of a formed protoneutron star are likely to be emitted (see, e.g., \\cite{HD}). From the study for oscillating and rotating neutron stars, we will be able to understand what oscillation modes are relevant for the emission of gravitational waves during core collapses. We here pay attention to two fundamental oscillation modes (quasiradial and quadrupole $p$ modes of no node for the density perturbation) which are candidates for the dominant modes in the oscillating and rotating stars formed after the collapse.  This paper is organized as follows. In Sec. II, our numerical implementations for axisymmetric general relativistic simulation are briefly reviewed. In Sec. III, the gauge-invariant wave extraction technique and the quadrupole formula adopted in the present work are described. Sec. IV presents numerical results of gravitational waveforms emitted from oscillating neutron stars. The simulations were performed both for nonrotating and rotating neutron stars using an axisymmetric code recently developed \\cite{S2002}. Sec. IV is devoted to a summary. Throughout this paper, we adopt the geometrical units in which $G=c=1$ where $G$ and $c$ are the gravitational constant and the speed of light, respectively. ",
        "conclusions": "We have studied gravitational waves from axisymmetrically oscillating neutron stars adopting the gauge-invariant wave extraction method as well as the quadrupole formula. It is found that several types of the nonwave components such as the stationary parts of metric and numerical errors are contained in the gauge-invariant variables. The numerical errors are generated due to an approximate treatment for the outer boundary conditions. We illustrate a method to subtract the dominant components of the numerical errors and demonstrate it possible to extract gravitational waves even from such noisy data sets with a residual of magnitude $\\sim 10^{-3}$.  The gravitational waveforms computed in the quadrupole formula agree well with those obtained from the gauge-invariant technique besides a systematic underestimation of the amplitude by $\\sim 20$\\%. An important point is that the evolution of the wave phase and the modulation of the amplitude are computed with a good accuracy. This indicates that for a study of gravitational waveforms from rotating stellar core collapse to a protoneutron star, the quadrupole formula will be a useful tool in fully relativistic simulations. It should be also addressed that the result in this paper supports the treatment in \\cite{HD} in which gravitational waveforms are computed using a quadrupole formula in approximate general relativistic simulations.  The gauge-invariant variables are extracted for various values of extraction radii. It is found that to extract gravitational waves within $\\sim 10\\%$ error, the extraction radius has to be larger than $\\sim 90\\%$ of the gravitational wave length. If the outer boundaries are located in the near zone with $L < \\lambda$, the amplitude of gravitational waves is overestimated: For $L \\approx 2\\lambda/3$, it is overestimated by $\\sim 20\\%$. For $L < 2\\lambda/3$, the factor of the overestimation is even larger.   In the present work, the amplitude of gravitational waves in a local wave zone is much larger than that of systematic numerical errors. This fact enables to subtract them from the gauge-invariant variables accurately. If the magnitude of the errors is much larger than that of the amplitude of gravitational waves, however, it would not be possible to carry out an accurate subtraction. For example, in rotating stellar core collapse, the amplitude of gravitational waves in the local wave zone at $r \\sim \\lambda$ is at most $\\sim 10^{-5}$ according to gravitational waveforms calculated by a quadrupole formula \\cite{HD}. To extract gravitational waves of such small amplitude, it is necessary to reduce the magnitude of the numerical errors. To achieve that, we need to impose more accurate outer boundary conditions. Developing such conditions is crucial in computing gravitational waves of small amplitude of $O(10^{-5})$ from raw data sets of metric.  Another possible method for computing accurate gravitational waves of small amplitude is to adopt a quadrupole formula taking into account higher-order post Newtonian terms. As indicated in this paper, the simple quadrupole formula underestimates the amplitude of gravitational waves by $O(M/R)$. In rotating stellar core collapse, the error in the amplitude will be $\\sim 10\\%$. To compute the amplitude within $\\sim 1\\%$ error, we should take into account higher general relativistic corrections. In quadrupole formulas with the higher post Newtonian corrections (as derived in \\cite{BDS}), it may be possible to obtain gravitational waveforms within 1\\% error. Such formulas will be useful to extract gravitational waves of small amplitude from rotating stellar core collapses and from oscillating neutron stars.  In addition to the study for gravitational wave extraction, oscillation modes of rotating neutron stars are analyzed. It is found that two modes (the fundamental quadrupole and quasiradial modes) are dominantly excited due to the global oscillation. The frequency of the quadrupole mode is proportional to $\\sqrt{GM/R^3}$, and is higher than that of the quasiradial one for the typical values of mass and radius of neutron stars. It is shown that the amplitude of the quadrupole mode decreases with time due to an incoherent nature of the oscillation, but that of the quasiradial mode is not damped quickly, hence being the dominant mode after several dynamical timescales. We expect that in rotating stellar core collapse to a protoneutron star in a nearly quasistationary state, these two modes may be the main components in the burst phase of gravitational waves. The quadrupole mode will be damped within a few dynamical timescales and subsequently the quasiradial mode will be the dominant component to be longterm quasiperiodic waves.  \\begin{center} {\\bf Acknowledgments} \\end{center}  Numerical computations were carried out on the FACOM VPP5000 machine in the data processing center of National Astronomical Observatory of Japan. This work is in part supported by Japanese Monbu-Kagakusho Grant (Nos. 13740143, 14047207, 15037204, and 15740142).  \\appendix"
    },
    "0402/astro-ph0402467_arXiv.txt": {
        "abstract": "We present high angular resolution ($\\sim0\\rlap.{''}1$-$0\\rlap.{''}4$) VLA observations at 2 and 6 cm made in 1983, 1986, and 1995 toward the ultracompact bipolar H~II region NGC~7538~IRS1. We find, at both wavelengths, clear evidence of a decrease in the emission from the lobes. This decrease, of orden 20-30\\%, has not been observed previously in any ultracompact H~II region. Most likely, it is due to recombination of the ionized gas in the lobes as a result of a decrease in the available ionizing photon flux. It is unclear if this decrease in the ionizing photon flux is due to an intrinsic change in the exciting star or to increased absorption of ionizing photons in the optically-thick core of the nebula. ",
        "introduction": "The ultracompact H~II (UC HII) regions are small (diameters $<$ 10$^{17}$ cm) and dense (electron densities $>$ 10$^{4}$ cm$^{-3}$) structures of gas that are maintained ionized by deeply embedded, recently formed O stars. They have a few recurring morphological types, one of them is the bipolar type that comprises only a handful of objects (Churchwell 2002). The bipolar UC HII regions have an hourglass shape when projected in the sky, the waist is believed to be produced by confinement by a disk or torus of neutral gas, while the lobes contain outflowing gas.  One of the most studied bipolar UC HII regions is NGC 7538~IRS1. It was first found in the infrared  by Wynn-Williams, Becklin \\& Neugebauer (1974), and  later its free-free emission was clearly resolved  at centimeter wavelengths by Campbell (1984). Its radio spectrum becomes fully optically thin only above $\\sim$100 GHz, with a total flux density of $\\sim$1.6 Jy at this frequency (Akabane et al. 1992). At a distance of 2.8 kpc (Sandell, Wright, \\& Forster 2003), this flux density implies an ionizing photon rate of $1.5 \\times 10^{48}$ s$^{-1}$, that could be provided by an O8.5 ZAMS star with luminosity of $5 \\times 10^{4}~L_\\odot$ (Thompson 1984).  There is a large scale molecular outflow in the region (CO: Campbell \\& Thompson 1984; HCN$^+$: Batrla, Pratap \\& Snyder 1988), that was proposed to be associated with NGC~7538~IRS1, but observations made by Batrla, Pratap \\& Snyder (1988) suggest that the source that is driving the outflow is located $\\sim$15$''$ south of IRS1. Gaume et~al.  (1995) made H66$\\alpha$ recombination line  observations  and found that this source has one of the most broad profiles ($\\Delta v \\simeq$ 150 km s$^{-1}$) seen in UC HII regions. There are also several masers associated with the source (e.g. OH: Hutawarakorn \\& Cohen, 2003; CH$_3$OH: Minier, Booth \\& Conway, 2000 ; H$_2$O: Kameya et~al. 1990; $^{15}$NH$_3$: Gaume et~al. 1991), among them there is the rare formaldehyde (H$_2$CO) 4.83 GHz maser (Hoffman et~al. 2003). Observations in HCN, HCO$^+$ (Pratap, Batrla \\& Snyder 1991) and CS (Kawabe et~al. 1992) show the presence of denser material around the source; this can account for the confinement of the outflow seen in radio, especially in the south direction.  In this paper we present the analysis of archival VLA continuum observations of NGC 7538~IRS1 made with high angular resolution. The main goal of our study was to search for proper motions or time variability in this compact object. In $\\S$2 we describe the observations; in $\\S$3 we discuss them, and finally in $\\S$4 our main conclusions are given. ",
        "conclusions": "Our main conclusions can be summarized as follows.  1) We analyzed data taken at 2 cm toward the bipolar UC HII region NGC 7538 IRS1, finding that its lobes show a decrease in flux density in the order of 20-30\\% over a time interval of 11 years. The 6 cm data confirm this result.  2) This relatively large decrease in the emission from the lobes is tentatively interpreted as due to recombination of the ionized gas, caused by a decrease in the available ionizing photon flux.  3) At present it is unclear if this decrease in the ionizing photon flux of the lobes is due to an intrinsic change in the exciting star or to increased absorption of ionizing photons in the core of the object."
    },
    "0402/hep-th0402050_arXiv.txt": {
        "abstract": "{ We analyze diffeomorphism invariance in inflationary spacetimes regulated by a boundary at late time. We present the action for quadratic fluctuations in the presence of a boundary, and verify that it is gauge invariant precisely when the correct local counterterms are included. The scaling behavior of bulk correlation functions at the boundary is determined by Callan-Symanzik equations which predict scaling violations in agreement with the standard inflationary predictions for spectral indices of the CMB. } ",
        "introduction": "\\label{sec:introduction} It is useful to think of cosmological evolution in terms of a family of spatial slices, with time appearing as a parameter identifying the individual slices. The cosmological wave function \\cite{Hartle:1983ai} is then a functional which, in the Schr\\\"{o}dinger picture, takes the schematic form: \\begin{equation} \\Psi[\\phi] \\sim e^{iS[\\phi]} \\label{cosmowave} \\end{equation} where $\\phi$ is a set of variables defined on a three dimensional equal-time slice. In the semiclassical approximation $S[\\phi]$ can be identified as the Hamilton-Jacobi (H-J) functional, a versatile tool in cosmology \\cite{Salopek:1990jq,Salopek:1994sq,Salopek:1997he,Parry:mw}. The H-J functional is defined as the on-shell action, interpreted as a functional of the dynamical variables on the equal time slice; so the H-J form of the dynamical problem involves gravitational physics on a `bulk' manifold which ends at a `boundary', the equal time slice under consideration. It is therefore well suited to the study of gravitational physics on manifolds with a boundary, a problem which also has many other applications, such as in brane-world models.  In the present paper we study the gravity-scalar system on a manifold which, for definiteness, we take as an inflationary spacetime. The on-shell actions for these spacetimes contain late time divergences which can be regulated by truncating the manifold at a late time, resulting in a boundary. As a result the action in \\eqref{cosmowave} is a functional of the `boundary data', the variables $\\phi$ evaluated on the spatial slice corresponding to the late time cut-off. Our main results are: \\begin{itemize} \\item Diffeomorphism Invariance is not automatic in the presence of such a boundary. The simplest way to preserve diffeomorphism invariance is to introduce local counterterms on the boundary. We determine their form. \\item We compute the quadratic action for fluctuations around a manifold with a boundary. We present our result in terms of gauge invariant variables. \\item We interpret the spectral indices of the Cosmic Microwave Background (CMB) in terms of scaling violations of a `boundary theory'. This perspective is holographic in character, since a three-dimensional theory controls the four-dimensional physics. \\end{itemize}  The starting point for our discussion is the straightforward and explicit computation of the H-J functional~\\cite{Maldacena:2002vr,Larsen:2003pf}. The result of this computation suffers from a divergence as the time of the slice is taken to future infinity. This divergence is dominated by large wavelengths, and so it can be cancelled by adding a local boundary term --- a counterterm --- to the action.  However, the na\\\"{i}ve computation suffers from additional, and seemingly more serious, problems. As we will explain, the presence of an arbitrary boundary, introduced to regulate divergences, renders the H-J functional inconsistent with the full set of four dimensional diffemorphisms. This failure of local reparametrization invariance can be remedied by supplementing the standard action with a boundary term. We will show that this boundary term, designed to restore diffeomorphism invariance, is in fact the same as the counterterm needed to cancel infrared divergences.  The gravity side of the AdS/CFT correspondence \\cite{Maldacena:1997re,Gubser:1998bc,Aharony:1999ti} involves the on-shell action on a (possibly deformed) AdS-space. It is well known that this action exhibits infrared divergences due to the behavior of the metric near the boundary at spatial infinity. These divergences are naturally cancelled by the introduction of boundary counterterms \\cite{Witten:1998qj,Balasubramanian:1999re,Emparan:1999pm,Kraus:1999di}. In the context of the AdS/CFT correspondence, the counterterms are interpreted in the dual conformal field theory as the usual counterterms needed to cancel ultraviolet divergences in quantum field theory \\cite{Balasubramanian:1999re}; but their origin on the gravity side is less clear. Although we work with cosmological spacetimes for definiteness, our results are valid for asymptotically AdS-spaces as well. This suggests a new perspective on the counterterms in AdS/CFT, which is rooted solidly in gravity: counterterms are needed to maintain diffeomorphism invariance.   In standard cosmological perturbation theory it is customary to implement diffeomorphism invariance by introducing gauge invariant physical observables \\cite{Lifshitz:em,Bardeen:kt,Mukhanov:jd,Mukhanov:1990me,Riotto:2002yw}. We will extend this result to include the boundary, and present the quadratic action for fluctuations in this more general case, again written in gauge invariant form. This is one of our main results.  Diffeomophism invariance also constrains the dependence of the boundary theory on the gauge invariant variables. The origin of these additional constraints are the diffeomorphisms acting on the direction normal to the boundary. These transformations are implemented as scaling symmetries on the three-dimensional equal time surface, and so their effect is to determine the scale dependence of the correlation functions.  We refer to the equations determining the scale dependence as the Callan-Symanzik equations.  The Callan-Symanzik equations can be solved using techniques that are standard from renormalization group theory. The result of this analysis is general formulae for the correlation functions of the theory, determined by symmetries alone. To exploit these formulae, one must add dynamical input, {\\it e.g.} from slow roll inflation. Given this input, our expressions are renormalization group improved versions of the more conventional results. We consider both scalar and tensor fluctuations and determine, in particular, the scalar and tensor mode spectral indices, $n_s$ and $n_t$, which characterize the scale dependence of the CMB.  The terminology introduced to describe consequences of diffeomorphism invariance --- counterterm, Callan-Symanzik equations, and the Ward identity --- is  that of a local quantum field theory on the equal time slice. In the context of the AdS/CFT correspondence our terminology fully justified but, in cosmology, it refers to a conjectured dS/CFT correspondence \\cite{Strominger:2001pn, Strominger:2001gp, Klemm:2001ea, Spradlin:2001nb, Witten:2001kn, Balasubramanian:zh}. Such ideas, though rather speculative, have been used to address inflation \\cite{Maldacena:2002vr, Larsen:2003pf, Larsen:2002et, vanderSchaar:2003sz}. It would be extremely interesting if a truly holographic theory of cosmology could be established. However, whether cosmological holography is true or not, the counterterms we discuss are a universal part of the gravitational action, determined from diffeomorphism invariance alone. It may be that gravity is characterized by ``infrared universality classes\" which would be similar to the ``ultraviolet universality classes\" familiar from quantum field theory. In quantum field theory truly short distances decouple from long distance physics and similarly it seems that, in cosmology, truly large distance physics, beyond the horizon, decouples from short distance physics, {\\it i.e.} observable cosmology \\cite{Schalm:2004qk}. The notion of ``infrared universality classes\" could develop into a framework for addressing the notorious fine-tuning problems in cosmology. This might apply not only to the fine-tuning problems normally associated with inflation, but also to other naturalness problems associated with the decoupling of long and short distance physics, such as the cosmological constant problem.   This paper is organized as follows. In section 2 we review the cancelling of infrared divergences {\\it via} the introduction of counterterms. We then present an argument that identifies the origin of the counterterms as diffeomorphism invariance. In section 3 we compute the quadratic action of fluctuations around the background. To do this, we review the standard notion of gauge invariance in the bulk theory, and show how this can be extended to the boundary, precisely when the correct boundary terms are introduced. In Section 4 we discuss the consequences of diffeomorphism invariance for the form of correlation function. The constraints are summarized by a master equations which, in a special case, reduces to the Callan-Symanzik equation. In section 5 we solve the Callan-Symanzik equation. In particular we determine the spectral parameters of cosmological inflation as the scaling violations of the theory. Finally, in section 6, we conclude with an outlook for further developments. ",
        "conclusions": ""
    },
    "0402/nucl-th0402001_arXiv.txt": {
        "abstract": "Highly precise data on the magnetic dipole strength distributions from the Darmstadt electron linear accelerator for the nuclei $^{50}$Ti, $^{52}$Cr and $^{54}$Fe are dominated by isovector Gamow-Teller-like contributions and can therefore be translated into inelastic total and differential neutral-current neutrino-nucleus cross sections at supernova neutrino energies. The results agree well with large-scale shell-model calculations, validating this model. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0402/astro-ph0402159_arXiv.txt": {
        "abstract": "We examine the variation of the fine structure constant in the context of a two-field quintessence model. We find that, for solutions that lead to a transient late period of accelerated expansion, it is possible to fit the data arising from quasar spectra and comply with the bounds on the variation of $\\alpha$ from the Oklo reactor, meteorite analysis, atomic clock measurements, Cosmic Microwave Background Radiation and Big Bang Nucleosynthesis. That is more difficult if we consider solutions corresponding to a late period of permanent accelerated expansion.    \\vskip 0.5cm ",
        "introduction": "The recent claim that the spectra of quasars (QSOs) indicates the variation of the fine structure constant, $\\alpha$, on cosmologically recent times \\cite{Murphy:2003hw,Murphy:2002ve,Murphy:2001nu,Murphy:2000ns} has raised considerable interest in examining putative sources of this variation. In most models a possible variation of the fine structure constant is studied by arbitrarily coupling fields to electromagnetism, as suggested by Bekenstein \\cite{Bekenstein}.  Thus the proposals put forward sofar consider a scalar field \\cite{Olive:2001vz,Gardner}  (with an additional coupling to dark matter and to a cosmological constant in the former case)  and quintessence \\cite{Anchordoqui,Mota}. In fact, as discussed in Ref. \\cite{Kostelecky}, the couplings of gravito-scalar fields, such as the axion or the dilaton, to electromagnetism naturally arise in $N=4$ Supergravity in four dimensions, making this model particularly interesting.  It is worth mentioning that, in the latter model, the mass of these scalars can drive the accelerated expansion of the Universe, transient or eternal \\cite{Bertolami}.  In this work, we shall consider the implications for the variation of $\\alpha$ of a two-field quintessence model \\cite{Bento:2001yv}, with the quintessence fields coupled to electromagnetism, as proposed by Bekenstein \\cite{Bekenstein}. There are several motivations for studying potentials with coupled scalar fields. Firstly, if one envisages to describe the Universe dynamics from fundamental theories, it is most likely that an ensemble of scalar fields (moduli, axions, chiral superfields, etc) will emerge, for instance, from the compactification process in string or braneworld scenarios. Furthermore, coupled scalar fields are invoked for various desirable features they exhibit, as in the so-called hybrid inflationary and reheating models \\cite{LindeBKB}.  The model of. Ref.~\\cite{Bento:2001yv} has the additional bonus of leading to transient as well as permanent solutions for the late time acceleration of the Universe. The former solutions are desirable given that it has been recently pointed out that an eternally accelerating universe poses a challenge for string theory, at least in its present formulation, since asymptotic states are inconsistent with spacetimes that exhibit event horizons \\cite{Hellerman}.  Moreover, it is argued that theories with a stable supersymmetric vacuum cannot relax into a zero-energy ground state if the accelerating dynamics is guided by a single scalar field \\cite{Hellerman}, a problem that can be circumvented in the two-field model we are considering \\cite{Bento:2001yv}.  We now turn to the available observational bounds on the variation of $\\alpha$.  Observations of the spectra of 128 QSOs with $z=0.2 - 3.7$ suggest that, for $z>1$, $\\alpha$ was smaller than at present \\cite{Murphy:2003hw,Murphy:2002ve,Murphy:2001nu,Murphy:2000ns}  \\beq {\\Delta \\alpha \\over \\alpha} \\equiv {\\alpha (z)-\\alpha_0 \\over \\alpha_0} = (-0.54 \\pm 0.12) \\times 10^{-5}~, \\label{Murphy} \\eeq at $4.7 \\sigma$.  The most recent data is from Chand {\\it et al.} \\cite{Chand1,Chand2} obtained via a new sample of Mg~II systems from distant quasars with redshifts in the range $0.4 \\le z \\le 2.3$ yield more stringent bounds ($3~\\sigma$), namely:  \\beq {\\Delta \\alpha \\over \\alpha} = (-0.06 \\pm 0.06) \\times 10^{-5}~, \\label{Chand1} \\eeq where terrestrial isotopic abundances have been assumed. If, instead, low-metalicity isotopic abundances are assumed, Chand {\\it et al.} obtain   \\beq {\\Delta \\alpha \\over \\alpha} = (-0.36 \\pm 0.06) \\times 10^{-5}~, \\label{Chand2} \\eeq in which case the statistical inconsistency with Murphy {\\it et al.}, Eq.~(\\ref{Murphy}), is clearly smaller. Notice that, in contrast with Webb {\\it et al.}, who use different lines from different multiplets and elements, Chand {\\it et al.} use mostly Mg II data yielding a smaller but better quality dataset.   On the other hand, the Oklo natural reactor provides a bound, at $95 \\%$ CL,  \\beq -0.9 \\times 10^{-7} < {\\Delta \\alpha \\over \\alpha} < 1.2 \\times 10^{-7}~, \\eeq for $z=0.14$ \\cite{Damour:1996zw,Fujii:2003gu,Fujii:1998kn}. Notice, however, that the use of an equilibrium neutron spectrum  has been criticized in Ref.~\\cite{Lamoreaux:2003ii}, where a lower bound on the variation of $\\alpha$ over the last two billion years is given by   \\beq {\\Delta \\alpha\\over \\alpha} \\geq +4.5 \\times 10^{-8}~. \\eeq  Estimates of the age of iron meteorites ($z=0.45$), combined with a measurement of the Os/Re ratio resulting from the radioactive decay $^{187}$Re~$\\to~^{187}$Os, gives \\cite{Olive:2003sq,Fujii:2003uw, Olive:2002tz}  \\beqa {\\Delta \\alpha \\over \\alpha} = (-8 \\pm 8) \\times 10^{-7}~, \\eeqa at $1\\sigma$, and \\beqa -24 \\times 10^{-7} < {\\Delta \\alpha \\over \\alpha} < 8 \\times 10^{-7}~, \\eeqa at $2\\sigma$.   Notice that, if the variation of the fine structure constant is linear, the Murphy {\\it et al.} observations of QSO absorption spectra and of geochemical tests from meteorites are incompatible given that the former yields $\\Delta \\alpha / \\alpha \\simeq 5 \\times 10^{-6}$ at $z =0.5$, while the latter leads to $\\Delta \\alpha / \\alpha \\le 3 \\times 10^{-7}$ at $z=0.45$ \\cite{Mota}. However, this problem may not exist if one uses the Chand {\\it et al.} dataset.  Moreover, observations of the hyperfine frequencies of the $^{133}$Cs and $^{87}$Rb atoms in their electronic ground state, using several laser cooled atomic fountain clocks, give, at present ($z=0$) \\cite{Marion:2002iw} (see also \\cite{Bize:bj})  \\beq {\\left\\vert {\\dot{\\alpha} \\over \\alpha}\\right\\vert} < 4.2 \\times 10^{-15}~\\mbox{yr}^{-1}~, \\eeq where the dot represents differentiation with respect to cosmic time. Tigher bounds arise from the remeasurement of the $1s-2s$ transition of the atomic hydrogen and comparison with a previous measurement with respect to the ground state hyperfine splitting in $^{133}$Cs and combination with the drift of an optical transition frequency in $^{199}$Hg$^+$, that is \\cite{Fischer}:  \\beq {\\dot{\\alpha} \\over \\alpha} = (-0.9 \\pm 4.2) \\times 10^{-15}~\\mbox{yr}^{-1}~. \\eeq   There are also constraints coming from  Cosmic Microwave Background Radiation (CMBR), where  \\beq \\vert {\\Delta \\alpha / \\alpha} \\vert \\leq 10^{-2}~, \\eeq at $z=10^3$ \\cite{CMBBBNconstr,CMBconstr}, and from Big Bang nucleosynthesis (BBN),  \\beq - 6 \\times 10^{-4} < {\\Delta \\alpha / \\alpha} < 1.5 \\times 10^{-4}~, \\eeq at $z=10^8~-~10^{10}$ \\cite{CMBBBNconstr,BBNconstr}.    In addition, we should take into account the equivalence principle experiments, which imply \\cite{Olive:2001vz}  \\beq {\\zeta _F \\over \\sqrt \\omega} \\leq 10^{-3}~, \\label{zetaf} \\eeq where $\\zeta _F$ is the coupling between the scalar and electromagnetic fields and $\\omega \\equiv {M_*^2 \\over 2 M^2}$,  $M$ being  the reduced Planck mass $(M\\equiv M_P/\\sqrt{8\\pi})$ and $M_*$ the corresponding analogue in the scalar sector.  Notice that not all models put forward sofar satisfy the abovementioned bounds.  For instance, quintessence models like the last one in \\cite{Anchordoqui} and the $N=4$ Supergravity models in four dimensions are not consistent with Webb {\\it et al.} data.  In the following, we shall present our model and show that it is consistent with all the available data for a suitable choice of the model parameters. ",
        "conclusions": "In this work, we have studied the variation of the fine structure constant in the context of a quintessence model with two coupled scalar fields. We find that transient acceleration models can fit the latest QSO data and comply with the upper bounds on $\\Delta \\alpha$ from the Oklo reactor, meteorite analysis and the atomic clock measurements.  For permanent acceleration models, however, it is more difficult to fit the QSO data and satisfy the Oklo, meteorite and BBN bounds simultaneously.  We have studied the sensitivity of our results to $\\zeta_1$ and $\\zeta_2$, the couplings of quintessence fields with electromagnetism, and we have found that, in order to be consistent with the data, these parameters must be at least one order of magnitude smaller than the upper bound implied by the Equivalence Principle.  On a more general ground, we could say that establishing whether there is a variation of the fine structure constant and, in the affirmative case, identifying its origin, remains a difficult task before a deeper analysis of the systematic errors of the observations and studies of the degeneracies with the various cosmological parameters.  This is particularly evident in what concerns the compatibility of the datasets of Murphy {\\it et al.} and Chand {\\it et al.}   \\begin"
    },
    "0402/astro-ph0402473_arXiv.txt": {
        "abstract": "The nature of SN\\,1961V has been uncertain.  Its peculiar optical light curve and slow expansion velocity are similar to those of super-outbursts of luminous blue variables (LBVs), but its nonthermal radio spectral index and declining radio luminosity are consistent with decades-old supernovae (SNe). We have obtained {\\it Hubble Space Telescope} STIS images and spectra of the stars in the vicinity of SN\\,1961V, and find Object 7 identified by Filippenko et al.\\ to be the closest to the optical and radio positions of SN\\,1961V. Object 7 is the only point source detected in our STIS spectra and only its H$\\alpha$ emission is detected; it cannot be the SN or its remnant because of the absence of forbidden lines. While the H$\\alpha$ line profile of Object 7 is remarkably similar to that of $\\eta$ Car, the blue color (similar to an A2\\,Ib supergiant) and lack of appreciable variability are unlike known post-outburst LBVs. We have also obtained Very Long Baseline Array (VLBA) observations of SN\\,1961V at 18~cm. The non-detection of SN\\,1961V places a lower limit on the size of the radio-emitting region, 7.6 mas or 0.34 pc, which implies an average expansion velocity in excess of 4,400 km~s$^{-1}$, much higher than the optical expansion velocity measured in 1961. We conclude the following: (1) A SN occurred in the vicinity of SN\\,1961V a few decades ago. (2) If the SN\\,1961V light maximum originates from a giant eruption of a massive star, Object 7 is the most probable candidate for the survivor, but its blue color and lack of significant variability are different from a post-outburst $\\eta$ Car. (3) The radio SN and Object 7 could be physically associated with each other through a binary system. (4) Object 7 needs to be monitored to determine its nature and relationship to SN\\,1961V. ",
        "introduction": "SN\\,1961V in the outskirts of the Sc galaxy NGC\\,1058 \\citep[distance = 9.3 Mpc;][]{T80,Setal96} is the prototype of Zwicky's Type~V supernovae (SNe), which are now classified as Type II Peculiar SNe. SN\\,1961V was unusual in many respects \\citep{BG71}. First, SN\\,1961V is one of the very few SNe whose progenitors are known. The progenitor of SN\\,1961V was visible as an 18 mag star from 1937 to 1960 \\citep{Z64}. Second, the optical light curve of SN\\,1961V was more complex and much broader than any other SN ever observed. It peaked at $\\sim$12 mag in late 1961, remained visible for several years, and faded to 21.7 mag in 1970 \\citep{BG71,BA70}. Third, the initial expansion velocity of SN\\,1961V, 2,000--3,000 km s$^{-1}$, was much lower than the typical expansion velocity of 15,000--20,000 km s$^{-1}$ for most SNe \\citep{Z64,BG71}.  These peculiar properties of SN\\,1961V have raised skepticism concerning its nature as a SN. Based on the extended optical light curve and the anomalously low expansion velocity, \\citet{Getal89} suggested that SN\\,1961V was a luminous blue variable (LBV) similar to $\\eta$~Car, and its ``SN explosion\" was actually a super-outburst. \\citet{Fetal95} obtained {\\it Hubble Space Telescope (HST)} WFC1 images of SN\\,1961V and identified a red star with $R$ = 24.55, their Object 6, as a candidate for the LBV survivor. Using archival {\\it HST} WFPC2 images, \\citet{VDFL02} suggested a fainter and redder star as an alternative candidate for the LBV survivor and called it Object 11 as an extension of the \\citet{Fetal95} object list. The LBV hypothesis cannot explain, however, the nonthermal radio spectral index of SN\\,1961V, $-0.4\\pm0.3$ \\citep{CHB88} or $-0.79\\pm0.23$ \\citep{Setal01b}. The radio emission from $\\eta$ Car and its ejecta is thermal and optically thick, as it shows a positive spectral index and a complete lack of polarization \\citep{Detal95,DWL97}. Furthermore, SN~1961V is more than 1,000 times more luminous than $\\eta$ Car at 20-cm wavelength \\citep{R83,Setal01b} and its radio light curve is similar to those of radio SNe \\citep{Setal01a}.  Existent images of SN\\,1961V show a complex environment. Narrow-band, emission-line images reveal two \\ion{H}{2} regions within 3$''$ to the north of SN\\,1961V \\citep{F85,CHB88}, while broad-band continuum {\\it HST} WFC1 images show massive stars in and around the \\ion{H}{2} regions and in the vicinity of SN\\,1961V \\citep{Fetal95}. SN\\,1961V is clearly associated with a star-forming environment, where SNe are expected to occur. Indeed, radio observations \\citep{CHB88,Setal01b} show a nonthermal source at the south edge of the eastern \\ion{H}{2} region, coincident to within 1$''$ of the optical position of SN\\,1961V, and a fainter nonthermal radio source in the western \\ion{H}{2} region, corresponding to an unrelated supernova remnant (SNR).  To determine the nature of SN\\,1961V, it is crucial to recover the optical counterpart of SN\\,1961V with a high degree of certainty, and to determine spectroscopically whether it is a surviving LBV or a SN turning into a SNR. Thus, we have obtained {\\it HST} imaging and spectroscopic observations of SN\\,1961V. We have also obtained high-resolution radio observations to determine the size of the radio source at SN\\,1961V. This paper reports these observations of SN\\,1961V and our analysis of its nature. ",
        "conclusions": "\\subsection{Stellar and Interstellar Environment of SN\\,1961V}  The broad-band STIS image in Figure 1 represents the deepest high-resolution image of SN\\,1961V and its surroundings ever obtained. In addition to the eleven bright stellar objects identified by \\citet{Fetal95} and \\citet{VDFL02}, many fainter stars are also detected. These eleven bright objects, marked in Figure 1c, are mostly supergiants with $V = $ 24 to 25.5. To illustrate the position of SN\\,1961V, radio contours extracted from the \\citet{Setal01b} 18 cm  VLA map are plotted over the STIS image in Figure 1d. The eastern radio source corresponds to SN\\,1961V, and our new astrometric calibration (\\S 2.1) now places Object 7 the closest to SN\\,1961V. Three nebular arcs are visible and form an apparent supershell structure about $5\\farcs3 \\times 3\\farcs0$ (or 240 pc $\\times$ 135 pc) in size. Encompassed within the supershell are SN\\,1961V and Objects 6, 7, 9, and 11, while projected along the supershell rim are Objects 5, 8, and 10.  To date, the best narrow-band images of SN\\,1961V are still those obtained by \\citet{F85} using the Kitt Peak 2.1 m telescope. His red continuum image had a limiting magnitude of $\\sim$ 21 mag, so it did not detect any of the stars in the vicinity of SN\\,1961V shown in our STIS image. Two \\ion{H}{2} regions are detected in the [\\ion{O}{3}] image, but they appear better resolved in the H$\\alpha$+[\\ion{N}{2}] image: the eastern \\ion{H}{2} region consists of a bright diffuse emission region and two faint emission knots to the south, while the western \\ion{H}{2} region is resolved into two emission patches with similar brightnesses.  Comparisons between our STIS image and Fesen's narrow-band images show the following correspondences. The bright diffuse emission from the eastern \\ion{H}{2} region is centered on the bright Object 8 projected against the northeast rim of the supershell, and the two H$\\alpha$-emission knots in the southern extension are coincident with Object 7 inside the supershell and Object 5 along the south rim of the supershell, respectively. In the western \\ion{H}{2} region, the eastern emission patch is coincident with the bright Object 10, but the western emission patch has no detectable stellar counterparts in our STIS image. The image of Object 10 is elongated and somewhat resolved, indicating that it consists of multiple stars; its location inside an \\ion{H}{2} region further indicates that it contains a group of massive stars. The stars in Object 10 are probably at least a few Myr old, as they are coincident with a radio SNR identified by \\citet{Setal01b}, the western radio source in Figure~1d.  In summary, the STIS image illustrates that SN\\,1961V is associated with a complex, extended star forming region. It is projected within the central cavity of a supershell, with bright \\ion{H}{2} regions distributed along the shell rim. The presence of both \\ion{H}{2} regions and SNRs indicates that a mixture of stellar populations at different ages exist within $\\sim$100 pc of SN\\,1961V.  \\subsection{The Optical Counterpart of SN 1961V}  The optical position of SN\\,1961V was determined by \\citet{K86} to an accuracy better than 0\\farcs1, and this position is coincident with a fading nonthermal radio source \\citep{BC85,CHB88,Setal01b}. As shown in Figure~1d, Object 7 is the closest to SN\\,1961V, the eastern radio source.  Our long-slit STIS spectroscopic observations have been designed to detect line emission, as the stellar sources are faint and the detection of their continuum requires unrealistically long exposures. Of all stellar sources in the slit, only one is detected in H$\\alpha$ emission (see Fig.~2). Aligning the STIS direct image with the H$\\alpha$ spectrogram, we find that the point source corresponds to Object 7. The wavelength scale in Figure~2 is calibrated for the slit center; as Object 7 is offset by 0\\farcs46 from the slit center, a wavelength correction of $-$5.1 \\AA\\ should be applied.\\footnote{The STIS spectrogram has a spectral scale of 0.554 \\AA~pixel$^{-1}$ and a spatial scale of 0\\farcs05 pixel$^{-1}$; therefore, a 0\\farcs46 offset from the slit center corresponds to a wavelength offset of $0.554 \\times 0.46 / 0.05$ = 5.1 \\AA.} The corrected wavelength of the H$\\alpha$ emission from Object 7 corresponds to a heliocentric velocity of $V_{\\rm hel} = 520 \\pm 10$ km~s$^{-1}$, which is consistent with the systemic velocity of the \\ion{H}{2} regions near SN\\,1961V measured by \\citet{Getal02}. The H$\\alpha$ line profile of Object 7, shown in Figure 3, appears to have a narrow core and broad wings, which can be measured up to $\\pm$550 km~s$^{-1}$ (limited by the noise level of the spectrum). The absence of forbidden lines such as [\\ion{O}{1}] $\\lambda$6300 and [\\ion{O}{3}] $\\lambda$5007 indicates that the H$\\alpha$ emission is stellar as opposed to originating from a SN/SNR transition object like SN\\,1979C or SN\\,1980K \\citep{Fetal99}.  Our STIS spectrograms also detected a patch of diffuse emission in the H$\\alpha$ and [\\ion{O}{3}] $\\lambda\\lambda$4959, 5007 lines. This emission originates from the northwest rim of the supershell and corresponds to the northern part of the western \\ion{H}{2} region identified by \\citet{F85}. The velocity of this diffuse emission cannot be accurately determined, but is consistent with $V_{\\rm hel} \\sim 520$ km~s$^{-1}$.  \\subsection{Radio Properties of SN\\,1961V}  Recent VLA observations of SN\\,1961V \\citep{Setal01b} show that it is a nonthermal radio source with a spectral index of $\\alpha =-0.79\\pm0.23$ ($S_{\\nu}\\propto \\nu^{\\alpha}$) and a time decay index of $\\beta =-0.69\\pm0.23$ at $\\sim$20~cm ($S_{\\nu}\\propto t^{\\beta}$). The spectral and temporal properties of SN\\,1961V are consistent with other decades-old radio SNe, for example, SNe 1970G, 1957D, and 1950B \\citep{Setal01a,S01}. These properties are not consistent with radio observations of LBVs, which have an overall flat radio spectrum when convolved to the VLA resolution limit of SN~1961V, for example, AG~Car, He~3-519, HR~Car, WRA~751, P~Cyg, and $\\eta$~Car \\citep{DW02,DW03,R83,DWL97}.  A more apt radio comparison may be made with SN\\,1954J in NGC\\,2403 \\citep[3.2 Mpc;][]{Fetal01}, which has been established to be an outburst of the LBV ``Variable 12'', misidentified as a SN \\citep{HD94,SHG01}. The region near SN\\,1954J was observed with the VLA at 20~cm about 31 yr following its outburst, but no source was detected with a $3\\sigma$ limit of 0.35 mJy beam$^{-1}$ \\citep{ECB02}. On the other hand, the VLA detection of SN\\,1961V at 18~cm about 38 yr after the explosion indicates that it would been detected with a flux density of 0.98 mJy beam$^{-1}$ were it in NGC\\,2403. Clearly, SN\\,1961V is more than three times more luminous than SN\\,1954J at 18-20~cm.  Our recent VLBI experiment (\\S2.2) allows us to make a more definitive radio analysis than previously possible with resolution-limited VLA studies. The VLBI non-detection of SN\\,1961V indicates that its radio emitting region is larger than the VLBI resolution element; thus the minimum diameter of SN\\,1961V is 7.56 mas, or 0.34 pc for a distance of 9.3 Mpc. Assuming that the nonthermal radio emission detected by the VLA observations originates from synchrotron radiation associated with SN shocks \\citep{C82,CF94}, this minimum size implies an average expansion velocity in excess of 4,400 km s$^{-1}$, between the time of the VLBI experiment and the reported SN explosion, i.e., $\\sim$38 yr. This expansion velocity is much higher than the optical expansion velocity of 2,000 km s$^{-1}$ measured in 1961 November, toward the end of the broad maximum in the light curve.  Such a large discrepancy in expansion velocities determined at optical and radio wavelengths has been seen previously in SN\\,1986J, another Type II Peculiar SN. The optical expansion velocity of SN\\,1986J measured in 1986, several years after its possible explosion in 1982-1983, was less than 1,000 km s$^{-1}$ \\citep{Retal87}. However, VLBI observations of SN\\,1986J in 1990 and 1999 suggest that its expansion velocity was 20,000 km~s$^{-1}$ at $t = $0.25 yr and 6,000 km~s$^{-1}$ at $t = $15.9 yr \\citep{BBR02}, well in excess of the optical expansion velocity. In the case of SN\\,1986J, its identification as a SN has been commonly accepted.  \\subsection{The Nature of SN\\,1961V}  Was the SN\\,1961V event a SN explosion or a super-outburst of an LBV? The strongest support for the SN hypothesis has been provided by radio observations. The nonthermal radio spectral index, high radio luminosity, and the temporal decline of radio luminosity are all consistent with the existence of a decades-old SN. Furthermore, these radio properties are not like those of any known LBV: radio emission from LBVs are predominantly thermal in origin, and LBVs (including SN\\,1954J) are much fainter than SN\\,1961V.  The size of SN\\,1961V's emitting region may further constrain its nature. Object 7 is unresolved in the {\\it HST} STIS images, placing an upper limit of 0\\farcs05 = 2.25 pc on its size. The VLBI non-detection places a lower limit of 7.6 mas = 0.34 pc on the diameter of the radio-emitting region. These sizes are significantly larger than the size of $\\eta$ Car's nebula in 1995, 0.09 pc $\\times$ 0.17 pc, $\\sim$150 yr after it was ejected during ``The Great Eruption'' in 1837-1860 \\citep{HD94,SG98}. Whether SN\\,1961V was a SN explosion or an LBV outburst, the ejected material expands and the expansion velocity can be estimated from the size and time lapse. The above lower and upper limits on size require that the expansion velocity is at least 4,000 km~s$^{-1}$ but at most 27,000 km~s$^{-1}$. Such high expansion velocity is consistent with a SN explosion, but not an LBV outburst.  The hypothesis that SN\\,1961V was an LBV outburst can be verified only if the LBV survivor can be convincingly identified. Our STIS observations suggest that Object 7 is the most likely candidate for the LBV survivor because it is the closest to the optical and radio position of SN\\,1961V and is the only stellar object with the H$\\alpha$ emission line detected. However, Object 7 is not as red as $\\eta$ Car after a super-outburst. Using {\\it HST} WFC1 images taken in 1991 December, \\citet{Fetal95} reported $V = 24.22$, $R = 23.37$, and $I = 23.79$ (errors $<$ 0.2 mag) for Object 7. Using {\\it HST} WFPC2 $F606W$ images taken in 1994 September and $F450W$ and $F814W$ images taken in 2002 July, \\citet{VDFL02} reported $B = 24.04$, $V = 23.85$, and $I = 23.83$ for Object 7 with errors $\\sim$ 0.14 mag. These can be compared to the LBV Variable 12 that was responsible for SN\\,1954J: $B = 22.7$, $V = 21.9$, $R = 21.1$, and $I = 20.9$ (errors $\\sim$ 0.2 mag) in 1999 February \\citep{SHG01}. It is clear that the color of Object 7, $B-I = 0.21$ at about 41 yr after the light maximum, is not as red as the color of Variable 12, $B-I = 1.8$ at about 45 yr after the outburst. Because of the color difference, Object 7 at 41 yr after is about 1 mag brighter in $B$ but 0.6 mag fainter in $I$ compared to Variable 12 at 45 yr after the outburst.  We can also compare the H$\\alpha$ line profile of Object 7 to that of $\\eta$ Car. At the distance of SN\\,1961V, 9.3 Mpc, $\\eta$ Car would not be resolved from its ejecta; therefore, we have extracted an integrated spectrum of $\\eta$ Car and its surrounding ejecta nebula.\\footnote{The echelle spectrograph on the 4m telescope at the Cerro Tololo Inter-American Observatory was used to map the kinematics of the ejecta nebula of $\\eta$ Car in the H$\\alpha$ and [\\ion{N}{2}] lines in 1996 January by Chu et al.  The spectral resolution, determined from the FWHM of the sky lines was $\\sim$ 12 km~s$^{-1}$. Parts of the data were published by \\citet{WDC99}.} The ejecta nebula of $\\eta$ Car shows pronounced [\\ion{N}{2}] lines relative to the H$\\alpha$ line, but the integrated spectrum of the star and the ejecta nebula is dominated by the stellar H$\\alpha$ emission and the nebular [\\ion{N}{2}] lines become negligible. Figure 3 shows that the H$\\alpha$ line profile of Object 7 is remarkably similar to that of $\\eta$ Car. The lack of a [\\ion{N}{2}] counterpart to the H$\\alpha$ emission from Object 7 indicates that the emitting material is dense and must be stellar, as is the case for LBVs or mass-losing stars in general. Unfortunately, the previous observations of the H$\\alpha$ line of SN\\,1961V  made by \\citet{Getal89} in 1986 were of much lower spectral and spatial resolution, so that the H$\\alpha$ and [\\ion{N}{2}] lines were blended and heavily contaminated by bright \\ion{H}{2} region emission. It is impossible to determine whether the H$\\alpha$ line width evolved from 1986 to present.  Is Object 7 an LBV? While its H$\\alpha$ line profile resembles that of $\\eta$ Car, its color is not red. In fact, based on its colors and magnitudes \\citet{Fetal95} assigned a spectral type of A2\\,Ib to Object 7. It is interesting to note that there exist A-type supergiants with similar H$\\alpha$ emission line profiles. For example, the A5\\,Ia-O star B324 in M33 shows an H$\\alpha$ emission line with a narrow core and broad wings, but M33 B324 is not known to exhibit variability and cannot be classified as an LBV \\citep{HMF90,Hetal94}. Similarly, the {\\it HST} photometric measurements of Object 7 by \\citet{Fetal95} and by \\citet{VDFL02} do not show significant variations from 1991 to 2002. The lack of large-amplitude variability from 30 yr to 40 yr after the light maximum is in sharp contrast with the variability of $\\eta$ Car after its Great Eruption. Therefore, even if Object 7 is an LBV in a dormant state, its blue color and lack of variability are unlike known LBVs that have gone through super-outbursts, such as $\\eta$ Car.  What is the nature of the SN\\,1961V event? The radio observations indicate the existence of a decades-old SN. It is possible that the SN exploded in 1961 and was responsible \\ for the light maximum of SN\\,1961V. It is also possible that the SN exploded prior to 1961 and did not contribute directly to the light curve of SN\\,1961V. In the latter case, the light maximum of SN\\,1961V might be caused by a catastrophic but not fatally-explosive event of a massive star, then the spatial coincidence would make Object 7 the most probable candidate for the survivor, but the event cannot have been a super-outburst from an $\\eta$ Car-like LBV. As the radio SN and Object 7 are both coincident with SN\\,1961V (well within 20 pc), there is a non-negligible possibility that all three objects are related to one another through a binary system similar to the situation in SN\\,1993J, where a massive binary companion of the SN progenitor is recently identified \\citep{Metal04}. We speculate that Object 7 could be a massive binary companion of the radio SN's progenitor. Either the SN itself or the SN ejecta impact on Object 7 may be responsible for the light curve of SN\\,1961V. The impact of SN ejecta on Object 7 injects energy into its atmosphere and causes it to expand and form the broad H$\\alpha$ emission line. Similar interactions have been suggested for the binary companions of Type Ia SNe \\citep[e.g.,][]{MBF00}.  We have learned from SN\\,1961V that the late evolution of massive stars is complex and confusing. Numerous parameters can affect the appearance of a SN or an outburst. Systematic studies of a large number of peculiar SNe are needed. In the meantime, it is necessary to monitor Object 7 photometrically and spectroscopically in the future for variability in order to achieve a better understanding of its nature and relationship with SN\\,1961V."
    },
    "0402/astro-ph0402645_arXiv.txt": {
        "abstract": "We study the acceleration of electrons and protons interacting with localized, multiple, small-scale dissipation regions inside an evolving, turbulent active region. The dissipation regions are Unstable Current Sheets (UCS), and in their ensemble they form a complex, fractal, evolving network of acceleration centers. Acceleration and energy dissipation are thus assumed to be fragmented. A large-scale magnetic topology provides the connectivity between the UCS and determines in this way the degree of possible multiple acceleration. The particles travel along the magnetic field freely without loosing or gaining energy, till they reach a UCS. In a UCS, a variety of acceleration mechanisms are active, with the end-result that the particles depart with a new momentum. The stochastic acceleration process is represented in the form of Continuous Time Random Walk (CTRW), which allows to estimate the evolution of the energy distribution of the particles. It is found that under certain conditions electrons are heated and accelerated to energies above $1\\,$MeV in much less than a second. Hard X-ray (HXR) and microwave spectra are calculated from the electrons' energy distributions, and they are found to be compatible with the observations. Ions (protons) are also heated and accelerated, reaching energies up to 10 MeV almost simultaneously with the electrons. The diffusion of the particles inside the active region is extremely fast (anomalous super-diffusion). Although our approach does not provide insight into the details of the specific acceleration mechanisms involved, its benefits are that it relates acceleration to the energy release, it well describes the stochastic nature of the acceleration process, and it can incorporate the flaring large-scale magnetic topology, potentially even its temporal evolution. ",
        "introduction": "}  Solar flares remain, after almost one hundred years of intense study, an unsolved problem for astrophysics. We define as {\\it flare} the sporadic transformation of magnetic energy to (1) plasma heating, (2) particle acceleration, and (3) plasma flows. It seems that the total magnetic energy released in a single flare is not equally spread to all three components. The energetic particles carry a very large fraction of the total energy released during a flare, reaching sometimes up to $50\\%$ \\citep{Saint02}.  Modeling the explosive energy release requires methods that can treat simultaneously the large scale magnetic field structures and the small-scale dissipation events. The convection zone actively participates in the formation and evolution of large scale structures by rearranging the position of the field lines and at the same time it adds new magnetic flux (emerging flux) and new stresses to the existing topologies. The loss of stability of several loop-like structures forms large scale disturbances (CME), which further disturb the pre-existing and so-far stable large scale structures. In this way, the 3-D magnetic topologies are constantly forced away from the potential state (if they ever reach one) due to slow or abrupt changes in the convection zone. All these non-potential magnetic stresses force the large-scale magnetic topology to form short-lived, small-scale magnetic discontinuities in order to dissipate the excess energy in localized current sheets. This conjecture was initially proposed by \\citet{Parker88}, and subsequently it was modeled using 3-D MHD numerical simulations by many researchers \\citep[see for example the work of][]{NorGal96}. \\emph{We emphasize here that the concept of sudden formation of a distribution of unstable discontinuities inside a well organized large-scale topology has not been appreciated or used extensively enough to model the solar flare phenomenon}.  The scenario of spatially distributed, localized, small-scale dissipation, which moreover evolves in time, is not least supported by the observations which indicate highly fragmented energy dissipation and particle acceleration processes. There is strong evidence that narrow-band milli-second spike-emission in the radio range is directly associated to the primary energy release events. The emission itself of the radio-spikes is fragmented in space and time, as is seen in radio-spectrograms and in spatially resolved observations. It must thus be concluded that also the energy release process is fragmented in space and time, to at least the same degree as is the radio spike-emission \\citep[see][]{Benz03}. Also type III burst radio-emission, caused by electron beams escaping from flaring regions, exhibits fragmentation as a strong characteristic \\citep[e.g.][]{Benz94}. It is a too simple interpretation of the available data that the large scale structures seen have a relatively simple topology down to all scales.  One approach which is capable to capture the full extent of this interplay of highly localized dissipation in a well-behaved large scale topology ('sporadic flaring') is based on a special class of models which use the concept of Self-Organized Criticality \\citep[SOC;][]{Bak87}. The main idea is that active regions evolve by the continuous addition of new or the change of existing magnetic flux on an existing large scale magnetic topology, until at some point(s) inside the structure magnetic discontinuities are formed and the currents associated with them reach a threshold. This causes a fast rearrangement of the local magnetic topology and the release of the excess magnetic energy at the unstable point(s). This rearrangement may in-turn cause the lack of stability in the neighborhood, and so forth, leading to the appearance of flares (avalanches) of all sizes that follow a well defined statistical law \\citep{Lu91,Lu93,Vla95,Isl00,Isl01}, which agrees remarkably well with the observed flare statistics \\citep{cros}.    The modeling of the solar atmosphere cannot be done exclusively with the use of high-resolution MHD numerical codes. The small scale discontinuities easily give rise to kinetic instabilities and anomalous resistivity, which play a very important role, they dramatically change the evolution of even the large scale structures. On the other hand, numerical codes following the evolution of charged particles in idealized, non-evolving, large-scale current-sheets also fail to capture the spatio-temporal evolution of the magnetic energy dissipation. The extraordinary efficiency of particle acceleration during solar flares questions the use of ideal MHD for the description of flares, since it misses kinetic plasma effects, which yet play a major role in the energy dissipation process and for the local state of the plasma. The coupling of the large scale and the small scale is extremely difficult to handle, not only in solar physics but in physics in general, and it is the main reason for not having resolved the solar flare problem for so many years \\citep{Carg02}.  Our inability to describe properly the coupling between the MHD evolution and the kinetic plasma aspects of the driven flaring region is the main reason behind our lack of understanding of the mechanism(s) which causes the acceleration of high energy particles. Let us now define more accurately the so called `acceleration problem'. \\textit{We need to understand the mechanism(s) which accelerate electrons and ions in relatively large numbers to energies well above the relativistic regime on a short time scale with specific energy-spectra for the different isotopes and charge states}.  Let us summarize very briefly the main observational constraints for the acceleration processes in solar flares.   \\textbf{Electron} energies well in excess to 100 keV, and occasionally up to tens of MeV, are inferred. Electrons reach $100$keV in $<1$s and higher energies in a few seconds. The number of electrons required above 20 keV is large for x-class flares and can reach up to $10^{38}$ electrons per second, although this number is model dependent and not very accurate \\citep[see e.g.][]{Mill97}. The observed spectra in the hard X-rays can be fitted with single or double power-laws, combined often --- but not always --- with a thermal emission spectrum at the lowest energies \\citep{Hol03,Hol03b,Piana03}. The spectra in the microwaves roughly show a power-law decay at high frequencies \\citep{Hol03}.   \\textbf{Ions} (especially protons) are inferred to have energies above $1\\,$MeV and up to $1\\,$GeV per nucleon. The 10 MeV ions are accelerated on the same time scales as the electrons, with the high energy ions delayed up to $10\\,$s.  The total number of ions accelerated may carry the same energy as the electrons. The continuous component of the $\\gamma$-ray spectra are usually power laws. An important finding is the massive enhancement of $^3_2$He in very impulsive flares \\citep[see e.g.][]{Mill97}.  \\noindent \\textbf{Acceleration mechanisms:} Numerous books and reviews have been devoted to the challenging problem of particle acceleration \\citep{Hey,vla84,Mel,vla94,Kir,vla96,Kuip,Mill97}. The proposed models usually address parts of the problem, and almost all acceleration mechanisms have no clear connection to the large-scale topology and to the magnetic energy release mechanism(s). The most prominent mechanisms are shock waves \\citep{HolPes83,Bla87,Ell85,Dec88}, MHD turbulence \\citep{Fermi,Mill95}, and DC electric fields \\citep{BenHol92,Mo85,Mo90}.  \\noindent \\textbf{Mixing acceleration mechanisms:} Several inquiries have been made in which different acceleration mechanisms had been mixed.  \\citet{Dec} analyzed the role of Shock Drift Acceleration (SDA) when the shock is surrounded by a turbulent spectrum. As we already mentioned, the SDA is fast but not efficient since the particles drifting along the shock-surface's electric field quickly leave the shock. The presence of turbulence reinforces the acceleration process by providing a magnetic trap around the shock surface and forcing a particle to return many times to the shock surface. The particle leaves the shock surface, travels a distance $s_i$ inside the turbulent magnetic field, returns back to the shock surface with velocity $v_i$, drifts a distance $l_i$ along the shock electric field $E_{sc}$, changing its momentum by $\\Delta p_i\\sim E_{sc}\\cdot (l_i/v_i)$, it escapes again, travels a distance $s_{i+1}$ before returning back to the shock and drifting along the electric field, in other words, the acceleration-cycle has begun again. The process repeats itself several times before the particle gains enough energy to escape from the turbulent trap around the shock surface. Let us note some very important characteristics of this acceleration: (1) The distances $s_i$ traveled by the particle before returning back to the shock are only indirectly relevant to the acceleration, they basically delay the process and influence the overall timing, i.e.\\ the \\textit{acceleration time}, another important parameter of the particle acceleration process. (2) The energy gain depends critically on the lengths $l_i$ the particle drifts along the shock surface, in a statistical sense though, i.e.\\ on the distribution of the $l_i,\\,i=1,2,3\\,...$. (3) The times $\\tau_i$ a particle spends on the shock surface are again crucial for the energy gain, and also, together with the $s_i$, for the estimation of the acceleration time. (4) For the total acceleration problem, which concerns the energies reached and the times needed to reach them, all three variables, $s_i,\\,l_i,\\,\\tau_i$, are of equal importance.   \\citet{Amb} discussed a similar problem, placing a current sheet into a region of Alfv\\'enic turbulence. Also here, the ability of the associated DC electric field to accelerate particles is enhanced by the presence of the MHD turbulence. The acceleration process is again of a cyclic nature, as in the case of turbulent SDA, and the process is again characterized by the three variables $s_i,\\,l_i,\\,\\tau_i$. The turbulent current sheet has several avenues to enhance the acceleration-efficiency since the plasma inflow is dynamically driven and causes a variety of new and still unexplored phenomena. \\citet{Azn} analyzed the mixture of stationary MHD turbulence with a DC electric field. The trapping of the particles inside the turbulent magnetic field causes a new `collision scale', and, in some circumstances, acceleration becomes dependent on an alternative 'Dreicer field', in which particle collisions are replaced by collisions with magnetic irregularities. Actually also the diffusive shock acceleration described above is of a mixed type, having as elements a shock and magnetic turbulence, although turbulence plays a more passive role of just scattering the particles. It seems that most acceleration mechanisms are more or less of a mixed type.   We can conclude that the mixture of mechanisms enhances the acceleration-efficiency and removes some of the draw-backs attached to the different, isolated mechanisms. A second, main conclusion we draw is that \\textit{cyclic processes}, e.g.\\ through trapping around the basic accelerators, are important elements --- if not the presupposition --- of efficient and fast acceleration in space plasmas.  \\noindent \\textbf{Can the UCS naturally provide the unification of all the above mechanisms~?} Unstable Current Sheets (UCS) are the regions where magnetic energy is dissipated, and it is natural to ask if they can become the actual source of the energetic particles. According to the existing understanding of UCSs, several potential mechanisms for particle acceleration co-exist at a UCS. Plasma flows driving turbulence, shock waves, and DC electric fields are expected to appear simultaneously inside and around a driven and evolving UCS. If the UCS is located in the middle of a turbulent magnetic topology, all these phenomena will be enhanced and the sporadic external forcing of the plasma inflow into the UCS will create bursts of sporadic acceleration. The scenario of a single UCS currently enjoys very large popularity \\citep[e.g.][]{Litv03,FletMart98,Mart88}, and the question is: Can one single UCS be the answer to the acceleration problem~?  We believe that this is impossible since a single, isolated UCS must be enormously large $(10^{9}\\times 10^{9}\\times10^5$ m), remain stationary for a long time, and continuously accelerate particles with extreme efficiency in order to provide the required numbers of accelerated particles and the observed acceleration times. From the Earth's magnetic tail, it is known that large UCS break up quickly, creating a network of smaller scale UCS, with a specific probability distribution $P(l)$ of their characteristic scales, a process which just is a manifestation of turbulence \\citep{Angelo99}. The formation of a large scale helmet above certain loops, driven by erupting filaments cannot be excluded entirely and represents a special class of very energetic phenomena \\citep{Masu}. We though believe that such a current sheet breaks down on a very short time scale, and the formation of smaller scale current sheets will be unavoidable. Eventually, even in this very special occasion the acceleration probably takes place in an environment similar to the one discussed in this article.  There is a second reason,  based on a result on the statistics of flares, for questioning the idea of a single, large-scale UCS that is associated with just one or at most two loops, as in the scenario of the sometimes called 'standard model' \\citep[see e.g.][]{Shibata95}. \\citet{With00} determined the frequency-distributions in total emitted energy of flares occurring in the \\textit{same, individual} active region (number of flares per unit energy), and they find featureless power-law distributions, extending over many decades. If a single UCS would be the basic mechanism behind a flare, then the energy output must be expected to be related to the physical properties of the individual active region, e.g.\\ to the linear dimensions of the active region and the associated UCS, so that it is at least difficult to imagine how a single-site reconnection model could be able to produce a featureless energy distribution that extends over many decades.    \\noindent \\textbf{From single to multiple, small-scale UCS acceleration:} Our attempt in this article is to take advantage of the positive properties of isolated UCS as accelerators, but at the same time to assume that the dissipation happens at \\textit{multiple}, \\textit{small-scale} sites.  The 3-D magnetic topology, driven from the convection zone, dissipates energy in localized UCS, which are spread inside the coronal active region, providing a natural fragmentation for the energy release and a multiple, distributed accelerator. In this way, the magnetic topology acts as a host for the UCS, and the spatio-temporal distribution of the latter defines the type of a flare, its intensity, the degree of energization and acceleration of the particles, the acceleration time-scales etc. Evolving large-scale magnetic topologies provide a variety of opportunities for acceleration which is not restricted to flares, but can also take place before a flare, and after a flare, being just the manifestation of a more relaxed, but still driven topology. Depending thus on the level to which the magnetic topology is stressed, particles can be accelerated without a flare to happen, and even long-lived acceleration in non-flaring active regions must be expected to occur. Consequently, the starting point of the model to be introduced below is a driven 3-D magnetic topology, which defines a time-dependent spatial distribution of UCS inside the active region. In this article, we will focus on the interaction of the UCS with particles. The details of the mechanisms involved in the acceleration of particles inside the UCS are not essential in a stochastic modeling approach.  Ideas related to the basic approach of this article have been presented earlier \\citep{Ana91,vla93,Ana94,Ana97}, but remained conceptually simpler, they left many points open for future development, being not that general and complete as approaches as the model we present here. In the approach we follow here, we make extensive use of new developments in the theory of SOC models for flares. Also taken into account are ideas from the theory of Complex Evolving Networks \\citep{Alb,Don}, adjusted though to the context of plasma physics: The spatially distributed, localized UCS can be viewed as a network, whose 'nodes' are the UCS themselves, and whose 'edges' are the possible particle trajectories between the nodes (UCS). The particles are moving around in this network, forced to follow the edges, and undergoing acceleration when they pass by a node (see also Fig.\\ \\ref{plottwo}). The network is complex in that it has a non-trivial spatial structure, and it is evolving since the nodes (UCS) are short-lived, as are the connectivity channels, which ever change during the evolution of a flare. This instantaneous connectivity of the UCS is an important parameter in our model, it determines to what degree multiple acceleration is imposed onto the system, which in turn influences the instantaneous level of energization and the acceleration time-scale of the particles.  In this article, we introduce a stochastic, multiple acceleration model for solar flare particle acceleration. We assume an ensemble of unstable current sheets (UCS), distributed in space in such a way that it reflects a relatively simple large scale topology in which turbulence at small scales is developing. Concretely, the UCS in their ensemble form a fractal, as it is the case in the SOC models \\citep{Isl03b,Mc02}. Particles move erratically around, and occasionally enter a UCS in which they undergo acceleration. The acceleration process is taken as a simple DC electric field mechanism. The basic approach of the model is a combined random walk in position- and momentum-space. The frame-work we introduce is such that potentially all multiple acceleration models can be formulated in this frame-work, as long as the acceleration takes place in spatially localized, disconnected, small-scale regions. Our main interest here is in the diffusion, acceleration, and heating of the electrons. We will follow the evolution of the kinetic energy distribution in time and calculate from these distributions the hard X-ray (HXR) and microwave spectra, which can be compared to the observations. Moreover, we will shortly inquire the case of ion heating and acceleration.  The basic elements of our model are presented in Section 2. In section 3, we present some of our results, and in section 4 we discuss our findings and propose ways for continuing the exploration of the problem along the path of our modeling approach. A conclusion is presented in section 5. ",
        "conclusions": "In this article, we  shifted the emphasis of the acceleration process in active regions away from the details of specific mechanism(s) involved and focused on the global aspects of the active region, its evolution in space and time, and on the stochastic nature of the acceleration process. Our basic assumptions are: (1) acceleration is a local process in and near the UCS, (2) the UCS are distributed in a complex way inside the large-scale 3-D magnetic topology of the active region, and (3) acceleration is the result of multiple interactions with different UCS. These assumptions can be summarized in the statement that energy dissipation and particle acceleration in flares are fragmented. We achieved for the first time to connect the accelerator to the 3-D magnetic topology and the energy release process. From our results, we can safely conclude that the complexity of the 3-D magnetic field topology in active regions, in combination with the stresses imposed by the convection zone onto it, forms a highly efficient accelerator.  The adequate tool for modeling the stochastic nature of particle acceleration in flares is continuous time random walk in position- and velocity-space. This new approach opens up the way for the understanding of a variety of acceleration phenomena, e.g.\\ acceleration before and after a flare, acceleration without a flare, long lasting acceleration etc.  We feel that the road opened in this article is new and still unexplored in astrophysics. As mentioned, it can provide answers to a number of open problems which have remained unsolved when it was attempted to use one UCS (above or inside a flaring loop) and one (uncorrelated with the energy release) acceleration mechanism to explain phenomena of acceleration. Many new questions are still open and we plan to return to these issues in a forth-coming article."
    },
    "0402/astro-ph0402535_arXiv.txt": {
        "abstract": "In the SDSS, several unusual QSOs have been discovered that have very blue rest-frame UV spectra but no discernible emission lines.  Their UV spectra strongly resemble that of the newly discovered quasar PHL 1811 ($z=0.192$; $\\rm M_{V}=-25.9$).  With magnitudes of $B = 14.4$ and $R = 14.1$, PHL 1811 is the second brightest quasar known with $z>0.1$ after 3C 273.  Optically it is classified as a Narrow-line Seyfert 1 galaxy (NLS1).  Objects of this class are generally strong soft X-ray emitters, but a BeppoSAX observation of PHL 1811 showed that it is anomalously X-ray weak.  The inferred $\\alpha_{ox}$ was 1.9--2.1, much steeper than the nominal value of 1.6 for quasars of this optical luminosity, and comparable to the X-ray weakest quasars.  Follow-up {\\it Chandra} observations reveal a variable, unabsorbed X-ray spectrum and confirm that it is intrinsically X-ray weak.  {\\it HST} STIS spectra of PHL 1811 reveal a very blue continuum with little evidence for absorption or scattering intrinsic to the quasar. High-ionization lines are very weak; \\ion{C}{IV} has an equivalent width of only $\\sim 5$\\AA\\/.  Neither forbidden nor semiforbidden emission lines are detected. \\ion{Fe}{II} is the dominant line emission in the UV.  High metallicity is implied by the large \\ion{Fe}{II} to \\ion{Mg}{II} ratio and relatively strong \\ion{N}{V}. Low-ionization emission lines of \\ion{Al}{III}, \\ion{Na}{I} D, and \\ion{Ca}{II} H \\& K are present, implying high optical depth.  We demonstrate that the emission-line properties of PHL 1811 are a direct consequence of the UV-peaked continuum and weak X-ray emission. We propose that these properties are a consequence of high accretion rate, which powers the UV emission from an optically thick accretion disk, while suppressing the formation of a hot corona.  This is an extreme case of the same mechanism which is thought to be responsible for luminous NLS1s.  Based on the similarity between PHL 1811 and the lineless SDSS quasars, we propose that the lineless quasars discovered in the SDSS are the high-z counterparts of local high-luminosity NLS1s. ",
        "introduction": "In the Sloan Digital Sky Survey, a number of high redshift quasars have been discovered that are remarkable for their lack of emission lines.  A detailed study of one of these objects, SDSSp~J153259.96$-$003944.1 (hereafter referred to as SDSS~J1532$-$0039) was presented by Fan et al.\\ 1999.  The optical spectrum, redward of 6800 \\AA\\/, is a featureless blue power-law continuum; there are no emission lines.  Blueward of 6800 \\AA\\/, there are features due to Ly$\\alpha$ absorption that lead to a redshift estimate of $4.52$. Featureless continua are frequently found in Bl Lac objects. However, SDSS~J1532$-$0039 does not appear to have other characteristic properties of Bl~Lac objects: it was not detected in a deep radio observation, doesn't vary in the optical, and was not found to be optically polarized.  Follow-up observations with {\\it Chandra} failed to detect the quasar, indicating an $\\alpha_{ox}>1.74$ \\footnote{$\\alpha_{ox}$ is defined as the point-to-point slope between 2500\\AA\\/ and 2 keV.}. This steep $\\alpha_{ox}$ makes it moderately X-ray weak (Vignali et al.\\ 2001).  PHL 1811 is a luminous narrow-line quasar that was rediscovered in the VLA FIRST survey (Leighly et al.\\ 2000).  It is a very bright ($B=14.4$, $R=14.1$), nearby ($z=0.192$) object, well away from the Galactic plane, so it was surprising that it was not detected in the ROSAT All Sky Survey. Two {\\it Chandra} observations show that it is X-ray weak, with inferred $\\alpha_{ox}$ between $1.9$ and $2.1$, much steeper than for a typical quasar of this luminosity ($1.6$; Wilkes et al.\\ 1994). The {\\it HST} spectrum is unusual in that it is dominated by \\ion{Fe}{II}, and there are no prominent broad emission lines.  We thought that the X-ray weakness and the lack of broad emission lines suggest a similarity to SDSS~J1532$-$0039. ",
        "conclusions": "In this contribution, we discuss the similarity between the UV spectrum of PHL 1811 and the spectra of several high-z lineless quasars discovered in the SDSS, and propose that the physical cause of their spectra is the same.  At the same time, we acknowledge that the SDSS is quite likely to find lineless quasars whose spectra have a different origin; for example, they may be Bl~Lac objects.  If it is true that these spectra are a result of the same phenomenon, what are the potential implications?  Because of their steep X-ray spectra and high amplitude X-ray variability, it has been proposed that NLS1s are characterized by a high accretion rate (e.g., Leighly 1999 and references therein).  Since emission-line widths should depend on the black hole mass (Laor 2000), luminous NLS1s should have exceptionally high accretion rates compared with their Eddington value.  It is important to understand the accretion rate in high redshift objects, as luminous quasars have large black holes, and in order to grow large in the short amount of time implied by the high redshift, they should be accreting at a rapid rate.  Naturally, people are searching for a connection between high redshift quasars and NLS1s (e.g., Mathur 2000), although the evidence so far is mixed (Constantin et al.\\ 2002).  Based on the similarities between PHL 1811 and the high-z lineless quasars, we suggest that they are the early Universe counterparts of luminous Narrow-line Seyfert 1 galaxies.  What causes the weak emission lines?  In PHL 1811, the higher-ionization emission lines are probably weak because the continuum is very soft overall.  Such a soft continuum can produce copious H$^0$ and Fe$^+$ ions, but few higher ionization species. Also, the equivalent widths of the lines will appear small against the strong, blue optical/UV continuum."
    },
    "0402/astro-ph0402229_arXiv.txt": {
        "abstract": "We present data from a {\\it Chandra} observation of the nearby cluster of galaxies Abell 576.  The core of the cluster shows a significant departure from dynamical equilibrium.  We show that this core gas is most likely the remnant of a merging subcluster, which has been stripped of much of its gas, depositing a stream of gas behind it in the main cluster.  The unstripped remnant of the subcluster is characterized by a different temperature, density and metalicity than that of the surrounding main cluster, suggesting its distinct origin.  Continual dissipation of the kinetic energy of this minor merger may be sufficient to counteract most cooling in the main cluster over the lifetime of the merger event. ",
        "introduction": "\\label{sec:intro}  The study of the cores of clusters of galaxies has undergone a renaissance in the past few years with the launch of the {\\it Chandra} and {\\it XMM-Newton} observatories.  While previous observatories lacked the spatial resolution necessary to resolve structure within the cores of clusters, this new generation of telescopes has revealed an astonishing level of complexity in the structure of the intracluster medium (ICM).  Many clusters previously thought to be relaxed, regular systems have proven to be far from dynamical equilibrium, particularly in their cores \\citep*[e.g.][]{mvm01,mem03}.  Abell 576 is no exception.  Earlier data, especially optical spectra of the cluster's galaxy population, hinted at dynamical complexity in the cluster core \\citep{mgf+96}, but earlier X-ray observations showed the cluster to be quite regular and even cooler in the center, suggesting either a cooling flow or multi-phase gas with a very cool component \\citep{rvm+84,mgf+96}.  With its low redshift, Abell 576 makes excellent use of the capabilities of {\\it Chandra}, allowing us to examine in detail the very core of the cluster.  In this paper, we focus on the dynamical activity in the core of cluster.  The cluster shows strong evidence, first suggested by \\citet{mgf+96} from an analysis of the galaxy population, of the remnant core of a small merged subcluster.  We demonstrate that the X-ray data are consistent with this picture, and even suggest it as the most likely origin for the non-equilibrium gas at the center of Abell 576.  In fact, the subcluster may still be in the process of settling into the center of the main cluster's potential.  Throughout this paper, we use the cosmological parameters derived from the first release WMAP results \\citep{bhh+03}, so $1\\arcsec=0.738$ kpc at $z=0.0377$.  All errors are quoted at 90\\% confidence unless otherwise stated. ",
        "conclusions": "\\label{sec:conclusions}  While earlier X-ray observations of Abell 576 have shown it to be quite regular on large scales, we have demonstrated that the core of the cluster is far from dynamical equilibrium.  We found multiple surface brightness edges in the cluster center which we have demonstrated to be indicative of jumps in both density and abundance.  Of the two most likely explanations for the existence of these edges, our analysis favors the hypothesis that they are formed by gas stripped from a merging subcluster.  Most of the gas appears to have been stripped from the subcluster, leaving a core only $\\sim$30 kpc in radius.  The stripped gas has been found to have both a lower temperature and a higher abundance than the gas in the rest of Abell 576.  We find no evidence of gas cooling from the ambient temperature of the main cluster, but do find some suggestion of very cool gas at a temperature expected of gas that had condensed out of the ICM of the subcluster.  The simple cooling rate derived from the gas mass and cooling time is an order of magnitude larger than the spectroscopically measured $\\dot{M}$.  We have demonstrated that dissipation of the kinetic energy of the observed remnant core of an infalling subgroup may be sufficient to reduce cooling to the observed rate, if that energy is dissipated over a timescale of $\\lesssim 10^9$ years."
    },
    "0402/astro-ph0402065_arXiv.txt": {
        "abstract": "I discuss the classical cosmological tests-- angular size-redshift, flux-redshift, and galaxy number counts-- in the light of the cosmology prescribed by the interpretation of the CMB anisotropies.  The discussion is somewhat of a primer for physicists, with emphasis upon the possible systematic uncertainties in the observations and their interpretation.  Given the curious composition of the Universe inherent in the emerging cosmological model,  I stress the value of searching for inconsistencies rather than concordance, and suggest that the prevailing mood of triumphalism in cosmology is premature. ",
        "introduction": "The traditional cosmological tests appear to have been overshadowed by observations of the anisotropies in the cosmic microwave background (CMB).  We are told that these observations accurately measure the geometry of the Universe, its composition, its present expansion rate, and the nature and form of the primordial fluctuations \\cite{speal03}. The resulting values for these basic parameters are very similar to those deduced earlier from a variety of observations-- the so-called ``concordance model''-- with about 30\\% of the closure density of the Universe comprised of matter (mostly a pressureless, non-baryonic dark matter), the remainder being in negative pressure dark energy \\cite{osst95}. Given the certainty and precision of these assertions, any current discussion of observational cosmology must begin with the question:  Is there any room for doubt?  Why should we bother with lower precision cosmological tests when we know all of the answers anyway?  While the interpretation of the CMB anisotropies has emerged as the single most important cosmological tool, we must bear in mind that the conclusions drawn do rest upon a number of assumptions, and the results are not altogether as robust as we are, at times, led to believe. One such assumption, for example, is that of adiabatic initial fluctuations-- that is, 100\\% adiabatic.  A small admixture of correlated isocurvature fluctuations, an aspect of braneworld scenarios \\cite{maart}, can affect peak amplitudes and thus, the derived cosmological parameters.  A more fundamental assumption is that of the validity of traditional Friedmann-Robertson-Walker (FRW) cosmology in the post-decoupling universe.  Is the expansion of the universe described by the Friedmann equation?  Even minimal changes to the right-hand-side, such as the equation of state of the dark energy component, can alter the angular size distance to the last scattering surface at z=1000 and the luminosity distance to distant supernovae. But even more drastic changes to the Friedmann equation, resulting from modified gravitational physics, have been proposed in attempts to remove the unattractive dark energy \\cite{def01,ceal}.  Such suggestions reflect a general unease with the concordance model-- a model that presents us with a universe that is strange in its composition.  The most abundant form of matter consists of, as yet, undetected non-baryonic particles originally postulated to solve the problems of structure formation and of the missing mass in bound gravitational systems such as galaxies and clusters of galaxies. In this second respect, it is fair to say that it has failed-- or, to be generous, not yet succeeded-- because the predicted density distribution of dark halos which emerge from cosmic N-body \\cite{nfw} simulations appears to be inconsistent with observations of spiral galaxies \\cite{mdb98} or with strong lensing in clusters of galaxies \\cite{treal}.  Even more mysterious is the ``dark energy'', the pervasive homogeneous fluid with a negative pressure which may be identified with the cosmological constant, the zero-point energy density of the vacuum.  The problem of this unnaturally low energy density, $10^{-122}$ in Planck units, is well-known, as is the cosmic coincidence problem:  why are we observing the Universe at a time when the cosmological constant has, fairly recently, become dynamically important \\cite{car01}? To put it another way, why are the energy densities of matter and dark energy so comparable at the present epoch? This is strange because the density of matter dilutes with the expanding volume of the Universe while the vacuum energy density does not.  It is this problem which has led to the proposal of dynamic dark energy, quintessence-- a dark energy, possibly associated with a light scalar field-- with an energy density that evolves with cosmic time possibly tracking the matter energy density \\cite {rapeb}.  Here the difficulty is that the field would generally be expected to have additional observational consequences-- such as violations of the equivalence principle at some level, possibly detectable in fifth force experiments \\cite{car01}.  For these reasons, it is even more important to pursue cosmological tests that are independent of the CMB, because one might expect new physics to appear as observations inconsistent with the concordance model. In this sense, discord is more interesting than concord; to take a Hegelian point of view-- ideas progress through dialectic, not through concordance.  It is with this in mind that I will review observational cosmology with emphasis upon CMB-independent tests.  Below I argue that the evolution of the early, pre-recombination universe is well-understood and tightly constrained by considerations of primordial nucleosynthesis.  If one wishes to modify general relativity to give deviations from Friedmann expansion, then such modifications are strongly constrained at early times, at energies on the order of 1 MeV.  However, cosmological evolution is much less constrained in the post-recombination universe where there is room for deviation from standard Friedmann cosmology and where the more classical tests are relevant.  I will discuss three of these classical tests: the angular size distance test where I am obliged to refer to its powerful modern application with respect to the CMB anisotropies; the luminosity distance test and its application to observations of distant supernovae; and the incremental volume test as revealed by faint galaxy number counts.  These classical tests yield results that are consistent, to lower precision, with the parameters deduced from the CMB.  While one can make minimal changes to standard cosmology, to the equation of state of the dark energy for example, which yield different cosmological parameters, there is no compelling observational reason to do so.  It remains the peculiar composition and the extraordinary coincidences embodied by the concordance model that call for deeper insight. Such motivations for questioning a paradigm are not unprecedented;  similar worries led to the inflationary scenario which, unquestionably, has had the dominant impact on cosmological thought in the past 25 years and which has found phenomenological support in the recent CMB observations.  I am not going to discuss cosmological tests based upon specific models for structure formation, such as the form of the luminous matter power spectrum \\cite{peak} or the amplitude of the present mass fluctuations \\cite{hoek}. I do not mean to imply that such such tests are unimportant, it is only that I restrict myself here to more global and model-independent tests.  If one is considering a possibility as drastic as a modification of Friedmann expansion due, possibly, to new gravitational physics, then it is tests of the global curvature and expansion history of the Universe that are primary.  I am also going to refrain, in so far as possible, from discussion of theory-- of new gravitational physics or of any other sort.  The theoretical issues presented by dark matter that can only be detected gravitationally or by an absurdly small but non-zero cosmological constant are essentially not problems for the interpretive astronomer. The primary task is to realistically access the reliability of conclusions drawn from the observations, and that is what I intend to do. ",
        "conclusions": "In these lectures I have been looking for discord, but have not found it. The classical tests return results for cosmological parameters that are consistent with but considerably less precise than those implied by the CMB anisotropies,  given the usual assumptions. It is fair to say that the numbers characterizing the concordance model, $\\Omega_m \\approx 0.3$, $\\Omega_\\Lambda \\approx 0.7$ are robust {\\it in the context of the framework of FRW cosmology}.  It is, in fact, the peculiar composition of the Universe embodied by these numbers which calls that framework into question.  Rather small changes in the assumptions underlying pure FRW cosmology (with only an evolving vacuum energy density in addition to more familiar fluids) can make a difference.  For example, allowing $w=-0.6$ brings the number counts and z-distribution of faint galaxies into agreement with a Universe strongly dominated by dark energy ($\\Omega_Q = 0.9$).  The same also true of the high-z supernovae observations \\cite{tonry}).  Allowing a small component of correlated iso-curvature initial perturbations, as expected in braneworld cosmologies, can affect the amplitudes and positions of the peaks in the angular power spectrum of the CMB anisotropies \\cite{maart}, and therefore the derived cosmological parameters.  But even more drastic changes have been suggested.  Certain braneworld scenarios, for example, in which 4-D gravity is induced on the brane \\cite{dvali} imply that gravity is modified at large scale where gravitons begin to leak into the bulk \\cite{deff}. It is possible that the observed acceleration is due to such modifications and not to dark energy.  More ad hoc modifications of General Relativity \\cite{ceal} have also been proposed because of a general unease with dark energy-- proposals whereby gravity is modified in the limit of small curvature scalar.  My own opinion is that we should also feel uneasy with the mysterious non-baryonic cold dark matter, because the only evidence for its existence, at present, is its gravitational influence; when the theory of gravity is modified to eliminate dark energy, it might also be found that the need for dark matter vanishes.  In general,  more attention is being given to so-called infrared modifications of gravity (e.g. \\cite{arkan}), and this is a positive development.  High energy modifications, that affect the evolution of the early Universe, are, as we have seen, strongly constrained by considerations of primordial nucleosynthesis (now, in combination with the CMB results).  It is more likely that modifications play a role in the late, post-recombination evolution of the Universe, where the peculiarities of the concordance model suggest that they are needed.  The fact that the same rather un-natural values for the comparable densities of dark energy and matter keep emerging in different observational contexts may be calling attention to erroneous underlying assumptions rather than to the actual existence of these ``ethers''.  Convergence toward a parameterized cosmology is not, without deeper understanding, sufficient reason for triumphalism.  Rather, it should be a motivation to look more carefully at the possible systematic effects in the observations and to question more critically the underlying assumptions of the models. \\parskip 5truemm \\parindent 0pt  I thank Rien van de Weygaert, Ole M\\\"oller, Moti Milgrom, Art Wolfe, Jacob Bekenstein, and Scott Trager for useful comments on the manuscript.  I also thank Gary Steigman, Wendy Freedman, and Luis Ho for permission to use Figs.\\ 1 and 2. I am very grateful to the organizers of the Second Aegean Summer School on the Early Universe, and especially, Lefteris Papantonopoulos, for all their work and for inviting me to the very pleasant island of Syros."
    },
    "0402/astro-ph0402586_arXiv.txt": {
        "abstract": "{ In this paper, we have made an accurate investigation of proton acceleration in GRBs and we have predicted a possible signature of cosmic rays, in a sufficiently baryon-loaded fireball, via GeV $\\gamma$-ray emission produced by $\\pi^{0}$-meson decay. If two ungrounded assumptions are removed, namely, Bohm's scaling and a slow magnetic field decrease, the usual Fermi processes are unable to generate ultra high energy cosmic rays (UHECRs) in GRBs. We propose to develop another scenario of relativistic Fermi acceleration in the internal shock stage. We present the results of a realistic Monte-Carlo simulation of a multi-front acceleration which clearly shows the possible generation of UHECR. The amount of energy converted into UHECRs turns out to be a sizeable fraction of the magnetic energy. ",
        "introduction": "Gamma Ray Bursts (hereafter GRBs) have been considered as promising sources of Ultra High Energy Cosmic Rays (UHECRs) either through the external ultrarelativistic shock \\citep{Vietri95} or through the internal shocks \\citep{Waxman95}. The acceleration at the external shock does not achieve this goal when the external medium is a standard interstellar medium \\citep{Gallant1999}. As for the internal shocks, first and second order Fermi processes would achieve this goal with the extreme assumption of a cosmic ray mean free path of the order of its Larmor radius (so-called Bohm's scaling). In a previous paper \\citep{GialisPelletier03}, we have shown that such scaling is physically inconsistent for two reasons; first, the law is not validated by numerical simulations \\citep{Casse01} and second, even if inappropriately used, it would lead to tremendously fast magnetic energy depletion in the fireball. We also have considered a more realistic Kolmogorov scaling but a severe expansion loss limitation does not make it possible. These arguments are developed in Sects. \\ref{sec2} and \\ref{sec3} of this paper, together with the important issue of the magnetic field profile. Because both the global energy budget of the fireball and the Hillas criterium are in favour of the generation of UHECRs, we have made a detailed investigation of a different Fermi process proposed by \\cite{PelletierK00} and \\cite{Pelletier99}. In this kind of Fermi process, reviewed in Sect. \\ref{sec4}, the sheets invoked in the internal shock model are considered as relativistic hydromagnetic fronts that scatter cosmic rays. Thus, cosmic rays undergo a relativistic Fermi acceleration, similar to a kind of second order process through multiple interactions with all the fronts in the relativistic wind. Actually, this is not a second order Fermi process because no expansion is allowed in the relativistic regime, the energy jump being large at each scattering. The generation of cosmic rays resulting from these multiple scatterings is very different from generation obtained by a superposition of the individual first order contributions of each shock.  \\\\ In Sect. \\ref{sec5} we describe a realistic Monte-Carlo simulation of this process, taking into account several dynamical parameters in a conical expanding fireball, for several intensities and profiles of the average magnetic field. The results are presented in Sect. \\ref{sec6}; they suggest a good efficiency of the process of acceleration of UHECRs and give a prediction of the time at which they are created during the internal shock stage. Finally, we establish a possible diagnosis of cosmic rays in GRBs through $\\gamma$-ray emission by $\\pi^{0}$-meson decay which would be easily detected by HESS observatory, 5@5 experiment and GLAST. ",
        "conclusions": "We have shown through accurate investigation of scattering properties that the usual Fermi processes are unable to generate UHECRs in GRBs. First of all, it is well known that the highly relativistic external shock with a standard interstellar medium cannot achieve the goal of UHECR generation. The hope that it could do so was mostly based on the internal shock model but the irrelevant Bohm's scaling was used; moreover, a slowly decreasing magnetic field was assumed. Removing these extreme assumptions made the goal impossible \\citep{GialisPelletier03}. \\\\ However, because the global energy budget is in favour of copious cosmic ray generation, we thought it interesting to design an additional Fermi process that would stretch the cosmic ray distribution tail. Indeed, the acceleration through multi-scattering off magnetized relativistic fronts turns out to be more efficient even with a field decreasing in $r^{-2}$. In this paper, we found that a sizeable fraction ($\\sim 10^{-5}$) of the cosmic ray population reaches the UHE range, which is even more than necessary. \\\\ There is another important physical issue, namely the escape of the UHECRs from the fireball towards the observer. The relativistic Fermi process that we proposed generates UHECRs in a thin layer located at the very beginning of the internal shock phase. The question arises whether these particles suffer a further expansion loss. In fact, because the magnetic field decreases faster than $r^{-1}$, the interaction energy limit with the magnetic fronts decreases and the UHECRs are no longer scattered by them. Thus, because of the lack of scattering, they cannot experience the expansion loss, contrarily to models expecting $B\\propto r^{-1}$. Therefore, they directly travel across the shells, benefiting of the Lorentz boost. For a magnetic field decreasing as $r^{-2}$, the integrated spectrum is in $\\epsilon^{-2}$ over 4 decades.\\\\ In our previous paper \\citep{GialisPelletier03}, we predicted a possible double neutrino emission;  even though this improvement of the acceleration process does not significantly change the estimated fluxes, the collimation allows detectable events, especially with pp-neutrinos. However, we found a possible diagnosis of cosmic ray generation in some numerous class of GRBs, having $\\eta < \\eta_{\\star}/2$, associated with the signature of $\\pi^{0}$-decay in the GeV $\\gamma$-ray spectrum. This seems to be really observable, with a few percent of the GRB energy, for GRBs having a sufficient baryon load. This signature  (in the GeV range by HESS observatory, 5@5 experiment and GLAST) would be easily distinguished from the synchrotron bump and the SSC-bump with a flux as high as a few percent of the fireball global energy flux.\\\\ This study gives a new chance to GRBs as sources of UHECRs. But their contribution to the diffuse background spectrum might not be dominant compared to the AGN contribution. This is an important issue that is under intense debate in the literature \\citep{Stecker2000,Berezinsky03,BahcallWaxman03}."
    },
    "0402/astro-ph0402253_arXiv.txt": {
        "abstract": "In this paper we will discuss charged stars with polytropic equation of state, where we will try to derive an equation analogous to the Lane-Emden equation. We will assume that these stars are spherically symmetric, and the electric field have only the radial component. First we will review the field equations for such stars and then we will proceed with the analog of the Lane-Emden equation for a polytropic Newtonian fluid and their relativistic equivalent (Tooper, 1964)\\cite{1}. These kind of equations are very interesting because they transform all the structure equations of the stars in a group of differential equations which are much more simple to solve than the source equations. These equations can be solved numerically for some boundary conditions and for some initial parameters. For this we will assume that the pressure caused by the electric field obeys a polytropic equation of state too. ",
        "introduction": "The metric used here is the usual Schwarzschild metric, \\begin{equation} ds^2=e^{\\nu(r)}c^2dt^2-e^{\\lambda(r)} dr^2-r^2(d\\theta^2+sen^2\\theta d\\phi^2) \\end{equation} The time-independent gravitational field equations reduces to: \\begin{eqnarray} e^{-\\lambda}(-\\frac{1}{r^{2}}+\\frac{1}{r}\\frac{d\\lambda}{dr})+\\frac{1}{r^{2}}=-\\frac{8\\pi \\label{fe1} G}{c^{4}}(\\rho c^{2}+ \\frac{E^{2}}{8\\pi}) \\\\ e^{-\\lambda}(\\frac{1}{r}\\frac{d\\nu}{dr}+\\frac{1}{r^{2}})-\\frac{1}{r^{2}}=\\frac{8\\pi G}{c^{4}}(p-\\frac{E^{2}}{8\\pi}) \\label{fe2} \\end{eqnarray} The mixed energy-momentum tensor will include the terms of the Electromagnetic Field, giving us an equation of the form \\begin{equation} T_{\\nu}^{\\mu}=diag(-\\rho c^2-\\frac{E^2}{8\\pi},p-\\frac{E^2}{8\\pi},p+\\frac{E^2}{8\\pi},p +\\frac{E^2}{8\\pi}), \\label{emt} \\end{equation} where $p$ is the pressure, $\\rho$ the mass density and $E$ the radial electric field. The four-divergence of $T_{\\nu}^{\\mu}$ should vanish for it being a conserved quantity and so we get \\begin{equation} \\frac{dp}{dr} = -\\frac{1}{2}\\frac{d\\nu}{dr}(p+\\rho c^2)+\\frac{E}{8\\pi}(\\frac{dE}{dr}+\\frac{2E}{r}).\\label{eq:dp} \\end{equation} This is the Tolman-Oppeheimer-Volkoff  (TOV) equation for a spherically symmetric charged fluid. This equation represents the hydrostatic equilibrium for the fluid. The first term of the right hand side (r.h.s.) of equation (\\ref{eq:dp}) is derived from the normal matter-energy, but the second term is new. It represents the contribution from the coulombian force and the energy from the electric field. If we look at the energy-momentum tensor (equation (\\ref{emt})), we see that the effective pressure and density in the fluid is given for \\begin{equation} p_{ef}=p-\\frac{E^2}{8\\pi}, \\quad \\rho_{ef}=\\rho+\\frac{E^2}{8\\pi} \\end{equation}  Now the energy-momentum tensor takes the following form: \\begin{equation} T_{\\nu}^{\\mu}=diag(-\\rho_{ef} c^2,p_{ef},p+\\frac{E^2}{8\\pi},p +\\frac{E^2}{8\\pi}) \\end{equation}   We could have written the last two terms as function of the effective pressure and density, but we didn't because it is not advantageous . Now we have the modified TOV equation: \\begin{equation} \\frac {dp_{ef}} {dr}= -\\frac{1}{2}\\frac{d\\nu}{dr}(p_{ef}+\\rho_{ef} c^2)+\\frac{E^2}{2\\pi r}  \\label{TOVm} \\end{equation} ",
        "conclusions": ""
    },
    "0402/astro-ph0402123_arXiv.txt": {
        "abstract": "We describe results from time-dependent numerical modeling of the collisionless reverse shock terminating the pulsar wind in the Crab Nebula. We treat the upstream relativistic wind as composed of ions and electron-positron plasma embedded in a toroidal magnetic field, flowing radially outward from the pulsar in a sector around the rotational equator. The relativistic cyclotron instability of the ion gyrational orbit downstream of the leading shock in the electron-positron pairs launches outward propagating magnetosonic waves. Because of the fresh supply of ions crossing the shock, this time-dependent process achieves a limit-cycle, in which the waves are launched with periodicity on the order of the ion Larmor time. Compressions in the magnetic field and pair density associated with these waves, as well as their propagation speed, semi-quantitatively reproduce the behavior of the wisp and ring features described in recent observations obtained using the {\\it Hubble Space Telescope} and the {\\it Chandra X-Ray Observatory}.  By selecting the parameters of the ion orbits to fit the spatial separation of the wisps, we predict the period of time variability of the wisps that is consistent with the data. When coupled with a mechanism for non-thermal acceleration of the pairs, the compressions in the magnetic field and plasma density associated with the optical wisp structure naturally account for the location of X-ray features in the Crab. We also discuss the origin of the high energy ions and their acceleration in the equatorial current sheet of the pulsar wind. ",
        "introduction": "The mechanisms through which compact objects excite surrounding synchrotron nebulae are one of the long-standing problems of high-energy and relativistic astrophysics. The loss of rotational energy from the central magnetized object underlies the synchrotron emission observed from pulsar wind nebulae (PWNs; e.g., Bandiera, Amato, \\& Woltjer 1998) and is one of the prime candidates for the energization of the jets in active galactic nuclei, radio galaxies (see Krolik 1999), and gamma ray bursts (e.g., Blandford 2002). PWNs form the nearest at hand laboratories for the investigation of the physics of such high-energy particle acceleration, and are the example where rotational energization clearly bears responsibility for the nonthermal activity.  Despite intensive study, however, the physics of how the rotational energy gets converted to the observed synchrotron emission has not been satisfactorily understood. Magnetohydrodynamic wind models (Kennel and Coroniti 1984a,b and many subsequent efforts) provide an adequate macroscopic description of the ``bubbles'' of synchrotron-emitting particles and fields surrounding young ($t \\sim 10^3$ yr) rotation-powered pulsars, but identification and quantitative modeling of the mechanisms that convert the unseen energy flowing out from the central pulsar into the visible synchrotron emission have remained elusive.  Most attention has focused on the idea that the outflow is a supersonic and super-Alfvenic wind, perhaps with embedded large-amplitude electromagnetic waves (Rees and Gunn 1974; Melatos and Melrose 1996). Kennel and Coroniti's model simplified the picture by assuming the wind to be steady, with no wave structure. In that case, most models attribute the conversion of relativistic outflow energy to the nonthermal particle spectra that emit the observed synchrotron radiation in the body of the PWN to a relativistic collisionless shock wave that terminates the flow. This idea was originated by Rees and Gunn, who also observed that the shock should form at a radius where the dynamic pressure of the outflow is approximately equal to the pressure of the surrounding nebula.  They also noted that in the case of the Crab Nebula, this radius approximately corresponds to the location of the moving ``wisps'' discovered 80 years ago (Lampland 1921) in optical observations of that PWN. This identification has motivated many studies of the wisps, as possible direct manifestations of the shock wave thought to underlie the transformation of outflow energy into the nonthermally emitting particles in the nebula.  Shortly after the discovery of the Crab pulsar, Scargle (1969) described these motions in more detail and attempted to interpret them as wave phenomena in the nebula stimulated directly by glitches in the pulsar's rotational spin-down (Scargle \\& Harlan 1970; Scargle \\& Pacini 1971). Subsequent observations made clear that the association of wisp activity to pulsar rotational activity was a spurious consequence of undersampling the variability of the wisp motions.  The advent of the magnetohydrodynamic model has focused attention on the wisps' structure and variability as a probe of the physics of the region where the pulsar's outflow interacts with the nebular plasma.  Substantially improved observations in the optical (Hester et al. 1995; Tanvir, Thomson, \\& Tsikarishvili 1997; Hester 1998a; Hester et al., 2002), radio (Bietenholtz, Frail, \\& Hester 2001) and X-ray (Mori et al. 2002; Hester et al. 2002) now make it possible to test relatively refined theories of the shock's structure and particle acceleration properties.  A number of ideas have been suggested for the physics behind the observed dynamic behavior.  Hester (1998a) suggested that the wisp structures reflect synchrotron cooling at constant pressure of the particles whose outflow momenta randomize at a shock in a magnetized electron-positron plasma to a power-law distribution of particles with slope in energy space flatter than $E^{-2}$, an idea that requires the X-ray spectrum of the plasma in the wisp region to be flatter than $\\varepsilon^{-1.5}$ (X-ray flux in photons/keV-cm$^2$-sec). However, only the particles with energy high enough to emit nebular gamma rays ($\\varepsilon > 10 $MeV) have synchrotron loss times short enough to create a cooling instability whose time scales are comparable to the observed variations in the optical and X-ray, and the gamma ray spectrum of the Crab (presumably arising in a compact region around the pulsar no larger than the X-ray source) has a spectral index steeper than 1.5. Chedia et al. (1997) suggest that a drift instability of a subsonically expanding plasma (no shock wave) can explain the wisp structure -- how the outflow from the pulsar escapes catastrophic adiabatic losses is not addressed, however. Begelman (1999) suggests that the wisps represent the nonlinear phase of a Kelvin-Helmholtz instability between subsonic layers expanding in the rotational equator of the pulsar and flow at higher absolute latitude moving with a different velocity.  All of these ideas follow Kennel and Coroniti in assuming that the pulsar's outflow decelerates in a shock wave of unobservably small thickness, with the wisp dynamics due to phenomena in the post shock plasma (or, in the case of the Chedia et al. model, do without a shock completely).  In contrast, Gallant and Arons (1994, hereafter called GA), building on earlier work on the structure and particle acceleration characteristics of relativistic shock waves in collisionless, magnetized electron-positron (Langdon, Arons, \\& Max1988; Gallant et al. 1992) and electron-positron-ion (Hoshino et al.,  1992) plasmas, attributed the observed wisp structure to the {\\it internal} structure of an electron-positron-ion shock wave. The idea is motivated by the fact that the gyration of the heavy ions within the shock transfers energy to positrons and electrons through emission of ion cyclotron waves (formally, the magnetosonic mode of the magnetized electron-positron plasma). This leads to the formation of nonthermal particle spectra in rough agreement with the particle spectra required to explain the optical, X-ray, and gamma-ray emission, if the ions carry a large fraction of the pulsar's total energy loss, at least in the equatorial sector illustrated in Figure \\ref{figwisps}. The deflection of the ions' outflow by the magnetic field deposits a large amount of compressional momentum in the magnetic field, which is transverse to the flow, and into the $e^\\pm$ plasma frozen to the field.  These compressions appear as brightenings in the surface brightness of the nebula.  GA showed that their steady flow model gives a good account of a single high-resolution {\\it I}-band image of the wisps (van den Bergh \\& Pritchett 1989). However, the wisps are known to be highly variable, on a time-scale of weeks to months.  An adequate account of their behavior requires a time-dependent theory. \\begin{figure*}[t] \\centering \\plotone{f1.eps} \\caption{(a) HST optical image of the inner Crab Nebula (after Hester 1995), (b) Zoom in on the wisp region. (c) X-ray nebula from {\\it Chandra} (after Weisskopf 2000; [b] and [c] are nearly to scale); (d) Illustration of the polar jet and the equatorial flow geometry of the Crab wind. {\\it This figure is also available in color in the electronic edition of the {\\it Astrophysical Journal}}}. \\label{figwisps} \\end{figure*}  In this paper we present a model for time-variability of the wisp region in the Crab by extending the work of GA to include the time dependence in the flow. As was realized early on by Hoshino et al. (1992) and Gallant et al. (1992), the coherent ion orbits invoked in the model of GA are unstable to gyrophase bunching and emission of magnetosonic waves. We find that this instability introduces a time dependence into the shock in the pair-ion plasma that causes the variable surface brightness to have substantial similarity to the observed variability of the wisps. Because of the presence of the continuous stream of fresh ions passing through the pair shock, the relativistic ion cyclotron instability provides a mechanism for sustained periodicity and wave emission dynamics that closely reproduces time-resolved observations of the wisps.  In \\S \\ref{sec2} we describe the model and the underlying assumptions. In \\S\\S \\ref{shockdyn} and \\ref{applic}, we present the results of the simulations and the dynamical predictions for the wisp region and compare them to observations.  In \\S \\ref{disc}, we discuss the model's successes, with emphasis on its applicability to other PWNs that have wind-nebula dynamics not very affected by radiation losses, as well as the model's limitations, pointing out further improvements needed in order to make fully quantitative comparisons with observations. We also discuss some observational tests of the ideas at their current level of development in this section. In \\S \\ref{origin} we describe a scenario for the origin of the high-energy ions and their acceleration in the equatorial region of the pulsar wind. Our conclusions are summarized in \\S \\ref{conclu}.  The technical details of our hybrid numerical approach are left for the Appendix. ",
        "conclusions": "\\label{conclu} We have studied the internal structure and time variability of a collisionless shock in an electron-positron pair plasma with an energetically significant admixture of ions. Upon crossing the shock the ion component undergoes relativistic cyclotron instability that renders the shock structure very dynamic. The first loop of the ${\\bf E}\\times {\\bf B}$ ion orbit achieves a limit-cycle behavior and acts as a quasi-periodic wave emitter with a period of half a local Larmor time of the ions. This time-dependence led us to apply our results to the collisionless termination shock in the inner Crab Nebula where we identify the wisps with turning points in the drift orbit of the ions. The limit cycle of the instability produces large amplitude magnetosonic waves on a time-scale of several months that proceed to move downstream from the shock. This time-scale is intrinsic to the model and associates the appearance of the wisps to the mechanism of variability within the shock structure.  This is in contrast to other models that generally do not explain the origin of the wisp perturbations, only their subsequent evolution. Our model naturally makes the region between wisp 1 and wisp 2 in the optical (or the inner {\\it Chandra} X-ray ring and the torus) behave like a near zone of a radiating antenna and prescribes the pattern of brightness fluctuations to wisps 1 and 2 and the moving wisps. This pattern has both a stationary component oscillating in brightness (inner {\\it Chandra} ring and the torus region) and moving wisps that cross the region. That is why the snapshots of the inner nebula taken with insufficient temporal sampling register motion yet on average always see two major wisps that do not seem to leave.  The morphological agreement between the behavior of the wisp region in the Crab and the results of our simulations is an important argument in favor of the presence of ions in the pulsar outflows, which are often presumed to consist only of electron-positron pairs.  Such ions are a natural candidate for the return current flow of the pulsar that starts at the auroral boundary of the polar cap and is then mapped into the equatorial current sheet. In order to get a fit to the nebular dynamics we require an ion injection rate of $10^{34.5}$ $\\rm{s}^{-1}$.  The associated current is of the same order of magnitude as the Goldreich-Julian current of the Crab pulsar. Thus, by studying the interaction of the pulsar wind with the nebula we are able for the first time to get a window on the enigmatic electrodynamics of pulsars.  Electrostatically, ions are the preferred return current carrier for pulsars with ${\\bf \\Omega} \\cdot {\\boldsymbol \\mu} > 0$, which is thought to be the geometry applicable to the Crab pulsar (Yadigaroglu \\& Romani 1995). This means that for a rotator with antialigned rotation and magnetic axes, the equatorial current would be carried by electrons, and their effect on the nebula might be different from the Crab case.  The model of the time variability of collisionless shocks presented here is fairly general and can be applied to other PWNs, as well as other astrophysical sources where collisionless shocks arise.  For the pulsar B1509-58 and the associated nebula Gaensler et al. (2002) estimate a variability period of on the order of 3-5 yr. For Vela the spacing of semicircular X-ray arcs implies a period of variability similar to the Crab, yet there are no similar wisps in that remnant. This could be explained if Vela rings are not features in the equatorial flow, but rather in the polar direction (similar to the Crab halo (Hester et al., 1995)). The differences might also be a consequence of the later evolutionary state of the Vela PWN, where the environment has a chance to influence the termination of the pulsar's wind (Chevalier 1998, 2004). This underscores how much we are dependent on a reliable deconvolution of the geometry and the evolutionary history in all attempts to characterize the physics of non-thermal nebula energization by central compact objects. For sources other than the PWNs, collisionless shocks are common in relativistic AGN jets and thought to occur in gamma-ray bursts. Both settings are known to be quite variable, and whether ion-cyclotron instability can operate in these shocks remains to be studied.  We acknowledge assistance from NASA grants NAG5-12031 and HST-AR-09548.01-A. J.A. thanks the Miller Institute for Basic Research in Science and the taxpayers of California for their support. A.S. acknowledges support provided by NASA through Chandra Fellowship grant PF2-30025 awarded by the Chandra X-Ray Center, which is operated by the Smithsonian Astrophysical Observatory for NASA under contract NAS8-39073.   \\begin{appendix}"
    },
    "0402/astro-ph0402409_arXiv.txt": {
        "abstract": "We demonstrate that the observed increase of some nebular line ratios with height above the midplane in the diffuse ionized gas (DIG) in the Milky Way and other galaxies is a natural consequence of the progressive hardening of the radiation field far from the midplane ionizing sources. To obtain increasing temperatures and line ratios away from the midplane, our photoionization simulations of a multi-component interstellar medium do not require as much additional heating (over and above that from photoionization) as previous studies that employed one dimensional, spherically averaged models.  Radiation leaking into the DIG from density bounded H~{\\sc ii} regions is generally harder in the H-ionizing continuum and has its He-ionizing photons suppressed compared to the ionizing source of the H~{\\sc ii} region.  In line with other recent investigations, we find that such leaky H~{\\sc ii} region models can provide elevated temperatures and line ratios, and a lower He$^+$ fraction in the DIG. For a composite model representing the relative spectral types of O stars in the solar neighbourhood, we find that additional heating less than $10^{-26} n_{\\rm e}$~ergs/s/cm$^3$ can reproduce the observed elevated line ratios in the DIG.  This additional heating is considerably less than previous estimates due to the natural hardening of the radiation field reaching large heights in our simulations. ",
        "introduction": "The presence of extended layers of diffuse ionized gas (DIG) in the Milky Way and other galaxies is inferred from faint emission in H$\\alpha$ and other nebular lines (e.g., Reynolds 1995; Rand 1998; Hoopes, Walterbos, \\& Greenawalt 1996; Domgoergen \\& Dettmar 1997; Otte \\& Dettmar 1999; Hoopes \\& Walterbos 2003; Wang, Heckman, \\& Lehnert 1998; Otte et al. 2001, 2002). The ionization source for the DIG is believed to be O and B stars close to the midplane of galaxies.  Photoionization models using OB stars can reproduce the gross features of the DIG ionization structure (e.g., Miller \\& Cox 1993; Dove \\& Shull 1994).  If extra heating is included, in addition to that from photoionization by O stars, 1D models can reproduce the DIG spectrum (Mathis 1986; Domgoergen \\& Mathis 1994; Mathis 2000; Sembach et al. 2000). The main problem with O stars as the ionization source of the DIG is that for a smooth density distribution the mean free path of Lyman continuum photons is very small, $\\sim 0.1$~pc for $n({\\rm H}^0)=1\\, {\\rm cm}^{-3}$. Thus, it is difficult for photons to traverse the large distances to ionize the extraplanar gas which has a scaleheight of around 1~kpc in the Milky Way (Haffner, Reynolds, \\& Tufte 1999).  However, three dimensional photoionization models show that a two component (Wood \\& Loeb 2000) or fractal ISM (Ciardi, Bianchi, \\& Ferrara 2002) can provide low density paths allowing Lyman continuum photons to reach large $|z|$ heights above their midplane sources.  Measurements of emission lines in addition to H$\\alpha$ provide probes of the abundances, temperatures, and densities in the DIG.  Some of the most studied lines in the Milky Way and other galaxies are [N~II] $\\lambda 6584$, [S~II] $\\lambda 6717$, [O~III] $\\lambda 5007$, [O~II] $\\lambda 3727$ doublet, and [O~I] $\\lambda 6300$. Observations of the He~I $\\lambda 5876$ line provide information on the spectrum of the ionizing sources (Reynolds \\& Tufte 1995).  Hereafter the above lines will simply be referred to as [N~II], [S~II], [O~III], [O~II], [O~I], and He~I. Some differences in the observed line ratios in the DIG compared to traditional H~{\\sc ii} regions are: [S~II]/H$\\alpha$ and [N~II]/H$\\alpha$ ratios increase with height above the plane; [S~II]/[N~II] is quite uniform with height (Haffner et al. 1999) and with latitude (Rand 1997); He is observed to be underionized with respect to H (Reynolds \\& Tufte 1995; Heiles et al. 1996).  Most models for the DIG employ smooth, one dimensional density distributions, and predict volume averages of the line strengths and ratios (e.g., Mathis 1986; Domgoergen \\& Mathis 1994; Collins \\& Rand 2001; Sembach et al. 2000; Mathis 2000).  These models generally fail to reproduce the observed line ratios that increase with $|z|$ above the midplane.  Haffner et al. (1999) and Reynolds, Haffner, \\& Tufte (1999) showed that the observed [S~II]/H$\\alpha$ and [N~II]/H$\\alpha$ line ratios may be explained if the gas temperature increases with height above the midplane.  Including additional heating over and above that of pure photoionization can reconcile the one dimensional models with observations (Reynolds et al. 1999; Mathis 2000).  Additional heating may plausibly arise from photoelectric heating from dust (Reynolds \\& Cox 1992), dissipation of turbulence (Slavin et al. 1993; Minter \\& Spangler 1997), and shocks (Raymond 1992). However, it is well known that in photoionized H~{\\sc ii} regions the radiation field hardens towards the edge of the Str{\\\"o}mgren sphere resulting in the highest temperatures occurring at the largest distances from the ionizing source.  This arises because low energy photons have relatively short mean free paths and are absorbed close to the source.  Higher energy photons have longer mean free paths, travel farther, and deposit more energy per photon at large distances from the source, giving rise to the increasing temperature away from the source (Osterbrock 1989).  Such a scenario as described above is almost certainly occurring in the DIG, with the ionizing spectra penetrating to large $|z|$ above the plane being significantly harder than the source spectra, naturally producing a temperature profile that increases with $|z|$.  Models presented by Bland-Hawthorn, Freeman, \\& Quinn (1997), Wang et al. (1998), and Rand (1998) take this radiation transfer effect into account by introducing a hardening of the radiation field with $|z|$. In these plane parallel models, the hardening of the radiation field leads to increasing temperatures at large distances from the illuminated face and corresponding increases in [S~II]/H$\\alpha$ and [N~II]/H$\\alpha$. Mathis (2000) also mentions this effect, stating that line ratios for lines of sight that pierce the outer edges of spherical models may reproduce the observations.  Radiation leaking into the ISM from traditional H~{\\sc ii} regions may also be hardened compared to the source spectrum.  Hoppes \\& Walterbos (2003) investigated photoionization by photons from leaky H~{\\sc ii} regions finding that the hardened spectrum leads to elevated temperatures and increased line ratios when compared to models that do not envoke hardening of the source spectrum.  In addition, leaky H~{\\sc ii} region models lead to a suppression of He-ionizing photons ($h\\nu > 24.6$~eV) and a corresponding decrease in ionization stages such as He$^+$ and N$^{2+}$.  In this paper we present Monte Carlo photoionization models for a multi-component ISM.   In what follows we describe the photoionization code, adopted ISM density structure, and spectra for the ionizing sources.  Due to the very large parameter space, we restrict this paper to two dimensional models of a single source ionizing a multi component, stratified ISM.  Note that although our models are for 2D systems, our simulations are run on 3D density grid. A future paper will present models for 3D geometries and illuminations, investigating the role of diffuse ionizing radiation, 3D ionization and temperature structures, and the resulting intensity maps. Some specific issues we address in this paper are: increased temperatures with increasing $|z|$ above the midplane; problems arising with fitting [S~II] emission due to undetermined dielectronic recombination rates; predictions of excess [O~I] emission compared with observations; the ionization of He within the DIG. ",
        "conclusions": "\\subsection{Sulfur and Nitrogen}  The model line ratios presented above appear to be in good agreement with observations of [S~II] and [N~II] in the local DIG (Haffner et al. 1999), NGC~891 (Rand 1998), and several other galaxies (e.g., Otte et al. 2001, 2002). The increase of [N~II]/H$\\alpha$ and [S~II]/H$\\alpha$ with $|z|$ is a natural consequence of increasing temperatures away from the ionizing source due to the hardening of the radiation field. Compared to one dimensional averaged models, our two dimensional simulations reduce the requirement for extra heating to explain the increasing temperatures and line ratios with height above the midplane.  The change of slope in the [S~II]/H$\\alpha$ vs [N~II]/H$\\alpha$ scatter plots (e.g., Fig.~8) are not observed in the Milky Way's DIG (Haffner et al. 1999).  The change of slope in our simulations is due to the increased S$^+$/S and N$^0$/N fractions at the edge of the ionized volumes. The fact that such slope changes are not observed suggests that the DIG is almost fully ionized and not density bounded like our single source models.  We will investigate multiple source models with overlapping ionized regions in a separate paper. Alternatively, it is quite likely that our models do not provide a good representation of the emission at the ionized/neutral interface as we do not include the effects of shocks or ionization fronts.  The role of interfaces in interpreting ISM observations is very important (Reynolds 2004), and the very large [S~II]/H$\\alpha$ and [N~II]/H$\\alpha$ line ratios in our simulations may be a result of incomplete physics in our simulations.  Some galaxies do show changes in slope of the [S~II]/H$\\alpha$ vs [N~II]/H$\\alpha$ and this has recently been interpreted by Elwert, Dettmar, \\& T{\\\"u}llmann (2003) as an indicator for chemical evolution in galaxies.  They suggest that increased [S~II]/H$\\alpha$ compared to [N~II]/H$\\alpha$ may arise from younger DIG layers since in the ISM nitrogen (from low mass stars, planetary nebulae, and stellar winds) is enriched more slowly than sulfur (from spuernovae type~II).  Our models adopt uniform abundances throughout the simulation grid and do not address this scenario.  \\subsection{Oxygen}  There is currently limited data on oxygen lines in the Milky Way's DIG, with most observations probing the DIG at $b=0^\\circ$.  However, in NGC~891 there are detailed observations of the dependence of [O~III]/H$\\alpha$ and [O~I]/H$\\alpha$ with height above the plane and Otte et al. (2001, 2002) have made measurements of [O~II]/H$\\alpha$ in several galaxies. Rand (1998) finds that [O~III]/H$\\alpha$ increases with height, which is opposite to what is seen in almost all of our simulations.  The increase of [O~III]/H$\\alpha$ with $|z|$ in NGC~891 is cited as evidence for O being ionized by a different mechanism, such as shocks (e.g., Collins \\& Rand 2001), instead of pure photoionization.  Our models do not consider dynamics or ionization fronts, so cannot address these effects. Note, however, that models with very hard spectra ($T_\\star = 50000$~K, Fig.~7) can produce increasing [O~III]/H$\\alpha$ at large $|z|$.  Otte et al. (2001, 2002) observe [O~II]/H$\\alpha$ and [O~III]/H$\\alpha$ to increase with height above the plane in five galaxies they studied, finding $0.5\\la {\\rm [O~II]/H}\\alpha\\la 5$. They also observed increases of [S~II]/H$\\alpha$ and [N~II]/H$\\alpha$ with $|z|$.  Our pure photoionization models predict $0.5\\la {\\rm [O~II]/H}\\alpha\\la 3$, and additional heating can raise this even further (e.g., Fig.~9).   It appears that multi-dimensional pure photoionization models can reproduce most of the [O~II]/H$\\alpha$ observations of Otte et al. (2001, 2002), though additional heating or a harder ionizing spectrum may be required for some of the largest [O~II]/H$\\alpha$ ratios.  The [O~I]/H$\\alpha$ line ratios in our models show that [O~I] increases in strength towards the edge of the ionized zone.  The strong charge-exchange coupling of O$^0$ to H$^0$ and the increased temperatures towards the edge of the ionized zone result in the increased [O~I]/H$\\alpha$.  This is generally seen as a problem with photoionization models, since in the few observations in the local DIG, albeit at $b=0^\\circ$, [O~I] is observed to be rather weak with [O~I]/H$\\alpha \\la 0.03$.  The [O~I] emission may be reduced if the region is fully ionized, or density bounded instead of radiation bounded (e.g., Mathis 2000; Sembach et al. 2000).  This is seen in the vertical cuts showing that [O~I]/H$\\alpha < 0.03$ in models where the gas is ionized beyond the maximum $|z|$ of our simulation box (e.g., Fig.~3). The increasing [O~I]/H$\\alpha$ with $|z|$ in some of our simulations (e.g., Fig.~6) do match the data for NGC~891, where [O~I]/H$\\alpha\\sim 0.1$ at large $|z|$ (Rand 1998).  More observations of [O~I] at larger $|z|$ in the Galactic DIG will determine whether there really is a difference in the [O~I] emission between the Milky Way and NGC~891.   \\subsection{Helium}  As with [O~I], the few measurements of He~I in the Galactic DIG are at $b=0^\\circ$.  The observations indicate that helium is underionized relative to hydrogen with ${\\rm He}^+/{\\rm H}^+ \\la 0.04$ (Reynolds \\& Tufte 1995; see also Heiles et al. 1996). The He~I observations probe the ionizing spectrum for the DIG and indicate a relatively soft spectrum, typically spectral type O8 or later, corresponding to $T_\\star \\sim 36000$~K. The situation is not as extreme in NGC~891 where He$^+$/H$^+ \\sim 0.06$, indicating a harder ionizing spectrum for its DIG (Rand 1997).  Our simulations using a composite spectrum, representative of the O stars in the solar neighbourhood, produce fairly low values for He~I/H$\\alpha$, in line with current observations at $b=0^\\circ$.  Further observations of He~I at higher latitudes will provide additional tests of our models and the hardening of the radiation field at large $|z|$.  We have also investigated an alternative mechanism for reducing the helium ionization, even when the sources are hotter than O8. The basic idea is that the radiation leaking out of midplane H~{\\sc ii} regions to ionize the DIG may have its helium ionizing photons suppressed (e.g., Hoopes \\& Walterbos 2003). Our leaky H~{\\sc ii} region models do indeed show that low values for He$^+$/H$^+$ may be obtained, even if the ionizing source is quite hot. For the same ionizing luminosity, these models also produce higher temperatures than our other models, since the leaking ionizing spectrum is harder in the H-ionizing continuum. Therefore, leaky H~{\\sc ii} regions provide a plausible mechanism for explaining the low helium ionization in the Galactic DIG."
    },
    "0402/astro-ph0402315_arXiv.txt": {
        "abstract": "\\noindent Neutrino emission drives neutron star cooling for the first several hundreds of years after its birth. Given the low energy ($\\sim$ keV) nature of this process, one expects very few nonstandard particle physics contributions which could affect this rate. Requiring that any new physics contributions involve light degrees of freedom, one of the likely candidates which can affect the cooling process would be a nonzero magnetic moment for the neutrino. To illustrate, we compute the emission rate for neutrino pair bremsstrahlung in neutron-neutron scattering through photon-neutrino magnetic moment coupling. We also present analogous differential rates for neutrino scattering off nucleons and electrons that determine neutrino opacities in supernovae. Employing current upper bounds from collider experiments on the tau magnetic moment, we find that the neutrino emission rate can exceed the rate through neutral current electroweak interaction by a factor two, signalling the importance of new particle physics input to a standard calculation of relevance to neutron star cooling. However, astrophysical bounds on the neutrino magnetic moment imply smaller effects. ",
        "introduction": "Neutrino physics plays a crucial role in the birth and subsequent evolution of neutron stars, beginning with the hot and dense environment of a supernova~\\cite{burrows}, where diffusing neutrinos are believed to trigger the explosive event, to the interior of a cold neutron star, whose cooling rate is determined principally by free-streaming neutrinos~\\cite{prakash}.  While various problems remain with the neutrino-driven supernova explosion mechanism despite recent intensive efforts (see~\\cite{buras} and refs. therein), uncertainties in neutrino emission from neutron stars is limited to a lack of knowledge of the precise underlying equation of state of supra-nuclear matter, where progress is tied to improving many-body calculations~\\cite{schwenk}. The important point is that the long-term cooling of neutron stars is controlled by neutrino emission, and this stage lasts up to about $10^5$ years of age, when cooling by emission of photons becomes more effective.  \\vskip 0.2cm  There is a host of well-known neutrino emission processes that operate in the crust and core of the neutron star. A comprehensive list of neutrino emitting reactions relevant to different regions of the star can be found in~\\cite{Yakovlev}. Which one dominates the cooling depends mainly on the temperature and to a lesser extent, on the density. It should be noted that neutron or proton superfluidity can also reduce or enhance neutrino emission dramatically through density-dependent gaps~\\cite{gnedin}. Since cooling rates are dependent on neutrino emissivities, it has proven useful to turn to neutron star cooling to help identify or constrain new physics contributions to these emissivities~\\cite{Yak03}. A typical example is the upper bound on the axion mass (or coupling) from axion bremsstrahlung emission in neutron-neutron collisions~\\cite{umeda} and stellar cooling~\\cite{georg}.  \\vskip 0.2cm  In this work, we are interested in a particular neutrino emission process involving the neutrino magnetic dipole moment. A magnetic moment for the neutrino was postulated almost immediately following Pauli's neutrino hypothesis. The astrophysical consequences of a neutrino magnetic moment, particularly on stellar cooling have been investigated previously (see~\\cite{Iwamoto95} and refs. 3-9 therein). These works have focused their attention on the plasmon decay process, which is the dominant cooling mechanism in red giants, and also for the crust of a neutron star until thermal conduction lowers temperatures significantly to the point when processes occurring in the core become more efficient. Since the crust is only a small volume fraction of the star, specially for softer equations of state, it is important to address neutrino emission from the core which could also receive important contributions from new physics. As an illustrative example of this, we present a calculation of the neutrino emissivity from neutrino bremsstrahlung in neutron-neutron collisions, where a neutron radiates an off-shell photon (via magnetic moment coupling of the neutron), which subsequently decays to a neutrino-antineutrino pair via magnetic moment coupling. Section IIA outlines the general features and expected significance of particle-physics corrections to known neutrino rates. In section IIB, we explain why neutrino bremsstrahlung from neutron-neutron collisions can be the dominant cooling process in the core, and present the computation of the total emissivity from magnetic moment interactions, which can be appreciable. We briefly discuss corrections to scattering rates of neutrinos by electrons and nucleons which dominantly determine neutrino opacities and mean free paths in supernovae environments. The conditions under which new particle physics corrections such as presented here can be quantitatively large are explored in section III. We conclude with a discussion of the relevance of our results in section IV.  \\vskip 0.2cm ",
        "conclusions": "We have explored the relevance of particle physics beyond the standard model to standard neutrino pair emission rates from neutron star interiors, focusing especially on the role of the neutrino magnetic moment. General arguments based on relevant energy scales imply that a neutrino magnetic moment provides the most interesting and feasible correction to known physics. In the low-energy approximate treatment of neutrino pair bremsstrahlung from neutron matter, we have estimated the correction to a benchmark calculation of the neutrino pair emissivity from electroweak interactions in neutron matter. The magnitude of the correction is estimated from limits on the neutrino magnetic moment $\\kappa$. \\\\  While stringent collider bounds on the $\\nu_e$ and $\\nu_{\\mu}$ electromagnetic moments exist, the $\\nu_{\\tau}$ is not as nearly strongly constrained, and employing a range set by collider and astrophysical inputs, we find that the standard emission rates may be revised by upto a factor of two or more. We expect similar results for neutrino scattering via magnetic moment interactions, leading to modified neutrino opacities. We have estimated the magnitude of such corrections by using elastic rates and free scatterers. For practical calculations in supernovae physics, we need to explore similar corrections to inelastic rates, which requires including medium effects and energy transfer. These studies are currently in progress. The magnetic moment interaction vertex remaining unaffected in medium, we would expect comparable enhancements upon inclusion of medium effects, therefore we have, in this work, presented ratios of rates rather than their absolute values. We observe that large corrections to the differential cross-section are obtained for neutrino-proton scattering over a wide range of angles and neutrino-electron scattering in the forward direction. The correction to neutrino-neutron scattering is expected to be at the few percent level.  \\vskip 0.2cm  It is also pertinent to mention here the consequences of these revised emissivities for the temperature versus time profile (cooling curve) of a typical neutron star. While stellar cooling is determined (for late times) principally by neutrino emission from the core, a temperature profile requires mapping the surface temperature to the core temperature, a procedure that is strongly dependent on important factors such as local surface temperature variations, magnetospheric emission as well as the composition of the stellar envelope and atmosphere. These details form an integral part of the surface temperature determination. Furthermore, knowledge of the temperature gradient within the core and upto the star's surface requires as input, the equation of state of nuclear matter at supra-nuclear density, and knowledge of related variables such as the incompressibility and the symmetry energy, which are poorly constrained in neutron matter. Model-dependent theoretical uncertainties tend to mask the effects of improved or revised estimates of standard emissivities when incorporated into a numerical simulation of neutron star cooling, and observational difficulties complicate the direct comparison of such numerical results to data. This uncertainty may be removed if a novel temperature or density dependence is discovered as in the case of the direct urca process or if exotic phases such as pion or kaon condensates exist from which cooling can be more efficient. It is unlikely that particle physics corrections to emissivities that are of ${\\cal O}(1)$ would considerably alter the overall cooling scenario. Nevertheless, the interesting possibility remains that large contributions from non-standard model physics may exist and yet remain hidden in the microphysics of neutron star cooling."
    },
    "0402/astro-ph0402664_arXiv.txt": {
        "abstract": "{ Among the hundred or so extrasolar planets discovered to date, 19 are orbiting a component of a double or multiple star system. In this paper, we discuss the properties of these planets and compare them to the characteristics of planets orbiting isolated stars. Although the sample of planets found in multiple star systems is not large, some differences between the orbital parameters and the masses of these planets and the ones of planets orbiting single stars are emerging in the mass--period and in the eccentricity--period diagrams. As pointed out by \\cite{Zucker02b}, the most massive short-period planets are all found in multiple star systems. We show here that the planets orbiting in multiple star systems also tend to have a very low eccentricity when their period is shorter than about 40 days. These observations seem to indicate that some kind of migration has been at work in the history of these systems. The properties of the five short-period planets orbiting in multiple star systems seem, however, difficult to explain with the current models of planet formation and evolution, at least if we want to invoke a single mechanism to account for all the characteristics of these planets. ",
        "introduction": "Studies of stellar multiplicity among solar-type stars of the solar neighbourhood have shown that about 40\\% of the G and K dwarfs can be considered to be real single stars \\citep{Duquennoy91,Eggenberger03c}. As the majority of solar-type stars belong to double or multiple star systems\\footnote{In this paper, double and multiple star systems will be called multiple star systems.}, it is of interest to consider the existence of planets in such an environment. Searches for extrasolar planets using the radial velocity technique have shown that giant planets exist in certain types of multiple star systems (see Table 1 for further details). The number of such planets is, however, still low, in part because close binaries are difficult targets for radial velocity surveys and were consequently often rejected from the samples. Due to the limitations of the available observational techniques, most detected objects are giant (Jupiter-like) planets; the existence of smaller mass planets in multiple star systems is thus still an open question.  \\begin{table*}  \\caption{Planets orbiting a component of a multiple star system with confirmed orbital or common proper motion (CPM stands for common proper motion and SB for spectroscopic binary).} \\begin{center} \\begin{tabular}{lcccccc} \\hline Star & a$_{{\\rm b}}$ & a$_{{\\rm p}}$ & M$_{{\\rm p}}\\sin{i}$ & e$_{{\\rm p}}$ & Notes & References\\\\ &  (AU)           &           (AU) &  (M$_{\\rm J}$)       &                &       &           \\\\ \\hline \\object{HD\\,40979}              & $\\sim$6400        & 0.811    & 3.32   & 0.23  & CPM$^{a}$& 12,11  \\\\ \\object{Gl\\,777\\,A}             & $\\sim$3000        & 4.8     & 1.33   & 0.48   & CPM           & 1,23 \\\\ \\object{HD\\,80606}              & $\\sim$1200        & 0.469   & 3.90   & 0.927 & CPM            & 22  \\\\ \\object{55\\,Cnc}                & $\\sim$1065        & 0.115   & 0.84   & 0.02  & CPM            & 8,21,19,2  \\\\ &                   & 0.24   & 0.21   & 0.34  &                &          \\\\ &                   & 5.9     & 4.05   & 0.16  &                &          \\\\ \\object{16\\,Cyg\\,B}             & $\\sim$850         & 1.6    & 1.5   & 0.634  & CPM            & 24,5,15  \\\\ \\object{Ups\\,And}               & $\\sim$750         & 0.059   & 0.71   & 0.034 & CPM            & 17,24,2,3\\\\ &                   & 0.83   & 2.11   & 0.18  &                &  \\\\ &                   & 2.50    & 4.61  & 0.44  &                &  \\\\ \\object{HD\\,178911\\,B}          & $\\sim$640         & 0.32    & 6.292  & 0.1243& CPM            & 28,30  \\\\ \\object{HD\\,219542\\,B}          & $\\sim$288         & 0.46    & 0.30   & 0.32  & CPM            & 7 \\\\ \\object{Tau\\,Boo}               & $\\sim$240         & 0.05    & 4.08   & 0.018  & orbit          & 13,24,2\\\\ \\object{HD\\,195019}             & $\\sim$150         & 0.14    & 3.51   & 0.03   & CPM            & 24,1,10 \\\\ \\object{HD\\,114762}             & $\\sim$130         & 0.35    & 11.03  & 0.34  & CPM            & 24,16,18  \\\\ \\object{HD\\,19994}              & $\\sim$100         & 1.54     & 1.78   & 0.33   & orbit          & 13,25,20 \\\\ \\object{HD\\,41004\\,A}           & $\\sim$23          & 1.33    & 2.5    & 0.39  & SB             & 29,31,27 \\\\ \\object{$\\gamma$\\,Cep}          & $\\sim$22          & 2.03    & 1.59   & 0.2   & SB             & 4,6,14 \\\\ \\object{Gl\\,86}                 & $\\sim$20          & 0.11    & 4.0    & 0.046 & CPM, SB$^{b}$ & 9,26 \\\\ \\hline \\end{tabular} \\end{center} \\label{tab1} {\\small { Notes: (a) According to \\citet{Halbwachs86}, this pair has only a probability of 60\\% to be physical. The physical nature of this binary has however been confirmed later on the basis of CORAVEL radial velocity measurements (Halbwachs, private communication); } (b) The multiplicity status of this system has still to be clarified.  \\vspace{1.5mm}  References: (1) \\citealt{Allen00}; (2) \\citealt{Butler97}; (3) \\citealt{Butler99}; (4) \\citealt{Campbell88}; (5) \\citealt{Cochran97}; (6) \\citealt{Cochran02}; (7) \\citealt{Desidera03}; (8) \\citealt{Duquennoy91}; (9) \\citealt{Els01}; (10) \\citealt{Fischer99}; (11) \\citealt{Fischer03}; (12) \\citealt{Halbwachs86}; (13) \\citealt{Hale94}; (14) \\citealt{Hatzes03}; (15) \\citealt{Hauser99}; (16) \\citealt{Latham89}; (17) \\citealt{Lowrance02}; (18) \\citealt{Marcy99}; (19) \\citealt{Marcy02}; (20) \\citealt{Mayor03}; (21) \\citealt{McGrath02}; (22) \\citealt{Naef01}; (23) \\citealt{Naef03}; (24) \\citealt{Patience02}; (25) \\citealt{Queloz00b}; (26) \\citealt{Queloz00}; (27) \\citealt{Santos02}; (28) \\citealt{Tokovinin00}; (29) \\citealt{Udry03spec}; (30) \\citealt{Zucker02}; (31) \\citealt{Zucker03b}.} \\end{table*}   The orbital characteristics and the mass distribution of extrasolar planets can give us an insight into their formation mechanisms and their subsequent evolution. In the first paper of this series, \\citeauthor{Udry03} (\\citeyear{Udry03}; \\citetalias{Udry03}) discussed the period distribution and the mass--period diagram for extrasolar planets orbiting single stars. As pointed out by \\citet{Zucker02b}, planets orbiting a component of a multiple star system seem to have different characteristics than planets orbiting single stars. \\citet{Zucker02b} showed that there is a significant correlation in the mass--period diagram for planets orbiting single stars, while there may be an anticorrelation in this same diagram for planets found in multiple star systems. The difference is mainly due to a paucity of massive planets with short periods, and to the fact that the most massive short-period planets are all found in binaries.  The characteristics of extrasolar giant planets have forced considerable modifications of the standard model of planet formation. It is now usually believed that planets form within a protoplanetary disc of gas and dust orbiting a central star, but the precise modes by which this formation takes place are still debated, especially for giant planets \\citep[e.g.][]{Pollack96,Boss97,Bodenheimer00,Boss00b,Wuchterl00,Boss03}. Two major models have been proposed to explain giant planet formation (see Sect.~\\ref{giantformation}), each with its advantages and limitations, but there is currently no consistent model that accounts for all the observed characteristics of extrasolar planets. Observational constraints are thus needed, not only to specify our understanding of planet formation and evolution, but also to possibly discriminate between the proposed models. In this context, the detection and the characterization of planets orbiting in multiple star systems, even if more difficult to carry out than the study of planets orbiting isolated stars, may bring new constraints and additional information.  This paper is organized as follows. The sample of planets found in multiple star systems is presented in Sect.~\\ref{obs}. Some trends seen in the statistics are then emphasized in Sect.~\\ref{stat}. Models of formation and evolution of giant planets in binaries are briefly reviewed in Sect.~\\ref{models} and their predictions are compared to the observations in Sect.~\\ref{discussion}. Our conclusions are drawn in Sect.~\\ref{conclusion}. ",
        "conclusions": "\\label{conclusion}  The characteristics of giant planets found in multiple star systems seem to be different from the ones of planets orbiting single stars, at least for the short-period planets. The major differences are: \\begin{itemize} \\item the most massive ($M_2\\sin{i}\\gtrsim 2$\\,M$_{\\mathrm J}$) short-period planets all orbit in multiple star systems; \\item the planets found in multiple star systems tend to have a very low eccentricity when their period is shorter than 40 days. \\end{itemize} These observations seem to indicate that migration has played an important role in the history of the short-period planets orbiting in multiple star systems and that migration may be induced differently in binaries than around single stars.  From the theoretical point of view, it has been shown \\citep{Kley00} that the presence of a companion star affects the properties of a Jupiter-mass planet still embedded in a disc around the primary component of a binary by increasing the migration and mass accretion rates. Alternatively, the Kozai mechanism may be (or have been) at work in binaries hosting planets and will modify some of the orbital parameters of the planet. This mechanism can be efficient in wide binaries and, coupled with tidal dissipation, it may also lead to inward migration. Even if these two mechanisms may be invoked to explain the characteristics of a few planets orbiting in binaries, none of them seem to be able to account for all the properties of the five planets orbiting in multiple stars systems with a period $P\\lesssim 40$ days. Nonetheless, it is also possible that diverse mechanisms may have been at work in these systems, but leading to a similar final state and similar planet properties. New studies dedicated to this issue will be needed to settle this question and to find a satisfactory explanation to the existence and the characteristics of the short-period planets found in multiple star systems.  From the observational point of view, a larger sample of planets orbiting in multiple star systems will be required to confirm or refute the preliminary trends emphasized in this paper. In this context, the search for planets in multiple star systems, even if more difficult to carry out than the search for planets around single stars is of importance. On the other hand, the characterization of the star systems susceptible of hosting planets is underway and could bring interesting constraints for the models, thus helping our understanding of giant planet formation.     \\appendix"
    },
    "0402/astro-ph0402387_arXiv.txt": {
        "abstract": "\\noindent We use the most recent Type-Ia Supernova data in order to study the dark energy - dark matter unification approach in the context of the Generalized Chaplygin Gas (GCG) model. Rather surprisingly, we find that data allow models with $\\alpha > 1$. We have studied how the GCG adjusts flat and non-flat models, and our results show that GCG is consistent with flat case up to $68\\%$ confidence level. Actually this holds even if one relaxes the flat prior assumption. We have also analyzed what one should expect from a future experiment such as SNAP. We find that there is a degeneracy between the GCG model and a XCDM model with a phantom-like dark energy component. ",
        "introduction": "Recent cosmological observations reveal that the Universe is dominated by two invisible components. Type-Ia Supernova observations \\cite{Riess1998,Garnavich1998,Perlmutter1999},  nucleosynthesis constraints \\cite{Burles2001}, Cosmic Microwave Background Radiation (CMBR) power spectrum \\cite{Balbi2000,deBernardis2000,Jaffe2001}, large scale structure \\cite{Peacock2001} and, determinations of the matter density \\cite{Bahcall1998,Carlberg1998,Turner2000} allow for a model where the clumpy component that traces matter, dark matter, amounts for about $23\\%$ of the cosmic energy budget, while an overall smoothly distributed component, dark energy, amounts for approximately $73\\%$ of the cosmic energy budget.  The most interesting feature of this dark energy component is that it has a negative pressure and drives the current accelerated expansion of the Universe \\cite{Riess1998,Garnavich1998,Perlmutter1999}. From the theoretical side, great effort has been devoted to model dark energy.  The most obvious candidate  is the vacuum energy, an uncanceled cosmological constant [see eg. \\citeN{Bento1999}, \\citeN{Bento2001b}] for which $\\omega_x \\equiv p_x / \\rho_x = -1$. Another possibility is a dynamical vacuum  \\cite{Bronstein1933,Bertolami1986a,Bertolami1986b,Ozer1987} or quintessence. Quintessence models most often involve a single scalar field \\cite{Ratra1988a,Ratra1988b,Wetterich1988,Caldwell1998,Ferreira1998,Zlatev1999,Binetruy1999,Kim1999,Uzan1999,Amendola1999,Albrecht2000,Bertolami2000,Banerjee2001a,Banerjee2001b,SenSen2001,Sen2001} or two coupled fields  \\cite{Fujii2000,Masiero2000,Bento2002a}. In these models, the cosmic coincidence problem, that is, why did the dark  energy start to dominate the cosmological evolution only fairly recently, has no satisfactory solution and some fine tuning is required.  More recently, it has been proposed that the  evidence for a dark energy component might be explained by a change in the equation of state of the background fluid, with an exotic equation of state, the generalized Chaplygin gas (GCG) model, rather than by  a cosmological constant or the dynamics of a scalar field rolling down a potential \\cite{Kamenshchik2001,Bilic2002,Bento2002b}. In this proposal, one considers the evolution of the equation of state of the background fluid instead of a quintessence potential. The striking feature of this model is that it  allows for an unification of dark energy and dark matter \\shortcite{Bento2002b}. Moreover, it is shown that the GCG model may be accommodated within the standard structure formation scenario \\shortcite{Bilic2002,Bento2002b}. Concerns about this point have been raised by \\citeN{Sandvik2002}, however in this analysis, the effect of baryons has not been taken into account, which was shown to be important and allowing compatibility with the 2DF mass power spectrum \\cite{Beca2003}. Also, the \\citeN{Sandvik2002} claim was based on the linear treatment of perturbations close to the present time, thus neglecting any non-linear effects.  Thus, given it potentialities, the GCG model has been the subject of great interest, and various attempts have been made to constrain its parameters using the available observational data. Studies include Supernova data and power spectrum \\cite{Avelino}, age of the Universe and strong lensing statistics \\cite{Dev1}, age of the Universe and Supernova data \\cite{Makler,Alcaniz}. The tightest constraints were obtained by \\citeN{Bento2003a} using the CMBR power spectrum measurements from BOOMERANG \\cite{Boomerang} and Archeops \\cite{Archeops}, together with  the SNe Ia constraints. It is shown that $0.74 \\lsim A_s \\lsim 0.85$, and $\\alpha\\lsim 0.6$, ruling out the pure Chaplygin gas model. From the bound arising from the age of the APM 08279+5255 source, which is $A_s\\gsim 0.81$ \\cite{Alcaniz}, one can get tight constraints, namely  $0.81 \\lsim A_s \\lsim 0.85$, and $0.2 \\lsim \\alpha\\lsim 0.6$, which also rules out the $\\Lambda$CDM model. These results were in agreement with the WMAP data \\cite{Bento2003b}. It was also shown that the gravitational lensing statistics from future large surveys together with SN Ia data from SNAP will be able to place interesting constraints the parameters of GCG model \\cite{Silva2003}. As we shall see in Sections 3 and 4, all these constraints are consistent with Supernova data at $95\\%$ confidence level.  Recently \\citeN{Choudhury2003b} have analyzed the supernova data with currently available 194 data points [see also \\citeN{Choudhury2003a}] and shown that it yields relevant constraints on some cosmological parameters. In particular, it shows that when one considers the full supernova data set, it rules out the decelerating model with significant confidence level. They have also shown that one can measure the current value of the dark energy equation of state with higher accuracy and the data prefers the phantom kind of equation of state, $\\omega_X < -1$ \\cite{Caldwell2002}. Moreover, the most significant  observation of their analysis is that, without a flat prior, the latest Supernova data also rules out the preferred flat $\\Lambda$CDM  model which is consistent with other cosmological observations. In a previous paper, \\citeN{Alam2003} have reconstructed the equation of state of the dark energy component using the same set of Supernova data and found that the dark energy evolves rapidly from  $\\omega_x \\simeq 0$ in the past to a strongly negative equation of state ($\\omega_x \\lsim -1$) in the present, suggesting that $\\Lambda$CDM may not be a good choice for dark energy. More recently, other groups  have also analyzed these recent Supernova data in the context of different cosmological models for dark energy \\cite{Gong2004,Perivol2004}.  In this paper, we analyze the GCG model in the light of the latest supernova data \\cite{Tonry2003,Barris2003}. We consider both  flat and non-flat models. Our analysis shows that the problem with the flat model, which has been discussed in \\citeN{Choudhury2003b}, can be solved in the GCG model in a sense that flat GCG model is consistent with the latest Supernova data even without a flat prior.We have also analyzed the confidence contours for a GCG model, that one expects from a future experiment such as SNAP. We find that there is a degeneracy between the GCG model and a XCDM model with a phantom-like dark energy component.  This paper is organized as follows. In Section 2 we discuss various aspects of the generalized Chaplygin gas model and its theoretical underlying assumptions. In Section 3 we describe our best fit analysis of the most recent supernova data in the context of generalized Chaplygin gas model. Section 4 contains our analysis for expected SNAP results. Finally, in Section 5 we present our conclusions. ",
        "conclusions": "We have analyzed the currently available 194 supernova data points within the framework of the generalized Chaplygin gas model, regarding GCG as a unified candidate for dark matter and dark energy. We have considered both, flat and non-flat cases, and used the best fit value for ${\\cal M} = -0.033$  which is independent of a specific model, throughout our analysis. For the first time, we have crossed the $\\alpha = 1$ limit for the GCG model and try to see whether the data actually allow it or not.  For the flat case, we have studied both cases, restricting $\\alpha$ to the range $0\\leq\\alpha\\leq1$ and also without any restriction on the upper limit $\\alpha$. From Figures 1 and 2, it is quite clear that data favours $\\alpha > 1$, although there is a strong degeneracy in $\\alpha$. Also the quality of fit improves substantially as one relaxes the $\\alpha = 1$ restriction. Moreover, the minimum values allowed for $\\alpha$ and $A_s$ at $68\\%$ confidence level are [0.78, 0.778], which excludes the $\\alpha = 0$, $\\Lambda$CDM case, although there is no constrain on $\\alpha$ at $95\\%$ confidence level.  Moreover, if one does not assume a flat prior for the analysis, our study shows that the flat GCG model for $\\alpha$ values sufficiently different from zero, is consistent with the Supernova data up to $68\\%$ confidence level. It also allows small, both positive and negative curvature, making the GCG a somewhat better description than the $\\Lambda$CDM model. This is consistent with recent result which shows that without a flat prior, a flat $\\Lambda$CDM model, which is otherwise consistent with different cosmological observations, is not a good fit to the supernova data \\cite{Choudhury2003b}. Moreover, the fact that GCG is a better fit to the Supernova data than $\\Lambda$CDM, is consistent with the result of \\citeN{Alam2003}, who also have reconstructed a similar kind of evolving equation of state for the dark energy from the latest Supernova data.  We have also studied the confidence contours for a GCG and XCDM model expected from the future SNAP observation assuming different fiducial universes. In this regard, the degeneracy between the GCG model and a phantom-like dark energy scenario has made obvious in Section 4, where we have shown that when fitting a XCDM model to a GCG universe, the data will favour a phantom energy component, and vice-versa. This degeneracy is also illustrated analytically through the expression for the luminosity distance $d_{L}$ as function of redshift. We have shown that for higher redshifts, GCG model is completely degenerate with a XCDM model with a phantom type of constant equation of state. We mention that it has already been noted in \\citeN{Maor2002} that time varying equations of state might be confused with phantom energy, and here we show that this is true for the GCG, without breaking the dominant energy condition and without a big rip singularity in the future. It should also be noted that with the exception of a cosmological constant, most dark energy models predict a time varying equation of state, therefore a constant dark energy equation of state might not be the best parametrization for dark energy. We have also shown that for a $\\Lambda$CDM fiducial model, confidence regions for a GCG model, which are expected from future SNAP experiment, are quite different from what we have shown in Section 3.1, suggesting that SN Ia data does not favour the $\\Lambda$CDM model.  Thus, our study shows that the generalized Chaplygin gas model is a very good fit to the latest Supernova data both with or without a flat prior. With future data, one expects the error bars to be reduced considerably, but we still expect that Supernova data will favour a generalized Chaplygin gas model with high confidence.  \\vskip 2cm    \\centerline{\\bf"
    },
    "0402/astro-ph0402452_arXiv.txt": {
        "abstract": "We present modeling to explore the conditions of the broad-line emitting gas in two extreme Narrow-line Seyfert 1 galaxies, using the observational results described in the first paper of this series. Photoionization modeling using {\\it Cloudy} was conducted for the broad, blueshifted wind lines and the narrow, symmetric, rest-wavelength-centered disk lines separately.  A broad range of physical conditions were explored for the wind component, and a figure of merit was used to quantitatively evaluate the simulation results. Of the three minima in the figure-of-merit parameter space, we favor the solution characterized by an X-ray weak continuum, elevated abundances, a small column density ($\\log(N_H)\\approx 21.4$), relatively high ionization parameter ($\\log(U)\\approx -1.2$--$-0.2$), a wide range of densities ($\\log(n)\\approx 7$--$11$), and a covering fraction of $\\sim 0.15$.  The presence of low-ionization emission lines implies the disk component is optically thick to the continuum, and the \\ion{Si}{3}]/\\ion{C}{3}] ratio implies a density of $10^{10}$--$10^{10.25}\\rm\\,cm^{-3}$.  A low ionization parameter ($\\log(U)=-3$) is inferred for the intermediate-ionization lines, unless the continuum is ``filtered'' through the wind before illuminating the intermediate-line emitting gas, in which case $\\log(U)=-2.1$.  The location of the emission regions was inferred from the photoionization modeling and a simple ``toy''  dynamical model. A large black hole mass ($1.3 \\times 10^8\\,\\rm M_\\odot$) radiating at 11\\% of the Eddington luminosity is consistent with the kinematics of both the disk and wind lines, and an emission radius of $\\sim 10^4 \\, \\rm R_S$ is inferred for both.  We compare these results with previous work and discuss implications. ",
        "introduction": "In 1992, it was demonstrated by Boroson \\& Green that the optical emission line properties in the region of the spectrum around H$\\beta$ are strongly correlated with one another.  A principal components analysis allowed the largest differences among optical emission line properties to be gathered together in a construct commonly known as ``Eigenvector 1''.  The strongest differences hinged on the strength of the \\ion{Fe}{2} and [\\ion{O}{3}] emission, and the width and asymmetry of H$\\beta$.  This strong set of correlations is remarkable, as it involves correlations among the dynamics and gas properties between regions separated by vast distances. Furthermore, the same set of correlations are observed in samples selected in many different ways.  This pervasiveness leads us to believe that we are observing the manifestation of a primary physical parameter. The most favored explanation is that it is the accretion rate relative to the black hole mass onto the active nucleus.  Narrow-line Seyfert 1 galaxies are identified by their optical emission line properties.  They typically have narrow permitted optical lines (FWHM of H$\\beta<2000\\rm\\, km\\,s^{-1}$), weak forbidden lines ([\\ion{O}{3}]/H$\\beta<3$; this distinguishes them from Seyfert 2 galaxies), and frequently they show strong \\ion{Fe}{2} emission (Osterbrock \\& Pogge 1985; Goodrich 1989).  These are the same properties that define the Boroson \\& Green (1992) Eigenvector 1, and indeed, NLS1s fall at one end of Eigenvector 1.  Thus, the study of NLS1s offers an attractive research opportunity: if we can understand the origin of the emission-line properties of NLS1s, then we may be in a position to  understand AGN emission in general.  Furthermore, because as least  some of the lines in NLS1s are narrow, identification of the  frequently strongly-blended lines is less ambiguous than it is in  broad-line objects.  This paper is the second in a series of two that explores the UV emission-line properties of NLS1s.  In the first paper (Leighly \\& Moore 2004; hereafter referred to as Paper I) we introduce the topic by discussing previous work on the UV emission-line properties of NLS1s.  We then  present a detailed analysis of the {\\it HST} STIS spectra of two NLS1s, IRAS~13224$-$3809 and 1H~0707$-$495, known for their extreme X-ray properties. We found that their continua are as blue as that of the average quasar.  We observed that the high-ionization emission lines (including \\ion{N}{5} and \\ion{C}{4}) are broad (FWHM$\\approx\\rm 5000\\rm \\,km\\,s^{-1}$), and strongly blueshifted, peaking at around $2500\\rm \\,km\\,s^{-1}$ and extending up to almost $\\sim 10,000\\rm \\, km\\,s^{-1}$.  In contrast, the intermediate- and low-ionization lines (e.g., \\ion{C}{3}] and \\ion{Mg}{2}) are narrow (FWHM 1000--1900$\\rm\\,km\\,s^{-1}$) and centered at the rest wavelength. \\ion{Si}{3}] is prominent, and other low ionization lines (e.g., \\ion{Fe}{2} and \\ion{Si}{2}) are strong.  Based on these observations, the working model that we adopted considers that the blueshifted high-ionization lines come from a wind that is moving toward us, with the receding side blocked by the optically thick accretion disk, and the intermediate- and low-ionization lines are emitted in the accretion disk atmosphere or low-velocity base of the wind\\footnote{We do not know with certainty the geometrical and physical origin of the emission lines in the objects we are discussing here.  However, for simplicity, we refer to the highly blueshifted high-ionization lines as originating in the ``wind'', and the narrow, symmetric low-ionization lines as originating in the ``disk''.}.  The strongly blueshifted \\ion{C}{4 } profile suggested that it is dominated by emission in the wind. Following Baldwin et al.\\ (1996), we used the \\ion{C}{4} profile to develop a template for the wind.  We then used this template, plus a narrow and symmetric component representing the disk emission, to model the other bright emission lines. We inferred that the high-ionization lines \\ion{N}{5} and \\ion{He}{2} are also dominated by wind emission, and a part of Ly$\\alpha$ is emitted in the wind. A part of Ly$\\alpha$ and the intermediate- and low-ionization lines \\ion{Al}{3}, \\ion{Si}{3}], \\ion{C}{3}], and \\ion{Mg}{2} are dominated by disk emission.  The 1400\\AA\\/  feature, comprised of \\ion{O}{4}] and \\ion{Si}{4} was difficult to model; however, it appears to include both disk and wind emission.  IRAS~13224$-$3809 and 1H~0707$-$495 have distinctive X-ray properties among NLS1s; in order to determine whether their distinctive properties carry over to the UV, we analyzed {\\it HST} archival spectra from 14 other NLS1s with a range of two orders of magnitude in UV luminosity.  We find indeed that these two objects are extreme in this sample in the following properties: strongly blueshifted \\ion{C}{4} line, the low equivalent widths of many of the lines, particularly \\ion{C}{4} and \\ion{He}{2}, high \\ion{C}{3}]/\\ion{C}{4}, \\ion{Si}{3}]/\\ion{C}{3}], \\ion{Al}{3}/\\ion{C}{3}], and \\ion{N}{5}/\\ion{C}{4} ratios, steep $\\alpha_{ox}$, and blue UV continuum.  Correlation analysis finds a number of strong correlations.  An anticorrelation between \\ion{C}{4} asymmetry and equivalent width suggests that the line is generally composed of a highly asymmetric wind component and a narrower symmetric component.  The anticorrelation between \\ion{C}{4} asymmetry and $\\alpha_{ox}$ and $\\alpha_{u}$, and with \\ion{He}{2} suggests that UV-strong and X-ray weak continua may be associated with a wind, as would be expected if the acceleration mechanism is radiation-line driving.  The dominance of \\ion{Al}{3} and \\ion{Si}{3}] over \\ion{C}{3}] suggests possibly that the continuum is transmitted through the wind before it illuminates the intermediate- and low-ionization line emitting gas.  In this paper, our goal is to use modeling to explore the physical conditions of the line-emission gas. In \\S 2, we compute the predicted equivalent widths and ratios of the bright UV emission lines using the photoionization code {\\it Cloudy} (Ferland 2001), and then compare those with the observed values from IRAS~13224$-$3809 to constrain the ionization parameter, density, covering fraction, column density, metallicity, and shape of the photoionizing continuum.  In \\S 3, we assume a black hole mass, and use the photoionization results to estimate the distance from the central engine to the line emitting region.  We obtain the distance directly for the disk emission, but need to construct a toy dynamical model to obtain the distance for the wind.  In \\S 4, we compare with previous results,  discuss broader implications, and present a speculative scenario to tie all the results together. ",
        "conclusions": "\\subsection{Summary of Inferences from Modeling}  The photoionization modeling and toy wind model presented in this paper lead us to draw several conclusions and inferences about the optical and UV emission regions in IRAS 13224$-$3809 and 1H~0707$-$495.    In these two objects, a continuum that is overall relatively soft and deficient in X-rays illuminates gas with enhanced metallicity.  A low column-density wind is accelerated that manifests itself through blueshifted high-ionization emission lines including \\ion{C}{4}, \\ion{O}{4}], Ly$\\alpha$ and \\ion{N}{5} (the wind component).  Representative physical conditions in the wind derived from the photoionization modeling are $\\log(U)=-0.8$, $N_H=10^{21.4}\\rm\\, cm^{-2}$, $n=10^8 \\rm\\, cm^{-3}$, and covering fraction 0.15.  Using a simple radiative line-driving model, and assuming that we see the gas that is accelerated (i.e., no invisible massive substrate is present), we find that the photoionization results are consistent with the dynamical model if the black hole mass is $1.3 \\times 10^{8}\\rm\\, M_\\odot$, and the radius of foot point of the wind is $\\sim 10,000\\rm R_S$. We infer that the object is radiating at about 12\\% of the Eddington limit.  The continuum also illuminates higher density, lower velocity, higher optical depth gas, producing narrow emission lines centered at zero velocity (the disk component).  If the continuum illuminates this gas directly, the ionization parameter cannot be higher than $\\log(U)\\approx -3$.  However, if the continuum passes through the wind before illuminating this gas, the resulting continuum, already weak in hard X-rays, is now also deficient in EUV and soft X-rays above the helium ionization edge at 54 eV.  The result is that the high-ionization lines in the disk component are very weak, if present at all, and thus the high-ionization lines in the spectrum are dominated by the blueshifted wind components.  The disk component is instead dominated by intermediate-ionization lines, including Ly$\\alpha$, \\ion{Si}{4}, \\ion{Al}{3}, \\ion{Si}{3},\\ion{C}{3}], and \\ion{Fe}{3}.  Prominent low-ionization lines, including \\ion{Mg}{2}, \\ion{Fe}{2}, and \\ion{Si}{2} are also present. The line ratios in the disk component indicate a slightly high density ($\\log(n)=10.25$) but we note that a shift toward relatively lower-ionization lines is expected due to the relatively hard spectrum that the filtered continuum presents (Casebeer \\& Leighly 2004).  If the continuum is filtered through the wind, the inferred ionization parameter is $\\log(U)=-2.1$, and the corresponding radius is $14,000 \\rm R_S$.  A relatively small covering fraction of $0.05$ is inferred.  \\subsubsection{A Few Comments on the Robustness of Results}  We note that we do not claim that we can firmly rule out alternative models on the basis of the results presented here.  The scenario presented above was chosen because it appears to be most consistent with the data.  We also caution that the models are very simple.  Thus, one may be justified in asking, how well do they plausibly conform to the real situation?  The ionization model for the wind may be quite generally applicable, because we searched a very large region of parameter space: we varied the density, column density, ionization parameter, and covering fraction, as well as the metallicity and continuum shape to some extent. Furthermore, we applied a figure of merit to quantitatively gauge the applicability of our results.  Small improvements will be attained by better constraints on the metallicity that will be possible using the \\ion{O}{6} emission (e.g., Hamann et al.\\ 2002) in the {\\it FUSE} data (Leighly et al.\\ in prep.).  The {\\it FUSE} data will also be very valuable in trying to determine whether there is any ionization variation with velocity.  How good is the toy wind model?  The intent of the model was to simply make an estimate of the distance to the wind.  If radiative-line driving is the acceleration mechanism, then the magnitude of the acceleration depends directly on the effective optical depth, which in turn depends directly on the volume filling factor, and the density, and inversely on the velocity gradient (Arav \\& Li 1994).  As is, we obtain the velocity from the conservation of momentum equation.  Then, the gradient of the velocity scales with footpoint radius: a wind starting close to the AGN will reach the terminal velocity on a small size scale, implying a relatively compact emitting region and large velocity gradient, while a wind starting far from the AGN will reach terminal velocity on a larger size scale, implying a small velocity gradient.  The other important constraint comes from the requirement that integrating over the wind gives the same column density as that inferred from the photoionization modeling.  Then, combining the size scale with the column density gives the filling factor.  Since the integrated column density must be the same no matter where the wind starts, the filling factor and velocity gradient counterbalance each other to make the effective optical depth constant.  Thus, the toy model almost boils down to a simple Eddington limit problem; we have determined the radius at which radiation pressure, including radiative-line driving appropriate for the ionization state, velocity gradient and filling factor, opposing gravity can accelerate the wind to yield the inferred terminal velocity of $10,000/\\cos(\\Theta)\\rm km\\,s^{-1}$.  How could this estimate be refined?  A more general velocity law that is frequently used is the so-called beta velocity law $v(r)=v_{\\infty} (1-r_f/r)^\\beta$, where $\\beta$ is an adjustable parameter that is equal to 0.5 in the case considered above.  The velocity gradient for the beta velocity law still scales with the footpoint radius, so if we make the same assumptions as above, that the integrated column is required to be that inferred from the photoionization modeling, and radiative line-driving provides the acceleration against gravity, the resulting inferred radius for the emission region would probably not be much different unless extreme values of $\\beta$ were used.  The situation would change if we imagine that there is a significant amount of gas that we do not see that is also being accelerated.  For example, as the flow accelerates, some of the gas may become too ionized to emit in the UV; such gas is currently not accounted for.  A more massive flow requires a smaller footpoint, because a larger radiation force is needed to accelerate it to the terminal velocity. Another potential complication is that the streamlines could change direction, so that the component of the velocity parallel to the line of sight changes.  Many other complications could be imagined, and at some point, dynamical models coupled with photoionization models are needed.  In summary, the photoionization modeling results appear to be quite general, because we searched a large region of parameter space.  If we believe that acceleration mechanism is radiative-line driving, and that there is no heavy invisible substrate, we may be confident that the radius estimation for the wind is fairly accurate.  It is not possible to put specific confidence limits on the estimate without much more work, which is beyond the scope of this paper.  \\subsection{Comparison with Previous Results}  \\subsubsection{Wilkes et al.\\ 1999 \\& Kuraszkiewicz et al.\\ 2000}  Wilkes et al.\\ (1999) investigate emission line strengths and widths of 41 quasars and compare them with their continuum properties.  The quasars are a heterogeneous group chosen because they had {\\it Einstein} spectra; therefore, there is an inherent bias against objects with large values of $\\alpha_{ox}$, as noted in their paper. Also, their sample is comprised almost half of radio-loud objects. Among other things, they find an anticorrelation between EW(\\ion{C}{4}) and $\\alpha_{ox}$, as we do. In their sample, the objects with the lowest infrared (0.1--0.2 $\\mu$) luminosity are five NLS1s and 2 broad absorption-line quasars.  These objects all have relatively weak \\ion{C}{4} lines; when removed from the sample, Wilkes et al.\\ recover a Baldwin effect for \\ion{C}{4}.  Kuraszkiewicz et al.\\ (2000) follow Wilkes et al.\\ (1999) with a more detailed analysis of the NLS1s in the Wilkes et al.\\ sample, supplementing those with a few additional objects.  They analyze some of the same {\\it HST} spectra that we did in Paper I, although they do not look at IRAS~13224$-$3809 and 1H~0707$-$495.  They include {\\it IUE} spectra as well.  They do not attempt a velocity deconvolution of the lines, but rather draw inferences from the total flux of the line; certainly, the low resolution of the {\\it IUE} data prevents detailed profile studies.  Thus, they perform 1-zone {\\it Cloudy} modeling of the line flux, rather than modeling of two regions, as we do here, although in the end they infer a stratified BLR.  They infer that a density of $10^{11-12} \\rm\\, cm^{-3}$ is required to produce the weak \\ion{C}{4}, and that this density is about 10 times higher than that inferred from reverberation mapping results ($10^{11}\\rm cm^{-3}$; Peterson et al.\\ 1985).  They also infer a low ionization parameter of $\\log(U)=-3$.  As noted in \\S 2.2, our Type I solution is quite similar to that inferred by Kuraszkiewicz et al.\\ 2000.  Thus, the origin of the differences between our final solution and theirs apparently lies in primarily the much larger region of parameter space that we explored. Also, the velocity deconvolution of the line profile allows us to examine the conditions of the intermediate- and high-ionization line regions separately.  Kuraszkiewicz et al.\\ (2000) assume solar abundances; we explore a few cases of enhanced metal abundances. We also use \\ion{N}{5} which is very important for constraining abundances and the ionization parameter.  Kuraszkiewicz et al.\\ (2000) choose the spectral energy distribution from PG~1211$+$143, including the infrared. This object appears to have strong X-ray flux and therefore rather powerful heating.  As discussed in \\S 2.2, the abundances and X-ray flux are very important for determining the survival of ions under high radiation fluxes.  \\subsubsection{Richards et al.\\ 2002}  Richards et al.\\ (2002) recently reported analysis of quasar spectra from the Sloan Digital Sky Survey (SDSS; York et al.\\ 2000).  They found that the high-ionization lines, in particular \\ion{C}{4}, are systematically blueshifted with respect to the low-ionization lines \\ion{Mg}{2} and [\\ion{O}{3}].  They found an anticorrelation between the shift of \\ion{C}{4}, and its equivalent width, similar to the result reported here.  Examination of the profiles of composite spectra constructed according to \\ion{C}{4} blueshift leads them to interpret the shift/EW anticorrelation as a lack of red wing emission rather than a real shift.  This is consistent with the interpretation presented here and in Leighly (2001).  We differ with Richards et al.\\ (2002) on the interpretation of this inference.  They propose that all quasars have the same components (e.g., disk, wind, emission line regions), and the anticorrelation is a consequence of differing orientations. Specifically, they propose that blueshifted lines are seen in objects that are edge-on, like a broad absorption line quasar (BALQSO), with the difference being that the flow is not directly in the line of sight.  The receding red side is blocked by some kind of screen.  They find two pieces of supporting evidence for this hypothesis. First, they find similarities in the spectra with the most strongly blueshifted \\ion{C}{4} (their Composite D) with low-ionization BALQSOs (loBals).  Thus, if all AGN have BAL winds, and loBALs are viewed at an angle large with respect to the symmetry axis, then Composite D objects may be viewed edge on.    However, there is another explanation. Richards et al.\\ (2002) note that both the loBALs and the Composite D objects have weak \\ion{He}{2}. This is classic evidence for weak soft X-ray continuum emission.  Thus, it may be that both Composite D and the loBALs are soft X-ray weak, allowing a strong, low-ionization wind to form.  The second piece of supporting evidence is the fact that there are significantly more radio-detected (FIRST survey) objects in the composite with the least blueshift (Composite A) compared with the objects with largest blueshifts (Composite D).  Since FIRST survey will resolve out lobe emission, they interpret this as evidence that the Composite D are edge-on.  However, another speculative explanation could be that Composite A objects may have radio jets instead of winds, and the base of the jet may provide additional illumination of the disk component, producing stronger emission near the rest wavelength.  A final difference in interpretation comes for the 1400\\AA\\ feature for which they note little difference between the composites.  The analysis presented here shows that 1400\\AA\\/ is composed of \\ion{Si}{4} from the disk and \\ion{O}{4}] from the wind.  It may be that the shift of the spectra from disk-dominated to wind-dominated is accompanied by a shift from \\ion{Si}{4} to \\ion{O}{4}] in a way that maintains an approximately constant line flux and profile.  \\subsubsection{Brotherton et al.; Wills et al.}  In a series of three papers, Wills and Brotherton explored the emission line properties of a sample of high signal-to-noise restframe UV spectra from 123 high-luminosity moderate-redshift AGN (Wills et al.\\ 1993; Brotherton et al.\\ 1994a,b).  They found a number of correlations that are similar to those found here.  For example, Wills et al.\\ (1993), in studying the \\ion{C}{4} line and 1400\\AA\\/ feature, found that as the width of the \\ion{C}{4} increased, the equivalent width, kurtosis, and peak-to-continuum intensity decreased.  We did not compile the widths of \\ion{C}{4}, but we can understand the asymmetry parameter to be nearly equivalent because of our ability to explain the equivalent width/asymmetry anticorrelation observed with the sum of a strongly blueshifted wind component and a narrow, symmetric component (Paper I).  They also see the structure of the 1400\\AA\\/ feature change in a way that can be interpreted as decrease in \\ion{Si}{4} to \\ion{O}{4}] ratio with the increasing width of \\ion{C}{4}.  This seems to be consistent with our idea that when the 1400\\AA\\/ feature is dominated by disk, corresponding to narrow \\ion{C}{4}, it should be primarily composed of \\ion{Si}{4}; when dominated by wind, corresponding to broad \\ion{C}{4}, the 1400\\AA\\/ feature should have a significant component of \\ion{O}{4}].  Brotherton et al.\\ (1994a) study the relationship between \\ion{C}{4}, the 1900\\AA\\/ feature comprised primarily of \\ion{C}{3}], and \\ion{Mg}{2}.  They find that as the FWHM of \\ion{C}{4} increases, the ratio of \\ion{C}{3}] to \\ion{C}{4} increases.  This is consistent with our scenario, because when \\ion{C}{4} is dominated by wind, and therefore is measured to be broad, there is no contribution from the disk in the \\ion{C}{4} line, so \\ion{C}{3}] appears to be relatively strong.  Based on these detailed analyses, (Brotherton et al.\\ 1994b) propose that the broad lines in active galaxies are composed of two components: a component from an intermediate-line region that consists of a relatively narrow (FWHM $\\sim 2000\\rm\\, km\\, s^{-1}$) line at the systemic redshift, and a much broader (FWHM $\\sim 7000\\rm \\,km\\,s^{-1}$), somewhat blueshifted component.  This is similar to our proposed deconvolution.  However, the physical interpretation differs.  They place the VBLR very close to the central engine, whereas we place it rather far away from the central engine.  Another difference seems to be that they don't consider the possible effect of enhanced abundances or continuum shape. Also, they place the intermediate-line region very far from the central engine, whereas we show that it can be brought into a smaller radius by filtering the continuum through the wind.  Interestingly, they derive similar covering fractions for both regions as we do.  \\subsubsection{Wills et al.\\ 1999; Wills et al.\\ 2000; Shang et al.\\ 2002}  In another series of three papers, Wills and coworkers report the results of analysis of {\\it HST} spectra from a sample of 22 quasars (Wills et al.\\ 1999; Wills et al.\\ 2000; Shang et al.\\ 2002). This sample is drawn from a complete sample of 23 optically-selected PG quasars, chosen to have low redshift and low Galactic hydrogen column densities (also, Laor et al.\\ 1997a).  Wills et al.\\ 1999 present a first look at the spectra, and describe an extension of the optical-X-ray Eigenvector 1 to the UV properties. Narrow lines (including \\ion{C}{3}]) are linked with a larger \\ion{Si}{3}]-to-\\ion{C}{3}] ratio, stronger low-ionization lines, weaker \\ion{C}{4}, and stronger \\ion{N}{5}.  These are just the properties that describe the spectra of IRAS~13224$-$3809 and 1H~0707$-$495. This is not surprising, since these two objects are Narrow-line Seyfert 1 galaxies and are sufficiently blue and point-like to have been classified as PG quasars had they been located in the surveyed portion of the sky, although, like I~Zw~1, they would not have been included in the Wills et al.\\ (1999) sample because their Galactic column is too large.  Thus, the model that explains our objects can also be applied to the interpretation of the UV extension of Eigenvector 1: \\ion{C}{4} is weak because it originates in a wind with little contribution from a narrow, symmetric component.  \\ion{N}{5} is strong because abundances are high, but also possibly because there is some Ly$\\alpha$ pumping of \\ion{N}{5} in the wind.  Low-ionization lines are strong because the continuum is soft, and also perhaps because the wind filters the continuum before it illuminates the low-ionization line emitting region, leaving it no alternative but to cool by low-ionization line emission.  Filtering and a soft continuum can also influence the \\ion{Si}{3}]-to-\\ion{C}{3}] ratio.  More detailed analysis is presented in Wills, Shang \\& Yuan (2000) and especially Shang et al.\\ 2002.  The UV spectra were combined with optical spectra and a spectral principal components analysis was performed over the unprecedented large wavelength range.  They find that 79\\% of the variance lies in the first three eigenvectors, but this time the first eigenvector describes correlations between the luminosity and the relatively narrow (FWHM$\\approx 2000\\rm\\,km\\,s^{-1}$) core of the line. This is interpreted as a manifestation of the Baldwin effect (Baldwin 1977).  We discuss the Baldwin effect in the context of our interpretation in \\S 4.6.2. The second eigenvector found by Shang et al.\\ (2002) is associated with the continuum slope, and may be influenced by reddening.  The third eigenvector is associated with the traditional first eigenvector discussed by Boroson \\& Green (1992). It contains all of the classic parameters associated with the Boroson \\& Green (1992) Eigenvector 1 and the UV extension discussed by Wills et al.\\ 1999.  An interesting feature is that \\ion{C}{4} is very broad in this eigenvector.  \\subsection{Other Considerations for the Radius of the High-ionization Line Emitting Region}  In \\S 2.2, we showed that the wind emission lines are consistent with three types of photoionization solutions: Type I, a high-ionization, high-density solution; Type II, an intermediate-ionization, intermediate-density solution, and Type III, a solution characterized by a large range of densities.  We favor the Type III solution, because it encompasses a large region of the photoionization parameter space, and therefore seems to be the least fine-tuned.  But because a large range of densities were allowed by the photoionization solution, the distance from the central engine to the emission region could not be constrained without the toy dynamical model presented in \\S 2.2. That analysis indicated a radius of $\\ga 10,000 \\,\\rm R_S$ for a $1.3 \\times 10^8\\rm \\, M_\\odot$, depending somewhat on the angle the flow streamline makes with the observer.  In this section, we comment on the consistency of our large radius with reverberation mapping results, take another look at the Type I solution, and investigate the possibility of decreasing the radius of the high-ionization emission-line region for the Type III solution by ``shielding'' the wind.  \\subsubsection{The Type III Solution and Reverberation Mapping Results}  The Type III solution combined with the toy dynamical model indicates an emission region for the wind of $\\ga 10,000 \\rm \\, R_S$ for a $1.3 \\times 10^8\\rm \\, M_\\odot$, depending somewhat on the angle the flow streamline makes with the observer.  This corresponds to a distance of $3.8 \\times 10^{17} \\rm cm$, or about 150 light days.  On the face of it, an origin of the high-ionization line-emitting wind at such a large radius seems to be in conflict with reverberation mapping results, which find evidence for an ionization-stratified broad-line region in which high-ionization lines are produced quite close to the central engine.  For example, in NGC 5548, an object with UV luminosity 4 times smaller than that of IRAS~13224$-$3809 and 1H~0707$-$495, the lag of the high-ionization lines behind the continuum is on the order of 1 week or less (Korista et al.\\ 1995). In fact, there may be no conflict.  As discussed in Paper I, the high-ionization lines in NLS1s can be considered to be comprised of two components: the broad, blueshifted component that is produced by the wind, and the narrow, symmetric component that is associated with the disk.  Reverberation mapping results are probably dominated by the narrow, symmetric core that should be produced at smaller radii. Observations support this hypothesis, since in some cases the cores of the lines vary more than the wings (e.g., Wandel, Peterson \\& Malkan 1999).  This can be considered to be evidence that the broad component is produced by optically thin gas, which is less responsive to continuum changes (Shields, Ferland \\& Peterson 1995). However, it could also be produced if the wind were physically far from the central engine, as is inferred here.  \\subsubsection{The Type I Solution Reconsidered}  In this section, we take another brief look at the Type I solution. As discussed in \\S 2.2, in this solution, the density is constrained to very high values, $\\log(n)\\ga 11.5$.  The high density results in partial thermalization of Ly$\\alpha$ and \\ion{C}{4} relative to \\ion{Si}{4}; thus, the observed low equivalent width of \\ion{C}{4} is explained.  This solution is quite similar to the result obtained by Kuraszkiewicz et al.\\ (2002), who analyzed UV spectra of NLS1s but only considered the integrated line fluxes and did not analyze the line profile.  The high density and high inferred ionization parameter ($\\log(U)=-0.4$) imply a high photon flux ($\\log(\\Phi)=22.5$) which then implies that the emission region is located quite close to the central engine at $R=6.1 \\times 10^{15}\\rm\\, cm$ for $H_0=50 \\rm\\, km\\,s^{-1}\\,Mpc^{-1}$ and $q_0=0.5$ (our default choice of cosmological parameters), or $2.7 \\times 10^{15}\\rm\\, cm$ for $H_0=70\\rm\\, km\\,s^{-1}\\,Mpc^{-1}$, $\\Omega_M=0.3$, and $\\Lambda_0=0.7$. A radius of $6.1 \\times 10^{15}\\rm\\, cm$ corresponds to $160\\,R_S$ for our larger black hole mass of $1.3 \\times 10^8\\, M_\\odot$.  This value is somewhat smaller, but the same order as the inferred location of the disk wind proposed by Murray et al.\\ (1995) of $600\\rm \\, R_S$.  It would be easy to differentiate between Type I and Type III solutions from the variability of the high-ionization lines.  First, the inferred radius for the Type I solution implies a light-travel time from the central engine to the emission region of 1--2.4 days; thus the emission lines could respond rapidly to changes in the continuum. For the Type III solutions, no variability would be expected on short time scales because the light travel time is large. Furthermore, the Type I and Type III solutions are characterized by different column densities and different continuum opacities.  Both are optically thin to the continuum; however, the Type III solution is thinner, so that while the highest ionization lines, \\ion{O}{6} and \\ion{N}{5} should vary in response to the continuum, since they are produced in front of the helium ionization front, \\ion{C}{4} is produced deeper in the photoionized slab, so that it should saturate in response to flux changes (e.g, Shields, Ferland, \\& Peterson 1995). For the Type I solution, \\ion{O}{6}, \\ion{N}{5}, and \\ion{C}{4} are all produced near the illuminated face of the cloud, implying they would all vary in response to the continuum. Finally, the dynamical time scale for the Type I solution, assuming a $1.3 \\times 10^8\\,\\rm M_\\odot$ black hole, is less than a year, so one might expect to also see changes in the line profile.  There may be an additional problem with the Type I solution.  When the densities are so high that the lines are thermalized, the cooling will be dominated by continuum emission, such as the Balmer continuum (Rees, Netzer \\& Ferland 1989).  Neither IRAS~13224$-$3809 nor 1H~0707$-$495 show an especially strong Balmer jump (note that the Balmer jump can be distinguished from the strong \\ion{Fe}{2} by the wavelength of onset; e.g.\\ Dietrich et al.\\ 2002).  \\subsubsection{Can Murray et al.\\  Shielding Reduce the Radius of the High-ionization Line Emitting Region?}  In 1995, Murray et al.\\ suggested that a wind, driven by resonance scattering, could have a relatively small footpoint ($\\sim 600\\, R_S$) if there is also present a region of highly ionized gas ($U=10$) that ``shields''\\footnote{As discussed in \\S 3.2, we differentiate between a ``shielded'' continuum, which is assumed to have been transmitted through highly ionized gas (e.g., Murray et al.\\ 1995), and a ``filtered'' continuum, which is assumed to have been transmitted through the wind while ionizing and exciting it before illuminating the disk and producing the observed intermediate- and low-ionization lines.} the wind from the strongly ionizing soft X-ray emission from the central engine.  They then suggest that this wind could emit broad emission lines under the conditions that $U=1$--10.    Recalling that our best-estimated ionization parameter for the Type III solution is $\\log(U)=-0.8$, this implies that this shielding would allow the wind emission to be produced where fluxes are $100-1000$ times higher, corresponding to radii more than 10 times smaller, assuming the same density.   However, based on our experience with ``filtering'', as applied to the disk emission lines, we suspect that this scenario is untenable because there would not be sufficient high-energy photons in the shielded continuum to create and excite the high-ionization ions that we see in the wind.  In this section, we numerically test the effect of shielding on the emission from the wind.  We devised the following numerical experiment.  We transmitted our hard low-flux continuum through shielding gas having a metal abundance and a nitrogen abundance a factor of 5 and 10 over solar, respectively, and then use that transmitted continuum to illuminate the wind.  The parameters describing the shielding gas are as follows: we uniformly use $U=10$, and consider a range of column density from $10^{21}$ to $10^{22.8}\\rm \\, cm^{-2}$.  The parameters describing the wind are the following: it has uniformly $N_H^{max}=22$, and a range of ionization parameters from $\\log(U)=-1$ to $\\log(U)=0.5$.  Note that, as before, since we are interested in obtaining the radius of the emission region, the $U$s that we quote are appropriate for the unabsorbed continuum.  We determine a reference point for the simulations as follows.  The transmitted continuum will be almost identical to the incident continuum for the extreme minimum of shielding column density of $10^{21}\\rm\\, cm^{-2}$, because for $U=10$, that gas will be almost completely ionized.  Then, since the adopted value of $N_H^{max}=22$ and extreme minimum value of $\\log(U)=-1$ for the wind are close to the best wind parameters derived in \\S 2.2, the results of this combination of parameters will be close to the observed values. Therefore, we adopt the results from this extreme combination of parameters as reference values with which to compare the results for larger values of shielding $N_H$, and larger values of wind ionization parameter of $\\log(U)$.  The reference points lie in the lower left corner of each panel in Fig.\\ 13.  We plot contours of the normalized equivalent width, defined as the predicted equivalent width divided by the equivalent width obtained for the reference set of parameters described above; thus, contours marked by ``1.0'' terminate necessarily in the lower left corner of each plot.  These contours show that as the shielding column density increases, from left to right across each  plot, the highest ionization lines decrease, as the photons required to create the emitting ions are removed from the continuum.  At the same time, the gas is still subject to an intense photoionizing continuum, and it responds by increasing emission of lower ionization lines.  As the radius decreases, the line emission decreases as ions become over-ionized.  \\placefigure{fig13}  We interpret these results in terms of the Murray et al.\\ (1995) and Murray \\& Chiang (1998) scenarios as follows: shielding by highly ionized gas can certainly increase the emission of intermediate- and low-ionization lines at higher fluxes than would be possible without shielding.  However, shielding does not affect high-ionization lines in the same way, because the photons absorbed out by the highly ionized gas are those required to create ions such as N$^{+4}$ (I.P.=77 eV) and O$^{+5}$ (I.P.=113 eV).  Therefore, we conclude that shielding cannot bring the wind closer to the nucleus.  We note that this conclusion is not dependent on the column density in the wind. We examined higher column densities and found that while the flux in other lines increased, the \\ion{N}{5} line contours stayed the same, which means, not surprisingly, that with $N_H^{max}=10^{22}\\rm\\,cm^{-2}$ we integrate through the entire N$^{+4}$ zone.  \\subsection{Constraints on the Intermediate-Ionization lines}  In Paper I, we discovered several correlations among the equivalent widths and ratios of the intermediate-ionization lines \\ion{C}{3}], \\ion{Si}{3}] and \\ion{Al}{3} in our sample of NLS1s.  Specifically, we find that both the \\ion{C}{3}]/\\ion{C}{4} and the \\ion{Si}{3}]/\\ion{C}{3}] ratios are correlated with \\ion{Si}{3}] and \\ion{Al}{3} equivalent widths, but anticorrelated with the \\ion{C}{3}] equivalent width.  As discussed in Paper 1, these correlations could potentially have an origin in a variation in density, ionization parameter, or continuum shape.  To investigate quantitatively the complex interdependencies of these three parameters, we run some {\\it Cloudy} models using the hard low-flux continuum.  We examine the equivalent widths of \\ion{C}{4}, \\ion{Al}{3}, \\ion{Si}{3}], and \\ion{C}{3}] as a function of density for a constant covering fraction of 0.05, and two ionization parameters ($\\log(U)=-2$ and $-3$). We assume that the gas is ionization-bounded.  We compute the equivalent widths under two assumptions: the continuum illuminates the emitting gas directly; and the continuum illuminates the gas after passing through the wind, as discussed in \\S 2.3.  The results are displayed in Fig.\\ 14.  \\placefigure{fig14}  Fig.\\ 14 shows the decrease in equivalent width as a function of density expected for the semiforbidden lines \\ion{C}{3}] and \\ion{Si}{3}].  The permitted line \\ion{C}{4} shows little dependence on density, while \\ion{Al}{3} shows a moderate increase, possibly due to its increasing role in cooling in the \\ion{C}{3}] and \\ion{Si}{3}] forming region.  This seems to suggest that if an increase in density drives the correlations, the \\ion{C}{3}] to \\ion{C}{4} ratio should decrease, contrary to the observations.  Filtering the continuum through the wind results in a dramatic decrease in \\ion{C}{4}.  In fact, there is a larger decrease between the filtered and nonfiltered continua for $\\log(U)=-2$ than there is between either continuum with $\\log(U)=-2$ and $\\log(U)=-3$.  For $\\log(U)=-2$, there is little effect on the intermediate-ionization lines \\ion{Al}{3}, \\ion{Si}{3}] and \\ion{C}{3}], although we see a slight increase for \\ion{Si}{3}] for the filtered continuum.  For $\\log(U)=-3$ and a filtered continuum, all the lines become weaker, with \\ion{C}{4} decreasing the most.  These simple simulations show that the origin of the correlations among the intermediate ionization line properties could be a filtered continuum, or it could originate from a trend in ionization parameter, but an origin in density variations does not seem to be supported.  \\subsection{Ideas and Speculations about the Emission Regions in NLS1s}  In this paper and in Paper I, we learned many things about the UV emission lines from two particular NLS1s, 1H~0707$-$495 and IRAS~13224$-$3809, and about the properties of those two objects compared with a heterogeneous sample of NLS1s.  In this section, we attempt to use these results to construct a scenario for the emission regions in NLS1s.  Perhaps the simplest scenario for the emission-line geometry consists of an accretion disk with a wind.  The blueshifted part of the lines is produced in the wind, while the portion centered at the rest wavelength is produced in the base of the wind, or in the accretion disk itself.  In some objects, there is no significant wind, in which case all the line emission comes from the base of the wind. In others, the wind emission is relatively strong; in these objects, high ionization lines have a blue wing arising from the wind.  Why do some objects have winds and others do not?  If the winds are accelerated by radiation-line driving, then objects with steep $\\alpha_{ox}$  may have winds, while objects with flat $\\alpha_{ox}$ may not. This is because a strong UV continuum is necessary to drive the wind; however, too much soft X-ray emission will easily overionized the gas, destroying the resonance-scattering ions.   A high metallicity may also encourage a wind, since that could potentially contribute more resonance-scattering ions.  This is the simplest scenario; however, it does not explain a number of aspects of the data.  It does not explain why \\ion{C}{4} is so weak in the wind-dominated NLS1s like IRAS~13224$-$3809 and 1H~0707$-$495; the scenario above only allows blue wings; the \\ion{C}{4} line core, and other high-ionization line cores,  produced in the base of the wind should still be very strong, and should dominate the line emission.  It also does not explain the extreme line ratios in the wind-dominated NLS1s, including relatively large \\ion{C}{3}]/\\ion{C}{4},   \\ion{Si}{3}]/\\ion{C}{3}], and \\ion{Al}{3}/\\ion{C}{3}] ratios.   Additional tweaks of the scenario are necessary to explain these properties.  How to explain the weak \\ion{C}{4} core in the wind-dominated NLS1s? Increasing the density will not help.  Although the high \\ion{Si}{3}]/\\ion{C}{3}] ratio suggests a higher density in these objects, the inferred density, well constrained by the \\ion{Si}{3}]/\\ion{C}{3}] to be $\\log(n)\\approx 10.25$, is not nearly high enough to suppress \\ion{C}{4}, a permitted line.  This point was discussed in detail in \\S 4.4.  Another possibility is to simply move the emission region very far from the nucleus such that the dominant ionization state for carbon in the base of the wind is C+2.  This possibility is explored in the context of photoionization modeling in \\S 2.3.  There are two potential objections to this.  First, the inferred distance to the emission region is very large, so that for reasonable black hole masses, the Keplerian velocities are quite low.  A more important problem would be, why is the emission region exceptionally far from the nucleus in the wind-dominated NLS1s alone?  One possible answer could be to associate the emission region with the breakup radius of the accretion disk (Collin \\& Hur\\'e 2001), as that should vary with accretion rate (Hur\\'e 2000).  However, this does not simply solve the problem, since the breakup radius depends inversely on accretion rate; the larger the accretion rate with respect to Eddington, the smaller the radius, just the opposite dependence required if we associate the wind-dominated NLS1s with a higher accretion rate.  In \\S 2.3 we present a scenario that explains these observational results naturally.  If the continuum is transmitted through the wind before it illuminates the intermediate- and low-ionization line emitting gas, it will lack sufficient photons in the helium continuum to excite high-ionization lines.  Thus, the presence of a wind naturally results in a weak core for the high-ionization lines.  The softening of the continuum also tends to create and excite ions with lower ionization potentials; this can partially explain the dominance of \\ion{Si}{3}] and \\ion{Al}{3} in the intermediate-ionization line-emitting region.  \\subsubsection{A Speculative Scenario}  The analysis presented in this paper, and in Paper I, leads us to infer that the spectral energy distribution is a primary factor in determining both the ionization and the dynamics of the line-emitting gas in AGN.  A continuum strong in the UV and weak in X-rays can drive a wind by resonance scattering without overionizing the ions required for the line driving; this results in blueshifted high-ionization lines.  This wind may then filter the continuum before it illuminates lower-velocity gas emitting intermediate- and low-ionization lines. The filtered continuum lacks high energy photons, so the low-velocity gas cools by emitting particularly lower-ionization lines, including such low-ionization lines as \\ion{Fe}{2} and \\ion{Si}{2}.  Conversely, if the X-rays are strong compared with the UV, a wind is not formed because the gas is overionized before it can be accelerated.  This unfiltered continuum illuminates the lower-velocity gas directly. This continuum is strong in the extreme UV and in X-rays, so higher-ionization lines, including \\ion{C}{4}, are produced in the low-velocity gas.  The scenario outlined above begs the question: what determines the continuum to begin with?  The continuum may be directly controlled by intrinsic parameters such as the black hole mass and accretion rate, but how?  We present a speculative scenario to explain this as follows.  Under conditions of moderate accretion rate, the accretion disk may be ``thin'', at least at large radii (e.g., Frank, King \\& Raine 1992).  The thin disk is able to radiate the the power that it generates, is optically thick and geometrically thin, and its spectrum is predicted to be $F(\\nu) \\propto \\nu^{1/3}$.  Increasing the accretion rate causes the characteristic temperature to increase; increasing the black hole mass causes the characteristic temperature to decrease (e.g., Ross, Fabian \\& Mineshige 1992).  Thus, the thin disk can produce a blue continuum in the optical and UV that is harder for smaller black hole masses and higher accretion rates.  At smaller radii, the thin disk is predicted to generate more energy than can be radiated through its surface.  Radiation pressure should become important, causing the disk to puff up and become geometrically thick.  The simplest radiation-pressure dominated disks are known to be unstable; however, detailed models show that this solution may be still viable  (e.g., Agol et al.\\ 2001; Blaes \\& Socrates 2001).  X-rays are thought to be produced close to the black hole by an optically thin, hot corona that intercepts optical/UV photons and upscatters them.  The amount of X-ray emission depends on the amount of corona; specifically, the covering fraction and optical depth.  The amount of corona may be depend on how much accretion energy is diverted into producing the corona, by, for example, magnetic reconnection of buoyant loops generated in the disk that escape through the surface.  Some recent models predict that the amount of energy produced in the corona should also depend upon accretion rate (e.g., Liu et al.\\ 2002), so that high accretion rate objects have a smaller fraction of energy emitted in the corona compared with the disk. This scenario was used by Bechtold et al.\\ (2003) to  explain the steepening of $\\alpha_{ox}$ with luminosity.  So, combining the thin disk, the radiation-supported thick disk,  and the corona, we suggest a speculative geometrical scenario that is illustrated in Fig.\\ 15 (not to scale).  The key feature that potentially explains the behavior of the UV emission lines in NLS1s is the spectrum emitted by the outer part of the geometrically thick disk, because that is the spectrum that initially illuminates the line-emitting gas.  If that continuum is dominated by photons below 13.6 eV, it can accelerate a wind without overionizing it.  That accelerated wind, which may originate in the region where the disk becomes gravitationally unstable, reaches a modest velocity, yet does not emit very much line emission.  As it rises above the lip of the radiation-pressure supported torus, it is exposed to the full UV continuum, is quickly accelerated and emits.  An analogy may be wind blowing snow off a cornice.  Beyond the wind may lie the clumps from the gravitationally unstable disk that would emit the lower-velocity, intermediate and low-ionization lines  when they are illuminated by the continuum filtered through the wind.  \\placefigure{fig15}  In contrast, if the spectrum of the outer part of the geometrically thick disk emits very much radiation at energies above 13.6 eV, the potential reservoir for the wind gas is exposed to the full ionizing continuum and is overionized before it can be accelerated. Furthermore, the thin disk will be illuminated by a strong photoionizing continuum, and it will cool by emitting strong permitted lines including \\ion{C}{4}.  This speculative scenario explains the emission lines, but it can also explain the continuum.  Quasars that have strong blueshifted lines tend to have blue optical and UV continua.  This may be a consequence of the optical/UV emission from the thin disk dominating to a small radius.  An example of such an object is PHL~1811, a very luminous quasar (Leighly et al.\\ in prep.)  Conversely, objects that do not have blueshifted high-ionization lines may have less contribution from the optically thick, geometrically thin disk; their continuum is stronger in the extreme UV and may peak in the soft X-rays.  An example of such an object is RE~1034$+$39 (Casebeer \\& Leighly 2004).  \\subsection{Other Considerations}  \\subsubsection{Influence of X-ray Properties}  We have noted already that IRAS~13224$-$3809 and 1H~0707$-$495 have extreme X-ray properties compared with other NLS1s, as discussed in Leighly 1999b.  These properties include the highest-amplitude X-ray variability and the most prominent X-ray soft excesses in their {\\it ASCA} spectra.  As discussed above, the photoionization modeling shows that a larger region of parameter space is available when the X-ray flux is low.  We justified our assumption of low X-ray flux using the fact that 1H~0707$-$495 has been frequently found in very low flux states, a factor of 10 lower than observed when the {\\it ASCA} spectrum was made (Leighly et al.\\ 2002; Leighly et al.\\ in prep.). Furthermore, in Paper I, we showed that there is an anticorrelation between the asymmetry of the emission line and $\\alpha_{ox}$, such that X-ray weak objects have blueshifted profiles.  It is also possible that the high-amplitude variability plays a role in the production of the emission lines.   If the high-amplitude variability extends from the X-ray into the UV, a variable radiation line-driving force may be present, or a variable photoionizing continuum.  A variable line-driving force may be interesting because it could produce shocks in the wind that may result in density enhancements and may help formation of filaments dense enough to emit the lines that we see.  A variable photoionizing continuum can effect the ionization balance of the emitting region; Nicastro et al.\\ (1999) found that the effect was such that the emission region was overionized compared with ionization expected for the particular ionization parameter.  There is no data yet that supports rapidly variable UV emission in these two objects.  Young et al.\\ 1999 do not see any optical variability on short time scales over three nights in IRAS~13224$-$3809.  No significant variability in the near-UV was detected during a 20,000 second {\\it XMM-Newton} OM observation. However, observed correlated optical and X-ray variability in NGC~5548 led Uttley et al.\\ (2003) to propose that relatively luminous Seyferts may be more variable in the UV than less luminous objects because of relatively cooler accretion disk expected for a larger black hole and corresponding shift of the UV emission region to smaller $R/R_S$ where dynamical time scales may be short. Thus, variability in the line-driving continuum or photoionization continuum may be possible.  \\subsubsection{The Baldwin Effect}  It is well known that the equivalent widths of emission lines in quasars are anticorrelated with their luminosity (the Baldwin effect; The Baldwin 1977; see Osmer \\& Shields 1999 for a review).  Baldwin effect has been variously attributed to variations in the broad-line cloud covering fraction, ionization parameter, and inclination.  Perhaps the most promising explanation is in terms of the shape of the continuum, because it can explain the fact that the higher-ionization lines experience a steeper decrease with luminosity than the lower-ionization lines.  Another interesting feature of the Baldwin effect is that it is stronger in the line cores. This fact was also observed in the spectral PCA analysis of the {\\it HST} sample of PG quasars (Shang et al.\\ 2002).  It is possible that our scenario can contribute to the interpretation of the Baldwin effect.  Luminosity may be correlated with the black hole mass, predicting a softer spectrum for larger black holes for a fixed accretion rate relative to Eddington (e.g., Ross, Fabian \\& Mineshige 1992).  Thus, more luminous objects, because of their softer spectrum, may produce a wind that may filter the continuum, reducing the availability of energetic ionizing photons to the low-velocity gas, reducing the line core.  Interestingly, Shang et al.\\ (2002) find that broad \\ion{He}{2}$\\lambda 4686$ is present in their first eigenvector that they interpret as having an origin in the Baldwin effect, and is correlated with the strength of the narrow component of the lines.  As discussed above, \\ion{He}{2} is the classic indicator of a continuum strong in soft X-rays.  It is also interesting to note that this interpretation contrasts with that of Shields, Ferland \\& Peterson (1995); they speculate that the covering fraction of the optically-thin component depends inversely on luminosity.  \\subsubsection{BALQSOs}  In this paper, we show that a number of the emission line properties of NLS1s are plausibly related to the presence or absence of a radiatively-driven wind.  The other class of active galaxy in which there is good evidence for a radiatively-driven wind is the class of broad absorption-line QSOs (BALQSOs).  Of course, there is an important difference between BALQSOs and NLS1s: in BALQSOs, the wind is in the line of sight to the nucleus, causing sometimes spectacular absorption lines to be imprinted upon the spectrum, while in NLS1s, the wind is not in the line of sight.  This difference could be simply a matter of orientation.  One sees emission lines regardless of the viewer orientation (disregarding anisotropic emission, for the moment), so if the wind has some effect on the emission line properties of NLS1s, then it may produce the same effects on the emission lines in BALQSOs.  Weymann et al.\\ (1991) present a comprehensive comparison between the emission-line properties of BALQSOs and non-BAL quasars.  As discussed in the introduction to that paper, their work was motivated by various reports of the following differences between BALQSOs and ordinary quasars: 1) \\ion{Fe}{2} is stronger; 2) \\ion{Al}{3} is stronger; 3) \\ion{C}{4} is weaker; 4) \\ion{N}{5} is stronger.  We note that these properties describe the differences between NLS1s and ordinary broad-line quasars.  Weymann et al.\\ (1991) then present a very careful and conservative analysis of samples of BALQSOs and non-BALQSOs. While they conclude that, excluding the low-ionization BALQSOs\\footnote{Other authors have remarked upon the similarity between NLS1s and the low-ionization BALQSOs, objects which have very strong \\ion{Fe}{3} and \\ion{Fe}{2} (e.g.\\ Leighly et al.\\ 1997; Boroson \\& Meyers 1992, for the case of I~Zw~1).}, the emission lines are very much the same in BALQSOs and non-BALQSOs.  However, a few subtle differences remain.  There is an enhancement of \\ion{N}{5}, and there is an enhancement around \\ion{Al}{3}, which they attribute to \\ion{Fe}{2} or \\ion{Fe}{3}. They also find a correlation between \\ion{Fe}{2} and ``balnicity'' index, a measure of the BAL-ness of a quasar.  In our scenario, these properties would be explained as follows.  The \\ion{N}{5} enhancement arises in the BAL flow, and it is perhaps enhanced by scattering (as noted by Weymann et al.\\ 1991). However, we also speculate that the strong low-ionization lines result from illumination by a continuum filtered through the BAL flow. We point out that in our scenario, in general the lines such as \\ion{C}{4} are composed of both wind and disk components, and thus differences between objects with and without winds may not be very spectacular.  If the radiatively-driven winds in BALQSOs are related to the winds that we discuss in this paper, we  might also expect that BALQSOs should be deficient in X-ray emission.  Measuring the X-ray emission from BALQSOs is difficult, because the X-rays are usually absorbed, and it is difficult to deconvolve the spectrum and the absorption robustly without making assumptions about the continuum shape. Sometimes, good evidence for normal spectral energy distributions have been found (e.g., Gallagher et al.\\ 2002) .  However, there is at least one instance in which evidence that the BALQSO is intrinsically X-ray weak has been found (Sabra \\& Hamann 2001).  \\subsubsection{Early Universe Counterparts of NLS1s}  Recently, several investigators have sought the early universe counterparts of NLS1s (e.g., Mathur 2000).  This is important, because now a significant number of quasars with redshifts $z>4$ have been discovered.  This epoch corresponds to $>90$\\% of the age of the universe; therefore, these quasars are young.  The accretion rate in these objects is of special interest.  Luminous quasars must have large black holes, and in order to grow large in such a short time, they ought to be accreting at a rapid rate.  The properties of the objects discussed in this paper may give us a hint about what kind of objects at high redshift correspond to Narrow-line Seyfert 1 galaxies.  There are some objects that have been found that have extremely strong and narrow emission lines (e.g., Constantin et al.\\ 2002).  The lines in these objects have huge equivalent widths; clearly the emitting region is very well illuminated.  If we consider the fact that the Baldwin effect slopes for NLS1s typically lies significantly below that of other quasars (Paper I; Wilkes et al.\\ 1999), then that argues against these high-equivalent width objects being early-universe counterparts of NLS1s.  On the other hand, several high-redshift objects have been discovered in the Sloan Digital Sky Survey that have very blue continua and no apparent emission lines (Fan et al.\\ 1999; Anderson et al.\\ 2001). One of the properties of NLS1s is their relatively low line equivalent width and frequently blue continua.  Thus, the line-less quasars seem much more similar to the high-luminosity NLS1s, and therefore they may be the sought-after early-universe counterparts (Leighly, Halpern, \\& Jenkins 2004; Leighly et al.\\ in prep.)."
    },
    "0402/astro-ph0402208_arXiv.txt": {
        "abstract": "The variation of dark energy density with redshift, $\\rho_X(z)$, provides a critical clue to the nature of dark energy.  Since $\\rho_X(z)$ depends on the dark energy equation of state $w_X(z)$ through an integral, $\\rho_X(z)$ can be constrained more tightly than $w_X(z)$ given the same observational data.  We demonstrate this explicitly using current type Ia supernova (SN Ia) data [the Tonry/Barris sample], together with the Cosmic Microwave Background (CMB) shift parameter from CMB data (WMAP, CBI, and ACBAR), and the large scale structure (LSS) growth factor from 2dF galaxy survey data. We assume a flat universe, and use Markov Chain Monte Carlo (MCMC) technique in our analysis. We find that, while $w_X(z)$ extracted from current data is consistent with a cosmological constant at 68\\% C.L., $\\rho_X(z)$ (which has far smaller uncertainties) is not.  Our results clearly show the advantage of using $\\rho_X(z)$, instead of $w_X(z)$, to probe dark energy. ",
        "introduction": "Recent observations of type Ia Supernovae \\citep{Riess98,Perl99} indicate that the universe is accelerating.  A fundamental quest in physics and cosmology is to identify the nature of the ``dark energy'' driving this acceleration. Possibilities include: (1) a cosmological constant, (2) a time dependent vacuum energy, or a scalar field known as ``quintessence'' that evolves dynamically with time \\citep{fafm,peebles88,Wett88,frieman,caldwell98,Zlatev99} \\footnote{See \\cite{Pad03} and \\cite{peebles03} for reviews with more complete lists of references).} or (3) modified Friedmann equation, e.g. the Cardassian models \\citep{freeselewis,freese03,mpcard,Wang03}, that could result as a consequence of our observable universe living as a 3-dimensional brane in a higher dimensional universe.  Other proposed modifications to the Friedmann equation include \\cite{Parker99,ddg,Bilic02,Ahmed02,Capo03,Carroll03b,Meng03,Puet04}.  The various dark energy models produce dark energy densities $\\rho_X(z)$ with different redshift dependences.  Hence, in order to differentiate between dark energy models, it is important that we allow the dark energy density to be an arbitrary function of redshift $z$ \\citep{Wang01a,Wang01b,Wang03}.   A powerful probe of dark energy is type Ia supernovae (SNe Ia), which can be used as cosmological standard candles to measure how distance depends on redshift in our universe.  The luminosity distance $d_L(z) = (1+z) r(z)$, with the comoving distance $r(z)$ given by \\begin{equation} r(z)= cH_0^{-1} \\int_0^z \\frac{dz'}{E(z')}, \\end{equation} with \\begin{equation} \\label{eq:E(z)} E(z) \\equiv \\left[ \\Omega_m(1+z)^3 + \\Omega_k (1+z)^2 + \\Omega_X \\rho_X(z)/\\rho_X(0) \\right]^{1/2}, \\end{equation} where $\\Omega_k \\equiv 1-\\Omega_m-\\Omega_X$, and $\\rho_X(z)$ is the dark energy density.  The dark energy equation of state, $w_X(z)$, is related to $\\rho_X(z)$ as follows \\citep{Wang01a}: \\begin{equation} \\label{eq:wrhoprime} w_X(z) =\\frac{1}{3}(1+z) \\frac{\\rho'_X(z)}{\\rho_X(z)} -1, \\end{equation} so that \\begin{equation} \\label{eq:rhoprimew} \\frac{\\rho_X(z)}{\\rho_X(0)} = \\exp\\left\\{ \\int_0^z \\frac{3 [1+w_X(z)]}{1+z} \\right\\}. \\end{equation} One can see that it is easier to extract $\\rho_X(z)$ from the data than to extract $w_X(z)$.  To obtain the dark energy density directly, one need only take a single derivative of the luminosity distance, whereas to extract $w_X(z)$, one needs to take a second derivative as well; from Eq.(\\ref{eq:wrhoprime}) one can see that $w_X(z)$ is on the same footing as $\\rho_X'(z)$.  Specifically, \\cite{Wang01a} argued that $\\rho_X(z)$ should be preferred since it suffers less from the smearing effect (due to the multiple integrals that relate $w_X(z)$ to $d_L(z)$) that makes constraining $w_X(z)$ extremely difficult \\citep{Maor01,barger01}.  \\cite{Tegmark02} came to the same conclusion. However, researchers have generally chosen to parametrize dark energy using its equation of state $w_X(z)$. Some have used $H(z)=H_0 E(z)$ (for example, see \\cite{Kujat02,Daly03,peri}, and references therein), which is similar to $\\rho_X(z)$, but measurements of which are not as straightforward to interpret, since $E(z)$ depends on $\\Omega_m$ (see Eq.(\\ref{eq:E(z)})).  In this paper, we explicitly demonstrate the advantage of using $\\rho_X(z)$, instead of $w_X(z)$, to probe dark energy. Sec.2 contains a comparison of $w_X(z)$ and $\\rho_X(z)$ parametrizations using current SN Ia, CMB, and LSS data. We give a recipe for parametrizing dark energy using $\\rho_X(z)$ in Sec.3. Sec.4 contains a summary and discussions. ",
        "conclusions": "The critical first step in solving the mystery of dark energy is to determine whether the dark energy density $\\rho_X(z)$ varies with time.\\citep{Wang01a} A definitive answer to this question can have profound implications for particle physics and cosmology.  Our main result is that one can learn more information by reconstructing $\\rho_X(z)$ rather than $w_X(z)$ from the data.  The two quantities are related by an integral, which in the case of $w_X(z)$ smears out much of the information one could otherwise learn. We show this explicitly by using a combination of SN Ia data from the Tonry/Barris sample as well as CMB (WMAP, CBI, and ACBAR) and large scale structure (2dF) data.  At 95\\% CL, both the $w_X(z)$ and $\\rho_X(z)$ reconstructions are consistent with a cosmological constant.  However at 68\\% CL, the $\\rho_X(z)$ reconstruction has smaller uncertainties and hence shows information that the $w_X(z)$ reconstruction cannot: the $\\rho_X(z)$ reconstruction is {\\it not} consistent with a time-independent dark energy, and the dark energy density appears to be increasing with redshift.  Future data will be required to resolve this question.  We have shown definitively the advantage of the $\\rho_X(z)$ parametrization over the $w_X(z)$ parametrization in determining the time-variation of $\\rho_X(z)$.  To help others apply the $\\rho_X(z)$ parametrization, we have given a recipe for using the $\\rho_X(z)$ parametrization in data analysis to probe dark energy (see Sec.3). Our methodology should be very useful in all data analysis aiming at unraveling the nature of dark energy."
    },
    "0402/astro-ph0402514_arXiv.txt": {
        "abstract": "A new scenario for neutron-star cooling is proposed, based on the correspondence between pion condensation, occurring in neutron matter due to critical spin-isospin fluctuations, and the metal-insulator phase transition in a two-dimensional electron gas.  Beyond the threshold density for pion condensation, where neutron-star matter loses its spatial homogeneity, the neutron single-particle spectrum acquires an insulating gap that quenches neutron contributions to neutrino-production reactions and to the star's specific heat.  In the liquid phase at densities below the transition point, spin-isospin fluctuations are found to play dual roles.  On the one hand, they lead to a multi-sheeted neutron Fermi surface that extends to low momenta, thereby activating the normally forbidden direct-Urca cooling mechanism; on the other, they amplify the nodeless $P$-wave neutron superfluid gap while suppressing $S$-wave pairing.  In this picture, lighter stars without a pion-condensed core experience slow cooling, while enhanced cooling occurs in heavier stars through direct-Urca emission from a narrow shell of the interior. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0402/hep-th0402190_arXiv.txt": {
        "abstract": "We perform a thorough phase-plane analysis of the flow defined by the equations of motion of a FRW universe filled with a tachyonic fluid plus a barotropic one. The tachyon potential is assumed to be of inverse square form, thus allowing for a two-dimensional autonomous system of equations. The Friedmann constraint, combined with a convenient choice of coordinates, renders the physical state compact. We find the fixed-point solutions, and discuss whether they represent attractors or not. The way the two fluids contribute at late-times to the fractional energy density depends on how fast the barotropic fluid redshifts. If it does it fast enough, the tachyonic fluid takes over at late times, but if the opposite happens,  the situation will not be completely dominated by the barotropic fluid; instead there will be a residual non-negligible contribution from the tachyon   subject to restrictions coming from nucleosynthesis. ",
        "introduction": "A phase of accelerated inflation in the early stages of our universe is favored by first-year WMAP data \\cite{WMAP}. Thus, it seems  inflation is here to stay as the dominant paradigm for structure formation. The quest for a string theory motivated explanation of cosmological inflation has resulted in the emergence of the proposal of inflation driven by a tachyon field. The idea strongly relies in the possibility of describing tachyon condensates in terms of perfect fluids within string theories \\cite{sen}. A plethora of papers studying cosmological consequences of such  fluids have appeared since, some in the framework of general relativity  \\cite{power-law,others, exponential, inverse}, some others in the brane-world scenario \\cite{brane}.  As happens with standard scalar fields, one's favorite inflationary behaviour is tailored by {\\it ad hoc} choices of the initial conditions and the shape of the potential, but it is important to investigate up to what extent the features of the model depend on those choices. One way to address that problem is to consider tachyon field dynamics, because  for a given potential such an analysis will provide us with constraints on the initial conditions.  The stability of tachyonic inflation  against changes in initial conditions  has been studied  for an exponential potential  \\cite{exponential} and for the inverse power-law potential \\cite{inverse}. Exact solutions for a purely tachyonic matter content with an inverse square potential are known \\cite{power-law}, but no solutions exist for cases which combine tachyonic and barotropic fluids, so a dynamical systems approach may be relevant. Interestingly, the inverse square potential plays the same role for tachyon fields as the exponential potential \\footnote{The general solution to the Einstein equations for a FRW spacetime with an exponential potential  was given in \\cite{solution}.} does for standard scalar fields \\cite{liddle,firstsem,exp,thirdsem}. On the one hand, those are the potentials that give power-law solutions.  On the other hand, only those potentials allow constructing a two-dimensional autonomous system  \\cite{thirdsem} using the evolution equations, whereas for any other potential the number of dimensions will be higher if the system is to remain autonomous.  As compared to \\cite{inverse},  we throw in one more ingredient in our study, because  we allow for the presence of a barotropic fluid, along with the tachyon fluid. The crucial consequence is the appearance of fixed-point solutions in which  the two fluids redshift at the same rate (tracking behaviour \\cite{liddle}), so that there is some sort of equilibrium.  Tracking solutions are particularly interesting because their dynamical effects mimic a decaying cosmological constant (see \\cite{thirdsem,firstsem,secondsem} for seminal references). Now, the fine-tuning problems posed by a cosmological constant would be waived precisely because of the independence on the initial conditions. Nevertheless, the contribution of such relics to the fractional energy density are bounded by nucleosynthesis.  On top of those interesting features, tracking solutions act as attractors at large, which means that the system is not any picky in what initial conditions are regarded.  In Section II we study the phase-plane, find its fixed points and characterize them. In Section III we discuss the cosmological consequences of the attractor solutions: in subsection III A we consider tachyon dominated  solutions, whereas in subsection III B we discuss the tracking ones. Finally, in Section IV we outline our main conclusions and future prospects. ",
        "conclusions": "The evolution equations of a spatially flat FRW universe containing a barotropic fluid plus a tachyon $T$ with an inverse-square potential $V(T)=\\beta T^{-2}$ define a two-dimensional flow. The evolution of such models has been investigated by studying the orbits of that flow in the physical state, which is in this case a subset of the Euclidean plane. The Friedmann constraint, combined with a careful choice of coordinates, renders this subset compact.  We have shown that the energy density of the tachyon dominates at late times for  $\\gamma>\\alpha(\\sqrt{\\alpha^2+4}-\\alpha)/2$, where $\\gamma$ is the barotropic index of the fluid and $\\alpha=4/3\\beta$. In constraint, for $\\gamma<\\alpha(\\sqrt{\\alpha^2+4}-\\alpha)/2$, the barotropic fluid does not dominate completely and the contribution of tachyonic energy density to the total one is not negligible. Nucleosynthesis imposes, then, tight bounds on the admissible values of $\\alpha$, but such restrictions can be relaxed if the locus of the initial solution is far from that of the tracking one and in a region where $\\Omega_T\\ll1$ (close to $O$), and only reaches $Q$ in the course of the evolution.  Finally, a possible generalization of this work would be considering generalized tachyon cosmologies like those presented in \\cite{Chimento}."
    },
    "0402/astro-ph0402272_arXiv.txt": {
        "abstract": "We present preliminary trigonometric parallaxes and proper motions for 22 L dwarfs and 18 T dwarfs measured using the ASTROCAM infrared imager on the U.S. Naval Observatory (USNO) 1.55--m Strand Astrometric reflector. The results presented here are based on observations obtained between September 2000 and November 2002; about half of the objects have an observational time baseline of $\\Delta$t~=~1.3 yr and half $\\Delta$t~=~2.0 yr. Despite these short time baselines, the astrometric quality is sufficient to produce significant new results, especially for the nearer T dwarfs. Seven objects are in common with the USNO optical CCD parallax program for quality control and seven in common with the ESO 3.5--m NTT parallax program. We compare astrometric quality with both of these programs. Relative to absolute parallax corrections are made by employing 2MASS and/or SDSS photometry for reference frame stars. We combine USNO infrared and optical parallaxes with the best available CIT system photometry to determine $M_J$, $M_H$, and $M_K$ values for 37 L dwarfs between spectral types L0 to L8 and 19 T dwarfs between spectral types T0.5 and T8 and present selected absolute magnitude versus spectral type and color diagrams, based on these results. Luminosities and temperatures are estimated for these objects. Of special interest are the  distances of several objects which are at or near the L--T dwarf boundary so that this important transition can be better understood. The previously reported early-mid T dwarf luminosity excess is clearly confirmed and found to be present at J, H, and K. The large number of objects that populate this luminosity excess region indicates that it cannot be due entirely to selection effects. The T dwarf sequence is extended to $M_J$~$\\approx$~16.9 by 2MASS J041519$-$0935 which, at d~=~5.74 pc, is found to be the least luminous [log(L/L$_{\\odot}$)~=~-5.58] and coldest (T$_{\\rm eff}$~$\\approx$~760 K ) brown dwarf known. Combining results from this paper with earlier USNO CCD results we find that, in contrast to the L dwarfs, there are no examples of low velocity (V$_{tan}$$<$ 20 km $s^{-1}$) T dwarfs. This is consistent with the T dwarfs in this study being generally older than the L dwarfs. We briefly discuss future directions for the USNO infrared astrometry program. ",
        "introduction": "After the era of photographic proper motion surveys (e.g. Luyten 1979) revealed late M stars close to the limit of stellar hydrogen burning, a long, and largely frustrating, pursuit of sub--stellar objects was begun by many researchers. These efforts were motivated by the overarching desire to understand the Galactic mass and luminosity distributions of putative objects that would bridge the gap between the lowest mass stars and giant planets and by the fact that no theory of star formation could be considered complete without accounting for the mass function of such objects. Becklin \\& Zuckerman (1988) identified GD 165B as the first object clearly cooler than an M dwarf, followed several years later by the discovery of the first `methane dwarf', Gliese 229B \\citep{nak95,opp95}; an object cold enough that its spectrum shows strong methane absorption, similar to the giant gas planets. These `brown dwarfs' became the prototypes for L dwarfs \\citep{kir99,mar99} and T dwarfs \\citep{bur02a,geb02}, respectively.  It was not, however, until deep, large--scale optical surveys (the Sloan Digital Sky Survey [SDSS]\\footnote{\\bf www.sdss.org}; \\citealp{york00,abazajian03}) and near--infrared surveys (the Two Micron All Sky Survey [2MASS\\footnote{\\bf www.ipac.caltech.edu/2mass}; Skrutskie et al. 1997] and the Deep Near Infrared Survey of the Southern Sky [DENIS\\footnote{\\bf cdsweb.u-strasbg.fr/denis.html}; Delfosse et al. 1997, Epchtein 1997]) of the sky were undertaken that significant numbers of field brown dwarfs were revealed. L dwarfs from these surveys have been identified by many authors \\citep{del97,kir99,kir00,fan00,haw02,geb02,sch02}, as have T dwarfs \\citep{bur99,str99,tsv00,leg00,bur02a,geb02,bur03a,kna04}, amongst many others. Comprehensive summaries of field brown dwarf discoveries are maintained at the web sites maintained by Kirkpatrick for L dwarfs ({\\bf spider.ipac.caltech.edu/staff/davy/ARCHIVE}), by Burgasser for T dwarfs ({\\bf www.astro.ucla.edu/adam/homepage/research/tdwarf}), and by Leggett for both L and T dwarfs ({\\bf www.jach.hawaii.edu/skl/LTdata.html}). Currently there are more than 250 L dwarfs and 58 T dwarfs known.  Unlike stars, brown dwarfs are not massive enough to sustain continuous hydrogen fusion in their cores, but cool continually from their birth. Somewhere between early and mid--L is the crossover between hydrogen--burning stars and brown dwarfs. Unfortunately, other than objects in clusters \\citep{bas00}, it is difficult to establish ages for brown dwarfs, since their spectra do not always exhibit a known direct indicator of age such as from Li destruction. This results in degeneracies amongst mass, age, and luminosity. However, for all but the youngest objects, brown dwarf radii are largely independent of mass and age, and all are similar to the radius of Jupiter \\citep{cha00}. Thus luminosity scales well with T$_{\\rm eff}^4$, with L dwarfs having surface temperatures in the range 2200--1400 K, while T dwarfs have temperatures down to about 700~K (e.g. Kirkpatrick et al. 2000; Leggett et al. 2001; Leggett et al. 2002; Dahn et al. 2002; Burgasser at al. 2002a; Golimowski et al. 2004; this paper). Obviously, an accurate measurement of the distances to these objects is required to determine their luminosities and temperatures, along with understanding many other issues such as their cooling curves and surface flux redistribution due to atmospheric dust formation.  In earlier work, U.S. Naval Observatory (USNO) optical CCD parallaxes and proper motions were presented for eight late M dwarfs, 17 L dwarfs, and three T dwarfs (Dahn et al. 2002, hereafter D02). Most recently Tinney, Burgasser, \\& Kirkpatrick (2003, hereafter TBK03) presented near--infrared parallaxes and proper motions of 9 T dwarfs. In this paper we present preliminary trigonometric parallaxes and proper motions, obtained at near--infrared wavelengths, of 22 L dwarfs and 18 T dwarfs plus four additional L or T dwarf companion objects in binaries. We feel compelled to present preliminary parallaxes and proper motions now due to the intense interest by the  community in distance determinations to these objects, rather than waiting approximately another two years of observational time baseline before final results would be available for most of the objects. Final results will be presented in later papers. ",
        "conclusions": "In this section we present selected infrared absolute magnitude versus spectral type and infrared color relationships. At the risk of being accused of astronomical chauvinism, we have chosen for the remainder of this paper to discuss only USNO--derived optical and infrared parallaxes (this paper and D02). This provides a self--consistent set of parallax and proper motion determinations using the same telescope and similar observing philosophies and reduction software, with only the detector being different. When we publish completed parallax solutions, it will then be appropriate to combine these results with those of other researchers. Four of the objects, 2MASS~J085035+1057AB, 2MASS~J122554$-$2739AB, 2MASS~J172811+3948AB, and 2MASS~J210115+1736AB  are known binaries. We discuss how the separated spectral types and photometry are determined in \\S 15.4, \\S 15.6, \\S 15.8, and \\S 15.9, respectively.  \\subsection{Absolute Magnitude versus Spectral Type}  In Figure~2 we plot J-band absolute magnitude ($M_J$) versus spectral type. The solid data points are the results from this paper, where we have combined the infrared parallaxes with the infrared photometry and spectral classifications listed in Table~6. The photometric errors have been convolved with the parallax uncertainties to produce the vertical error bars. The horizontal errors are $\\pm$ 0.5 spectral type for those objects with well-determined spectral classification and $\\pm$ 1.0 spectral type for those with less certain classifications. The open data points are from D02 where we have used the parallaxes, infrared photometry, and spectral types published in that paper. For the seven objects in common between D02 and this paper, we plot both the CCD- and infrared-derived absolute magnitudes using the photometry of Table~6. In order to be consistent photometrically, we plot T513$-$46546 (D02) using 2MASS All-Sky PSC photometry transformed to CIT values by the Carpenter (2001) transformations. Several of the objects, in both the infrared and optical parallax groups, are known binaries and we have plotted them in accordance with what is known about their binary natures. These objects are discussed individually in \\S~15 and \\S~16 below. In Figures 3 and 4 we plot the $M_H$ and $M_K$ absolute magnitudes, respectively, versus spectral type with the same considerations as for Figure~2.  Our results for spectral types earlier than about L5 do not provide much new information, since we have only a few early L dwarfs, some of which currently have large error bars. However, the results are consistent with narrow loci in all three diagrams. For objects between L5 and L9 the dispersion is clearly much greater than for earlier objects. The widths are about 1.5 mag, 1.3 mag, and 1.0 mag in J, H, and K, respectively, for the L5--8 objects.  While this could be due to an admixture of objects of different ages, masses, and gravities, we have looked, within a given spectral type, for correlations of $M_J$, $M_H$, and $M_K$ with tangential velocities (see \\S 14.3) as a potential age indicator, but have found none. More likely, the additional width is due to the complicated atmospheric physics expected for late L dwarfs \\citep{bur02b,ste03}, which also can explain why the spread is a function of wavelength. These models also predict significant variability due to rapid evolution or motion of cloud holes, which alone could be responsible for the apparent spread in absolute magnitude. Clearly photometric monitoring will be necessary to fully understand the L--T transition objects.  In all three diagrams the transition from L to T dwarfs is smooth. The excess in luminosity for T1-5 spectral types, previously noted by D02 and TBK03, is clearly substantiated by our enhanced database. While this excess is most evident in the J band, it is also seen in the H and K bands. The possibility that the early T dwarf luminosity excess is caused by contamination due to binary systems \\citep{bur01} now seems unlikely due to the sheer number of objects which participate in this hump. TBK03 point out that the amplitude of the hump, like the spread at late L, is also unlikely to be explained by an age selection effect \\citep{tsu03}. T dwarfs between T6 and T8 once again form a rather tight locus terminating at $M_{J,H,K}$~$\\approx$~16.6.  The additional data for T dwarfs allow a somewhat clearer picture of the L--T transition region in these diagrams. While there is no self--evident reason to believe that $M_J$, $M_H$, or $M_K$ should map linearly with spectral type (TBK03), we note that inspection of Figures 2 through 4 shows that the late T dwarfs (T6--T8) are on a rough extension of the absolute magnitude versus spectral type relation of the the early L dwarfs (L0--L5). Relative to a fiducial line drawn between the early Ls and late Ts, in J--band the L--T transition objects show a luminosity deficit of about 1.5 mag at L6--L8 and a luminosity excess of about 1.5 mag at T1--T5. In H--band the L6--L8 deficit has shrunk to about 0.5 mag and disappears at K--band, while in H-- and K--bands the excess for T1--T5 objects remains at 1.0--1.5 mag.  We note that the contrast between the large spread of absolute magnitudes at late L versus the apparently narrow locus of early T absolute magnitudes may not be significant. The narrow T dwarf locus may yet prove to be an artifact of small number statistics. Also, spectral typing may just map out an equivalent diversity of physics over a smaller range of spectral types at late L than at early T.  \\subsection{Absolute Magnitude versus Infrared Colors}  In Figure~5 we plot $M_J$ versus $J-H$ color. The solid data points are the results from this paper, where we have combined the infrared parallaxes with the infrared photometry and spectral classifications listed in Table~6. The photometric errors have been convolved with the parallax uncertainties to produce the vertical error bars, while the horizontal errors are from Table~6. The open data points are those from D02, where we have used the parallaxes, infrared photometry, and spectral types published in that paper. We treat the seven objects in common between D02 and this paper and the binary systems as described in \\S 14.1. In Figure~6 we plot $M_J$ versus $J-K$ color, while in Figures 7 and 8 we plot $M_K$ versus $J-H$ and $J-K$, respectively.  These figures show the well-known color trends for L dwarfs ranging from $J-H$~$\\approx$~0.7, $J-K$~$\\approx$~1.1 for L0 ($M_J$~$\\approx$11.0, $M_K$~$\\approx$~10.0) to $J-H$~$\\approx$~1.2, $J-K$~$\\approx$~2.1 for L8 ($M_J$~$\\approx$~15.0, $M_K$~$\\approx$~13.5). Late T dwarfs (T5 to T8) all have roughly $J-H$~$\\approx$~0, $J-K$~$\\approx$~0, while early T dwarfs (T0.5 to T4.5) have a wide range of colors transitioning between late L and late T colors. Several objects are now placed in the transition region between the loci of the L and T dwarfs, which is best shown in the $M_K$ versus $J-K$ diagram (Figure~8). Unfortunately, the locations of many of the transition objects and early T dwarfs are poorly known at this time due to uncertain distances. We will defer comparing evolutionary models with observations until we obtain final parallaxes and further USNO-CIT photometry. However, we note that the apparent brightening in $M_J$ across the L--T transition is consistent with the predictions of the cloud hole model of Burgasser et al. (2002b), while the fact that the brightening is an apparent trend argues against the hypothesis of Tsuji \\& Nakajima (2003) that it is simply an age effect.  \\subsection{Kinematics}  For stars in the solar vicinity, motion with respect to the Sun is an indicator of age, since older stars will have had time to be perturbed preferentially to different orbits by interaction with the Galactic disk. Because T dwarfs are thought to be the cooler and older analogs of at least some L dwarfs, it might be expected that T dwarfs will have a larger mean velocity, with respect to the Sun, than L dwarfs. The measured tangential velocities (V$_{tan}$) with respect to the Sun measured primarily for L dwarfs in D02 can be combined with those for L and T dwarfs in this paper (Table~2) to have sufficient objects in order to compare the velocity distributions. For the seven objects in common between the two papers we use the weighted mean values of V$_{tan}$. We remove three objects from consideration with exceptionally large V$_{tan}$ uncertainties, having both $\\sigma$(V$_{tan}$)~$>$~10~km~s$^{-1}$ and V$_{tan}$/$\\sigma$(V$_{tan}$)~$<$~3 (two L dwarfs: 2MASS J143535$-$0043 [31.7$\\pm$13.3 km s$^{-1}$] and 2MASS J095105+3558 [55.8$\\pm$32.7 km s$^{-1}$] and one T dwarf: SDSS~J083717$-$0000 [24.3$\\pm$11.8 km s$^{-1}$]; all from this paper). This leaves 33 L dwarfs and 17 T dwarfs to make the velocity comparison. The unweighted average values of tangential velocity for L and T dwarfs, respectively, is 30.0~$\\pm$~3.6 and 43.0~$\\pm$~4.8 km s$^{-1}$. The median values are, for L and T dwarfs respectively, 24.5 and 39.0 km s$^{-1}$. As D02 have discussed, the velocity of the L dwarfs is consistent with velocities for old disk M and dM stars.  The distributions of V$_{tan}$ are shown in Figure~9 for L dwarfs in the top panel and for T dwarfs in the bottom panel. Based on the Kolmogorov--Smirnov test, the null hypothesis that the two distributions are indistinguishable can only be rejected at the 73\\% level. The main difference is that the T dwarfs have no examples with V$_{tan}$~$\\le$~20 km s$^{-1}$, whereas the L dwarfs have 11 of 33 (33\\%) with V$_{tan}$~$\\le$~20 km s$^{-1}$. While there are fewer objects in the T dwarf subset, if the distributions were the same, 5.7 T dwarfs with V$_{tan}$~$\\le$~20 km s$^{-1}$ would be expected.    In this paper we have presented preliminary parallaxes and proper motions for 22 L dwarfs and 18 T dwarfs derived over time baselines of only $\\Delta$t~$\\approx$~1.3 or $\\Delta$t~$\\approx$~2.0 years. The resultant mean parallax uncertainties of 4.86 and 3.85 mas, respectively, which will be greatly improved by ongoing further observations, are nonetheless of sufficient quality to provide some significant new results for T dwarfs. We list here a summary of the more important conclusions we reach from our work.  A. The luminosity excess `hump' for early to mid T dwarfs in the absolute magnitude versus spectral type diagram is clearly confirmed. While seen most strongly at J--band, it is also evident in the H-- and K--bands. The possibility that the hump is due to a selection effect of binaries is likely ruled out by the large number of objects participating in the hump.  B. L5--L8 dwarfs have a significantly larger spread in the absolute magnitude versus spectral type diagrams than do earlier L dwarfs.  C. Late T dwarfs have a narrow locus in absolute magnitude versus spectral type diagrams, similar to the early L dwarfs. Relative to a straight line connecting the earlier and later objects, the L--T transition objects show systematic trends.  In J--band the late L dwarfs show a luminosity deficit and the early T dwarfs a luminosity excess. The late L dwarf luminosity deficit is less in H--band and is gone in K band, while the early T dwarf luminosity excess amplitude is somewhat less in H--band and K--band than at J--band.  D. The absolute magnitude behavior across the L--T transition described above exemplifies the critical role of condensate cloud evolution at these temperatures based on the most recent spectral models.  E. Using newly derived bolometric corrections for L and T dwarfs by Golimowski et al. (2004) we derive luminosities and T$_{\\rm eff}$ for L and T dwarfs with USNO--derived parallaxes either from this paper or D02.  F. 2MASS J041519$-$0935 is found to be the least luminous [log(L/L$_{\\odot}$)~=~$-$5.58] and therefore coldest (T$_{\\rm eff}$~$\\approx$~760 K ) brown dwarf yet found.  G. We find a broader distribution of L dwarf tangential velocities compared with that of the T dwarfs. While essentially the same between 20 to 60 km s$^{-1}$, the T  dwarfs do not have a low velocity population as do the L dwarfs. This is consistent with T dwarfs being, in general, older than L dwarfs."
    },
    "0402/astro-ph0402044_arXiv.txt": {
        "abstract": "We present a semi--analytic treatment of galactic winds within high resolution, large scale cosmological N--body simulations of a $\\Lambda$CDM Universe. The evolution of winds is investigated by following the expansion of supernova driven superbubbles around the several hundred thousand galaxies that form in an approximately spherical region of space with diameter $52 h^{-1}$ Mpc and mean density close to the mean density of the Universe. We focus our attention on the impact of winds on the diffuse intergalactic medium. Initial conditions for mass loss at the base of winds are taken from \\citet{shu}. Results are presented for the volume filling factor and the mass fraction of the IGM affected by winds and their dependence on the model parameters is carefully investigated. The mass loading efficiency of bubbles is a key factor to determine the evolution of winds and their global impact on the IGM: the higher the mass loading, the later the IGM is enriched with metals. Galaxies with $10^{9} < M_{\\star} < 10^{10}$ M$_{\\sun}$ are responsible for most of the metals ejected into the IGM at $z=3$, while galaxies with $M_{\\star} < 10^{9}$ M$_{\\sun}$ give a non negligible contribution only at higher redshifts, when larger galaxies have not yet assembled. We find a higher mean IGM metallicity than \\lya\\ forest observations suggest and we argue that the discrepancy may be explained by the high temperatures of a large fraction of the metals in winds, which may not leave detectable imprints in absorption in the Ly$\\alpha$ forest. ",
        "introduction": "Powerful outflows from star--forming galaxies have been detected throughout the history of the universe (\\citealt{heckman90}, \\citealt{hlsa}, \\citealt{adelberger}), providing, perhaps, the mechanism to transport metals from the interstellar medium (ISM) of galaxies to the low density intergalactic medium (IGM). This could at least partially explain the widespread level of chemical enrichment observed in the spectra of quasars (\\citealt{cowie}, \\citealt{joop}, \\citealt{ell}, \\citealt{simcoe}).  The energy necessary to power outflows on galactic scales is supplied by supernova explosions and winds from young massive stars in OB associations. Any episode of star formation may create a superbubble in the ISM and, if the rate of energy input is large enough, the superbubble can blow out of the ISM and create a wind. In local starbursts (\\citealt{phil}, \\citealt{cecil2}, \\citealt{walter}, \\citealt{sugai}), winds have been observed to extend to at least 10 kpc from their host galaxies. \\citet{ses} claim that winds can reach even larger distances, but are unobservable because of the low emissivity of the outflowing gas.  At present, it is difficult to predict which galaxies are responsible for seeding the IGM with metals or to establish the effects that supernova--driven blastwaves have on the galaxy formation process. While gravity does not influence the evolution of superbubbles in the ISM, it is crucial for determining the long term fate of winds. Since winds from dwarf galaxies form in shallower potential wells, they are the most likely to be able to disperse their metal content into the IGM. On the other hand, \\citet{ses} suggest that most of the energy from winds resides in a hot ($T\\sim 10^7$ K) low density component that can escape the galaxies even when the bulk of the outflowing mass is retained. \\citet{mac} demonstrate that metals are easily accelerated to velocities larger than the escape velocity, implying that a galaxy can lose a high fraction of its metals even with a relatively low mass ejection efficiency. Although winds may occur more frequently in dwarf galaxies, the metals ejected by massive galaxies may dominate the total budget. It is therefore not a trivial problem to assess which galaxies have been responsible for the pollution of the IGM and when the enrichment occurred.  Several groups have applied simple phenomenological prescriptions to simulations in order to investigate the effects of winds on the IGM and some important results have emerged. For example, \\citet{mfr} find that pregalactic outflows are an efficient mechanism for distributing the metals produced in stars over large cosmological volumes, prior to the reionisation epoch. \\citet{aguirre} argue that radiation pressure ejection or winds from relatively large galaxies at lower redshifts can account for the observed metallicity of the IGM and the intracluster medium. \\citet{tom}  and \\citet{croft} demonstrate that cavities evacuated by winds in the outskirts of galaxies may leave characteristic signatures in the Ly$\\alpha$ forest. In contrast, \\citet{tv} find that winds have little effect on the statistics of H I absorption lines and produce C IV absorption lines in reasonable agreement with observations.  The significance of galactic winds for the evolution of the IGM is still not fully established, however. Both hydrodynamic and semi--analytic simulations use phenomenological prescriptions for the physics of galactic winds and new parameters have to be introduced to account for the uncertainties that derive from a still uncomplete observational picture. In particular, no well founded relation is available to link the properties of the ISM and the morphology of its host galaxy to the structure and evolution of the outflows. Because of insufficient resolution and incomplete physics, numerical results often disagree with each other and the effects of winds on the Ly$\\alpha$ forest remain controversial, leaving the way open for further studies.  In this paper, we present a new implementation of the physics of galactic winds within the semi--analytic galaxy formation model of \\citet{semi}, and we apply it to a set of high resolution N--body simulations of structure formation in a $\\Lambda$CDM universe (\\citealt{felix}, \\citealt{bene}). By using a high resolution simulation of a spherical region of diameter 52 $h^{-1}$ Mpc, we investigate the long term evolution of winds and their effects on a typical region of the IGM. We solve the equation of motion for a spherical astrophysical blastwave to follow the evolution of winds after they escape the visible regions of galaxies. Our phenomenological model for winds uses the initial conditions proposed by \\citet{shu}, which parameterise the mass loss and the initial velocity of winds as a function of the star formation rate of the galaxy. Here we follow the evolution of galactic winds throughout most of the history of the universe and we outline their impact on the IGM by estimating the fraction of volume and mass of the IGM which they affect as a function of time and model parameters.  This paper is organised as follows: in Section \\ref{due} we present our set of high resolution N-body simulations and the semi--analytic prescriptions we adopt to model the physics of galactic winds; in Section \\ref{tre} we outline some global properties of winds as a function of the model parameters; in Sections \\ref{quattro} and \\ref{budget} we present the results for the volume filling factor and the fraction of mass affected by winds and in Section \\ref{ejection} our findings for the ejection of metals into the IGM; finally, in Section \\ref{mares} we discuss the dependence of our results on the numerical resolution of the simulations and we draw our conclusions in Section \\ref{cinque}. ",
        "conclusions": "\\label{cinque}  We have presented semi--analytic simulations of galaxy formation in a cosmological context, which include the physics of galactic winds. The semi-analytic prescriptions are applied to high resolution N--body simulations of a typical ``field'' region of the Universe.  The results of our model can be quite accurately interpreted as a consequence of the mass loading efficiency of winds. The mass accumulated in bubbles is directly linked to the amount of mass entrained from the ambient medium, set by the parameter $\\varepsilon$, and the ultimate fate of winds is strongly dependent on this swept--up mass. Bubbles that load little mass from the surrounding medium can escape the gravitational potential well of their host haloes more efficiently at every redshift. These bubbles are mostly composed of metal rich supernova ejecta and shocked ISM and need to spend little of their energy to accelerate the accreted gas. Since most of the energy injected by the starburst is available to power the expansion of the bubble, these winds have the highest probability to escape the gravitational attraction of haloes and expand into the IGM. The formation of highly mass loaded winds is instead suppressed in all kinds of galaxies, although the suppression is particularly strong in galaxies with $M_{\\star} \\lesssim 10^{9}$ M$_{\\sun}$. This is because the energy provided by star formation is not sufficient to overcome the ram pressure of the infalling material which adds to the gravitational pull of the galaxy.  Our estimates of the volume filling factor of winds (Section \\ref{quattro}) and of the fraction of IGM mass affected by winds (Section \\ref{budget}) suggest that galactic outflows are unlikely to significantly modify the properties of the Ly$\\alpha$ forest. No obvious correlation is found between $f_m$ and $f_v$. The volume filling factor is clearly dependent on the mass loading efficiency of bubbles, with low values of $f_v$ associated with highly mass loaded bubbles and viceversa. The fraction of IGM mass affected by winds is usually comparable to the volume filling factor. Only in models with high mass loading efficiency we find that $f_m > f_v$, which implies that the actual fraction of intergalactic mass affected by outflows may be large even when the winds physically fill a small region of space. This is a consequence of the clustering of matter on large scales and of the fact that galaxies form in high density regions, where their winds can sweep up a larger amount of material than they would if they were expanding inside a low density region.  The efficiency of winds in seeding the IGM with metals is investigated in section \\ref{ejection}. Galaxies with $M_{\\star} \\lesssim 10^{9}$ M$_{\\sun}$ play a role in the chemical enrichment of the IGM only at very high redshifts, when larger objects have not yet assembled. At $z=3$ most of the metals are ejected by galaxies with $10^9 \\lesssim M_{\\star} \\lesssim 10^{10}$ M$_{\\sun}$, while galaxies with $M_{\\star} \\gtrsim 10^{10}$ M$_{\\sun}$ contribute only about 10\\%--20\\% of the ejected metals. The result that metals are mostly ejected by relatively small galaxies qualitatively agrees with the predictions of e.g. \\citet{tom}, \\citet{tv}, \\citet{violence}.  Our estimates of the mean metallicity of the IGM are significantly higher than the observed values at $z\\sim 1$ to $z\\sim 5$ and we have argued that metals in the IGM might not be observable in absorption in the spectra of quasars because of the high temperatures of winds. In a forthcoming paper we will discuss the possibility of finding observable signatures of cooled wind shells in the Ly$\\alpha$ forest."
    },
    "0402/astro-ph0402334_arXiv.txt": {
        "abstract": "{Two alternative scenarios concerning the origin and evolution of extremely metal-poor halo stars are investigated. The first one assumes that the stars have been completely metal-free initially and produced observed carbon and nitrogen overabundances during the peculiar core helium flash typical for low-mass Population~III stars. The second scenario assumes that the initial composition resulted from a mixture of primordial material with ejecta from a single primordial supernovae. Both scenarios are shown to have problems in reproducing C, N, and O abundances simultaneously, and both disagree with observed $\\Iso{12}{C}/\\Iso{13}{C}$-ratios, though in different directions. We concentrate on the most iron-poor, carbon-rich object of this class, \\hechr, and conclude, that the second scenario presently offers the more promising approach to understand these objects, in particular because evolutionary tracks match observations very well. ",
        "introduction": "The extremely (or ultra) metal-poor stars (UMPS) of the galactic halo are believed to be the closest links of the Galaxy to the first generation of stars, to Population~III. We therefore hope to learn about the first epoch of star formation and the end of the Dark Ages because they either are members of Population~III themselves or because they carry the immediate imprint of massive Pop.~III stars or primordial Supernovae. They have received considerable attention in the recent past because of the fact that they are at the crossroads of stellar evolution, star formation, galactic chemical evolution and cosmology, notably here the recent CMB results and the question of reionization by Pop.~III stars.  The discovery of \\hechr, with a record low abundance of heavy elements of $\\mathrm{[Fe/H]} = -5.3$, which is about a factor of 10 below the previously known lowest value, raised our interest to link this star to Pop.~III. Remarkably, the total ``metallicity'' of \\hechr\\ is by far not metal-poor due to a carbon and nitrogen overabundance of $\\mathrm{[C/Fe]} = 4.0$ and $\\mathrm{[N/Fe]} = 2.3$, putting it into a large subgroup of UMPS with similar peculiar composition.  From the point of view of stellar evolution theory metal-free stars are interesting in themselves because of some aspects of their structure and evolution which differ drastically from that of ordinary Pop.~II or I stars. One of these pecularities is the fact that during the core helium flash mixing between the helium and hydrogen shell and the convective envelope can take place, resulting in a carbon- and nitrogen-rich envelope and a second red giant branch phase. Therefore, it is a plausible assumption that the UMPS are true Pop.~III stars and that the carbon-rich subgroup consists of stars that experienced the first, peculiar helium flash. We want to emphasize that in this paper we are concerned only with the abundances of the CNO-elements, since heavier elements are unaffected by the nuclear processes in low-mass stars.  In our previous papers, we have therefore investigated the scenario mentioned above to explain the observed chemical composition of carbon-rich UMPS qualitatively and thus to link them to Population~III. The main problem one faces is that the flash-induced mixing appears to result always, independent of the details and assumptions of the calculations, in the same amount of carbon and nitrogen, such that for $\\mathrm{[Fe/H]} \\lesssim -3$ the predicted carbon overabundance is $\\mathrm{[C/Fe ]} \\gtrsim 3$, which is at the upper limit of observed values. Since the C overabundance of \\hechr is as high as the value found in our previous calculations, it appeared to be particularly worthwhile to apply our approach to this star. Similar calculations have recently been performed independently by \\cite{pclp:2004}.\\footnote{Our results were presented at the First~Stars~II meeting in May 2003; \\\\ see {\\tt http://www.astro.psu.edu/users/tabel/II/presentations/weiss.pdf}.}  We therefore present a model for \\hechr, based on the flash-induced mixing (FIM) in Sect.~2, repeating the basic features of this event. In Sect.~3, we will then show calculations for an alternative scenario, which assumes that this star (and other UMPS) are formed directly from the ejecta of Pop.~III supernovae, i.e.\\ that UMPS are the immediate successors of true, massive Pop.~III objects. This scenario is also applied to other extremely metal-poor halo stars. Sect.~4 finally summarizes our conclusions. ",
        "conclusions": "In earlier papers (Papers II and III) we have applied our scenario for the evolution of initially metal-free Pop.~III stars with additional surface pollution to some known objects of the galactic halo, investigating the possibility that the observed severe carbon and nitrogen enhancements are due to internal production and mixing in the course of the first core helium flash. These models indicated that both the abudnances of C and N in the models are too high, and that the carbon isotope ratio is too close to equilibrium values. In addition, we noticed at that time that a statistically representative sample of UMPS is needed to verify that the number of carbon-enhanced objects is in agreement with the predictions from the models, which would be post-flash, and therefore short-lived compared to lower RGB stars. The detection of \\hechr\\ with the lowest iron abundance of $\\logab{Fe}{H} = -5.3$ and highest carbon enhancement of $\\logab{C/Fe} = 4.0$, allowed us to investigate this scenario more closely. We thus computed specific models using realistic SN yields from the literature.  The low iron content of \\hechr\\ can only be achieved by adding tiny amounts of SN-ejecta matter or to impose rigid mass cuts for the SN explosion. We find that, independent of the particular choice of polluting matter, the final carbon and nitrogen abundances, which result from the core helium flash and the subsequent mixing, exceed the observed abundances by orders of magnitude, similar to the case of the less extreme cases we modeled in our previous papers. It appears that both in nature, and in our models, the {\\em amount} of overabundant C and N is rather constant, such that with decreasing Fe-abundance the relative overabundance is increasing. This effect can actually been seen in Fig.~2 of \\cite{rbs:99} by looking at the upper envelope of the $\\logab{C}{Fe}$ vs.\\ $\\logab{Fe}{H}$ distribution. However, the models appear to produce 10--100 times too much carbon and nitrogen. Additionally, the carbon isotope ratio in \\hechr\\ is definitely far above equilibrium values, which are always obtained in our simulations. These results, with the exception of the oxygen abundance, are widely independent on whether we use solar or early SNe material for the pollution. Therefore, \\hechr\\ contradicts most strongly our flash-induced mixing models.  We then investigated an alternative scenario, assuming that \\hechr\\ and other extremely metal-poor stars are representing a kind of ``early Pop.~II'' or Pop.~II.5 class of objects. Their homogeneous initial composition is of low, but finite metallicity. However, contrary to standard Pop.~II stars, the material still carries the imprint of one or few individual SNe of Pop.~III. We restricted ourselves to one particular SN model, since in terms of CNO-elements the various models available (mass, explosion energy, author) do not vary drastically. The choice was made according to inferences based on reproducing the heavy elements in \\hechr.  We find that, after mass cut and dilution with pristine interstellar material have been fixed, the stars are carbon and oxygen rich already on the main sequence, produce large amounts of nitrogen in CN-cycling, and expose them as a consequence of standard first dredge-up. The evolution is quite standard, as for Pop.~II models. C and N abundances agree naturally very well with the observations, but carbon isotope ratios are in this scenario definitely {\\em higher} than for most stars under consideration (with the exception of \\hechr). Therefore, for $\\Iso{12}{C}/\\Iso{13}{C}$ we face the opposite problem of the one we have for the FIM-models. Oxygen is always, due to the composition of SN ejecta, enriched. It would therefore be necessary to obtain results for oxygen for UMP stars\\footnote{It appears that indeed several of them have $\\logab{O}{Fe} \\approx 2$, similar to \\hechr\\ (V.~Hill, private communication).} to put further constraints on the various possibilities for the nature of the UMPS. In case of the objects we tried to model and for which we had oxygen abundances, it appears that we can roughly reproduce the observation. However, in case of \\hechr\\ the SN yields predict too much oxygen. This problem can also be noticed in the models by \\cite{lcb:2003}, while it seems to be less severe in \\cite{umno:2002} using a less massive SN progenitor with only $0.3\\cdot 10^{51}\\,\\mathrm{erg}$ explosion energy. The composition of one of the objects we investigated could be matched satisfactorily.  The observed stars lie very nicely on the tracks for our Pop.~II.5 scenario, although the errors are too large to allow excluding the FIM-possibility on that ground. It is possible to identify different evolutionary stages for the observed objects. Generally, the later the evolutionary phase, the higher the N abundance and the lower the $\\Iso{12}{C}/\\Iso{13}{C}$ ratio, reducing some of the problems. If the star would be on the AGB, a better agreement is possible, in particular, if the 3rd dredge-up happens during the thermal pulses, because in this case, the carbon isotope ratio would be reduced strongly. However, we have not followed the evolution of the models this far. Again, statistically significant samples of UMPs would be necessary to further look into this question.  In spite of the remaining problems, we presently favor the idea that the observed extremely metal-poor stars of the galactic halo are stars formed directly from the ejecta of one or few Pop.~III SNe of intermediate mass ($15-60\\,\\msun$), which are diluted with metal-free primordial gas. Overall, the agreement between model and observations appears to be better, and there is still a large uncertainty concerning the SN yield composition. Also, whether one or two or a few SNe have contributed to the initial composition of an UMP \\citep[see the discussion in][]{lcb:2003}, allows further fine-tuning of models. Nevertheless, solid statistical samples are clearly needed for further progress. Finally, we point out that all SN models favored, indicate progenitor masses of $\\approx 20 - 60\\,\\msun$. Currently, the primordial star formation scenario is favouring much higher initial masses for Pop~III stars. This question, too, remains to be cleared. The extremely metal-poor stars with their particular composition, may guide us in this."
    },
    "0402/astro-ph0402428_arXiv.txt": {
        "abstract": "We discuss the derivation of the analytic properties of the cross-power spectrum estimator from multi-detector CMB anisotropy maps. The method is computationally convenient and it provides unbiased estimates under very broad assumptions. We also propose a new procedure for testing for the presence of residual bias due to inappropriate noise subtraction in pseudo-$C_{\\ell}$ estimates. We derive the analytic behavior of this procedure under the null hypothesis, and use Monte Carlo simulations to investigate its efficiency properties, which appear very promising. For instance, for full sky maps with isotropic white noise, the test is able to identify an error of 1\\% on the noise amplitude estimate. ",
        "introduction": "The cosmic microwave background (CMB) provides one of the most powerful ways of investigating the physics of the early Universe. The main CMB observable is the angular power spectrum of temperature anisotropy, which encodes a large amount of cosmological information. In the last decade, important advances in the measurement of the CMB angular power spectrum took place; this resulted in relevant progress in our understanding of physical cosmology. CMB temperature anisotropies were first detected by the COBE satellite in 1992 \\cite{smoot92}. This discovery fuelled a period of intensive experimental activity, focused on measuring the CMB power spectrum on a large range of angular scales. A major breakthrough was made in the past few years, when the MAXIMA \\cite{hanany00} and BOOMERanG \\cite{debe00} balloon-borne experiments independently produced the first high-resolution maps of the CMB, allowing a clear measurement of a peak in the power spectrum, as expected from theoretical models and previously detected by the ground based experiment TOCO \\cite{miller99}. Since then, many other experiments have confirmed and improved on these results: DASI\\cite{dasi02}, BOOMERanG-B98 \\cite{nett02,debe02,ruhl03}, BOOMERanG-B03 \\cite{masi05, jones05, fp05, montroy05, cmt05}, VSA \\cite{vsa03}, Archeops \\cite{benoit03}, CBI \\cite{cbi02}, ACBAR\\cite{acbar02}, BEAST \\cite{beast03}. Most notably, the NASA satellite mission WMAP, whose first year data were released in February 2003 (\\cite{wmap03} and references therein), provided the first high-resolution, full sky, multi-frequency CMB maps, and a determination of the angular power spectrum with unprecedented accuracy on a large range of angular scales. Much larger and more accurate data sets are expected in the years to come from ESA's Planck satellite.  In this paper, we shall concentrate on extracting the CMB power spectrum from full sky maps with foregrounds removed. We shall focus mainly on techniques for dealing with noise subtraction. In principle, and for Gaussian maps, noise subtraction can be performed by implementing maximum likelihood estimates. It is well known, through \\cite{bjk98,borrill99}, that maximum likelihood estimates require for their implementations a number of operations that scales as $N_{pix}^{3},$ $N_{pix} $ denoting the number of pixels in the map. For current experiments, $N_{pix}$ ranges from several hundred thousands to a few millions, and thus the implementation of these procedures is beyond computer power for the near future. Many different methods have been proposed for producing computationally feasible estimates; here we just mention a few of them, and we refer the reader to \\cite{efst03} for a more complete discussion on their merits. Some authors have introduced special assumptions on the noise properties and symmetry of the sky coverage, to make likelihood estimates feasible; see, for instance, \\cite{oh99,wgh01,wh01,ch01}. Reference~\\cite{sza01} adopted an entirely different strategy, extracting the power spectrum from the 2-point correlation function of the map. Others have used estimators based on pseudo-$C_l$ statistics and Monte Carlo techniques \\cite{efh01,balbi02}, or based on Gabor transforms \\cite{fh02}. For multi-detector experiments, an elegant method, based on spectral matching to estimate jointly the angular power spectrum of the signal and of the noise, was proposed in \\cite{dela}. Pseudo-$C_{\\ell}$ estimators were adopted by the WMAP team \\cite{wmapps}, which used the cross-power spectrum estimator and discussed the best combination of the cross-power spectrum obtained from single couples of receivers.  Our purpose in this paper is to derive some analytic results on the cross-power spectrum estimator, to perform a comparison with standard pseudo-$C_{\\ell}$ estimators, and to propose some testing procedures on the assumption that any noise bias has been properly removed, which is clearly a crucial step in any estimation approach. We shall also present some Monte Carlo evidence on the performance of the methods that we advocate. The plan of this paper is as follows. In Section \\ref{sec:ps} we derive the analytic properties for the cross-power spectrum estimator and we compare them with equivalent results on standard pseudo-$C_{\\ell}$ estimators. In Section \\ref{sec:haus} we propose a procedure (the Hausman test) for verifying appropriate noise subtraction in pseudo-$C_{\\ell}$ estimators, and we derive its analytic properties. In Section \\ref{sec:mc} we validate our results by using Monte Carlo simulations, which are also used to test the power of our procedure in the presence of noise which has not been completely removed. In Section \\ref{sec:con} we review our results and discuss directions for future research. ",
        "conclusions": "\\label{sec:con} We have discussed the analytic properties of the cross-power spectrum as an estimator of the angular power spectrum of the CMB anisotropies. The method is computationally convenient for very large data sets as those provided by WMAP or Planck and it provides unbiased estimates under very broad assumptions (basically, that noise is uncorrelated along different channels). It thus provides a robust alternative, where noise estimation and subtraction are not required. We also propose a new procedure for testing for the presence of residual bias due to inappropriate noise subtraction in pseudo-$C_{\\ell}$ estimates (the Hausman test). The test compares the auto- and cross-power spectrum estimators under the null hypothesis. In the case of failure, the more robust cross-power spectrum should be preferred, while in the case of success both estimators could be used, and the choice should result from a trade-off between efficiency and robustness. We derive the analytic behaviour of this procedure under the null hypothesis, and use Monte Carlo simulations to investigate its power properties, which appear extremely promising. We leave for future research some further improvements of this approach, in particular, the use of bootstrap/resampling methods to make even the determination of confidence intervals independent from noise estimation. Finally, the optimal combination of polarization power spectra and the application of the Hausman test to the Planck satellite are currently under investigation.     \\ack The authors would like to thank G. De Gasperis for providing realistic temperature and polarization simulated maps and F.K. Hansen, P. Cabella, G. De Troia, F. Piacentini and K. Ganga for useful discussions. We acknowledge the use of the MASTER, HEALPix, CMBFAST and FFTW packages. Research supported by MURST, ASI, PNRA.          \\appendix"
    },
    "0402/astro-ph0402102_arXiv.txt": {
        "abstract": "{We present XMM-Newton MOS imaging and PN timing data of the faint millisecond pulsars \\object{PSR J0751+1807} and \\object{PSR J1012+5307}.  We find 46 sources in the MOS field of view of \\object{PSR J0751+1807} searching down to an unabsorbed flux limit of 3 $\\times 10^{-15}$ ergs\\ cm$^{-2}$ s$^{-1}$ (0.2-10.0 keV).  We present, for the first time, the X-ray spectra of these two faint millisecond pulsars.  We find that a power law model best fits the spectrum of \\object{PSR J0751+1807}, $\\Gamma$=1.59$\\pm$0.20, with an unabsorbed flux of 4.4$\\times 10^{-14} {\\rm ergs\\ cm}^{-2} {\\rm s}^{-1}$ (0.2-10.0 keV). A power law is also a good description of the spectrum of \\object{PSR J1012+5307}, $\\Gamma$=1.78$\\pm$0.36, with an unabsorbed flux of 1.2$\\times 10^{-13} {\\rm ergs\\ cm}^{-2} {\\rm s}^{-1}$ (0.2-10.0 keV).  However, a blackbody model can not be excluded as the best fit to this data.  We present some evidence to suggest that both of these millisecond pulsars show pulsations in this X-ray band.  We find some evidence for a single broad X-ray pulse for \\object{PSR J0751+1807} and we discuss the possibility that there are two pulses per spin period for \\object{PSR J1012+5307}. ",
        "introduction": "X-ray emission from millisecond pulsars (MSPs) is thought to be from: charged relativistic particles accelerated in the pulsar magnetosphere (non-thermal emission indicated by a hard power-law spectrum and sharp pulsations); and/or thermal emission from hot polar caps; and/or emission from a pulsar driven synchrotron nebula; or interaction of relativistic pulsar winds with either a wind from a close companion star or the companion star itself \\cite[see e.g.][for a more thorough review]{beck02}.  Until now, X-ray pulsations and spectra have often not been observable for faint MSPs e.g. \\object{PSR J0751+1807} \\citep{beck96} and \\object{PSR J1012+5307} \\citep{halp97}.  Thus it has been difficult to discriminate between competing neutron star models.  However, taking advantage of the large collecting area of {\\em XMM-Newton} \\citep{jans01}, it is becoming possible to observe not only the X-ray spectra but also put a limit on the presence of X-ray pulsations of these faint MSPs.  We have observed two faint MSPs with {\\em XMM-Newton}.  \\object{PSR J0751+1807} was first detected in an EGRET source error box, in September 1993 \\citep{lund93}, using the radio telescope at Arecibo. \\cite{lund95} combined the mass function, eccentricity, orbital size and age of the pulsar, determined from radio data, to predict the expected type of companion star to the millisecond pulsar.  They proposed that the secondary star is a helium white dwarf, with a mass between 0.12-0.6M$_\\odot$, in a 6.3 hour orbit with the pulsar.  They determined a 3.49 ms pulse period, but from the period derivatives, the spin down energy indicated that the pulsar is not the source of the $\\gamma$-rays that were originally detected by EGRET.  \\object{PSR J0751+1807} was subsequently detected in the soft X-ray domain by \\cite{beck96}, using the ROSAT PSPC.  However, there were too few counts detected to fit a spectrum or detect pulsations. Using the HI survey of \\cite{star92}, they deduced an interstellar absorption of 4$\\times$ 10${\\rm^{20}}$ cm${\\rm^{-2}}$.  Then using their estimated counts and assuming a powerlaw spectrum of $dN/dE \\propto E^{-2.5}$, they determined an unabsorbed flux of 1 $\\times$ 10${\\rm^{-13}}$ erg cm${\\rm^{-2}}$ s${\\rm^{-1}}$ (0.1-2.4 keV).  The corresponding X-ray luminosity was then calculated to be L$_x$ = 4.7 $\\times$ 10${\\rm^{31}}$ erg s${\\rm^{-1}}$, for a distance of 2 kpc.  The distance they used was calculated using the radio dispersion measure and the model of \\cite{tayl93} for the galactic distribution of free electrons.  \\object{PSR J1012+5307}, a 5.26 msec pulsar, was discovered by \\cite{nica95} during a survey for short period pulsars conducted with the 76-m Lovell radio telescope.  From the dispersion measure they found a distance of 0.52 kpc.  \\cite{call98}, and references therein, confirmed a white dwarf secondary of mass 0.16$\\pm$0.02M$_\\odot$ in a 14.5 hour circular orbit with the pulsar. \\cite{halp96} associated a faint (L$_x \\approx $2.5$\\times$ 10$^{30}$ ergs s${\\rm^{-1}}$, 0.1-2.4 keV) X-ray source detected with the ROSAT PSPC (80$\\pm$24 photons), with the radio MSP \\object{PSR J1012+5307}. However, the number of photons was insufficient to determine a spectrum or any pulsations.  In this work we present the X-ray spectrum of both \\object{PSR J0751+1807} and \\object{PSR J1012+5307} for the first time.  We also present some evidence for X-ray pulsation from both of these faint pulsars and we compare their nature with other millisecond pulsars, e.g. \\object{PSR J0437-4715} \\citep[][]{beck93,zavl98,zavl02}. ",
        "conclusions": "We have presented XMM-Newton data of the faint millisecond pulsars \\object{PSR J0751+1807} and \\object{PSR J1012+5307}.  Both of these pulsars have a reasonably large magnetic field at the light cylinder radius, which could indicate that both of these MSPs should show pulsations in X-rays.  We present for the first time the X-ray spectra of these two faint millisecond pulsars.  We find that a power law model best fits the spectrum of \\object{PSR J0751+1807}, $\\Gamma$=1.59$\\pm$0.20, with an unabsorbed flux of 4.4 $\\times 10^{-14} {\\rm ergs\\ cm}^{-2} {\\rm s}^{-1}$ (0.2-10.0 keV). A power law is also a good description of the spectrum of \\object{PSR J1012+5307}, $\\Gamma$=1.78$\\pm$0.36, with an unabsorbed flux of 1.2 $\\times 10^{-13} {\\rm ergs\\ cm}^{-2} {\\rm s}^{-1}$ (0.2-10.0 keV).  However, a blackbody model can not be excluded as the best fit to the data.  We have also shown evidence to suggest that both of these MSPs may show X-ray pulsations.  \\object{PSR J0751+1807} appears to show a single pulse, whereas \\object{PSR J1012+5307} may show some evidence for two pulses per pulse period."
    },
    "0402/astro-ph0402116_arXiv.txt": {
        "abstract": "We present the optical spectra of four newly discovered gravitational lenses from the Cosmic Lens All-Sky Survey (CLASS). These observations were carried out using the Low Resolution Imaging Spectrograph on the W. M. Keck-I Telescope as part of a program to study galaxy-scale gravitational lenses. From our spectra we found the redshift of the background source in CLASS B0128+437 ($z_s=3.1240\\pm0.0042$) and the lensing galaxy redshifts in CLASS B0445+123 ($z_l =0.5583\\pm0.0003$) and CLASS B0850+054 ($z_l=0.5883\\pm0.0006$). Intriguingly, we also discovered that CLASS B0631+519 may have two lensing galaxies ($z_{l,1}=0.0896\\pm0.0001$, $z_{l,2}=0.6196\\pm0.0004$). We also found a single unidentified emission line from the lensing galaxy in CLASS B0128+437 and the lensed source in CLASS B0850+054. We find the lensing galaxies in CLASS B0445+123 and CLASS B0631+519 ($l,2$) to be early-type galaxies with Einstein Radii of $2.8-3.0~h^{-1}$~kpc. The deflector in CLASS B0850+054 is a late-type galaxy with an Einstein Radius of $1.6~h^{-1}$~kpc. ",
        "introduction": "Gravitational lensing has the ability to probe the internal mass distributions of galaxies at cosmological distances. The ability to form complete samples, based solely on mass, has allowed studies of the formation and evolution of early-type galaxies at intermediate redshifts (e.g. \\citealt*{keeton98,treu02,koopmans03}; \\citealt{rusin03}). Anomalous flux-density ratios observed in merging gravitational lens images have recently been used to argue for cold dark matter substructure within the haloes of lensing galaxies \\citep{metcalf01,keeton01,chiba02,metcalf02,dalal02,bradac02,keeton03}. The magnification of the lensed source has also allowed studies of star-formation at high redshift (e.g. \\citealt{ebbels96,ellis01}) and the dust content of quasar host galaxies to be estimated \\citep{barvainis02}. Furthermore, gravitational lenses have proved to be powerful tools for determining the cosmological parameters. Lenses with well-constrained mass models and accurate time-delay measurements have been used to calculate the Hubble parameter (e.g. \\citealt{schechter97,lovell98,biggs99,fassnacht99,koopmans99,fassnacht02a,treu02a}). The gravitational lensing statistics from large systematic surveys have provided complementary and independent constraints on the cosmological constant and density parameters (\\citealt{kochanek96a}; \\citealt*{falco98}; \\citealt{helbig99,chae02,chae03}) to those obtained from SN1a \\citep{riess98,perlmutter99}, large-scale structure \\citep{percival01,tegmark02} and cosmic microwave background \\citep{slosar03,sievers03,spergel03} observations. In addition, the nature of the curvature parameter can also be determined from the image-separations and source redshifts via the $\\Delta\\theta$-$z_{s}$ relation \\citep*{turner84,gott89,park97,williams97,helbig98}. However, each of these applications of gravitational lensing is critically dependent on the redshifts of the lensing galaxy and lensed source being known.  With the objectives outlined above in mind, a spectroscopic survey of galaxy-scale gravitational lenses using the W. M. Keck and Palomar Observatories has been underway over the last few years \\citep{kundic97,fassnacht98,lubin00}. In this paper we present our latest results. The observing sample consisted of four gravitational lenses discovered during the course of the Cosmic Lens All-Sky Survey (CLASS; \\citealt{myers02,browne02}). In Section \\ref{spec-targets} a short review of each of these gravitational lenses is given. The observations with the W. M. Keck-I Telescope are presented in Section \\ref{spec-obs} and the optical spectra are presented in Section \\ref{spec-results}. The resulting implications for each gravitational lens are discussed in Section \\ref{spec-disc}. Finally, in Section \\ref{spec-conc} a summary of the observations and analysis presented in this paper is given. Throughout we adopted an $\\Omega_{M}=0.3$, $\\Omega_{\\Lambda}=0.7$ flat-universe, with a Hubble parameter $H_{0}=100~h$~km~s$^{-1}$~Mpc$^{-1}$. ",
        "conclusions": "\\label{spec-conc}  We have presented new optical spectra of four recently discovered gravitational lenses from CLASS. The redshift of the background source in CLASS B0128+437 was found to be $z_s=3.1240\\pm0.0042$ and the lensing galaxy redshifts in CLASS B0445+123 and CLASS B0850+054 were found to be $z_{l}=0.5583\\pm0.0003$ and $z_{l}=0.5883\\pm0.0006$ respectively. Interestingly, we discovered that CLASS B0631+519 may have two lensing galaxies. One is a low luminosity galaxy at $z_{l,1}=0.0896\\pm0.0001$ which has little effect on the lensing potential, and the other is the primary lensing galaxy at $z_{l,2}=0.6196\\pm0.0004$. The other possibility is that the low redshift galaxy is a dark/dust enshrouded gravitational lens. However, the scenario in which the main lens is at $z = 0.6196$ is currently favoured. We find the lensing galaxies in CLASS B0445+123 and CLASS B0631+519 ($l,2$) to be early-type galaxies, with Einstein Radii of $2.8$ and $3~h^{-1}$~kpc respectively. The lensing galaxy in CLASS B0850+054 was found to be consistent with a late-type galaxy with an Einstein Radius of 1.6~$h^{-1}$~kpc.  For the redshift data presented in this paper to be useful for studying galaxy formation at high redshift or investigating the cosmological parameters, they must be coupled with the redshifts of the missing components. We have identified isolated emission lines in the CLASS B0128+437 and CLASS B0850+054 spectra which we hypothesised identifications for. However, these tentative identifications will only be confirmed by further observations. Of the 22 gravitational lenses discovered by CLASS, 17 have measured lensing galaxy redshifts and 12 have background source redshifts recorded. Obtaining spectra of the remaining gravitational lenses with unknown source or lens redshifts is an essential step in the follow-up of the CLASS gravitational lensing survey. To achieve this goal, further deep spectroscopic observations are required in the optical and near-infrared."
    },
    "0402/astro-ph0402320_arXiv.txt": {
        "abstract": "Strong intergalactic shocks are a natural consequence of structure formation in the universe. These shocks are expected to deposit large fractions of their thermal energy in relativistic electrons ($\\xi_e\\simeq 0.05$ according to supernova remnant observations) and magnetic fields ($\\xi_B\\simeq 0.01$ according to cluster halo observations). We calculate the synchrotron emission from such shocks using an analytical model, calibrated and verified based on a hydrodynamical \\LCDM simulation. The resulting signal composes a large fraction (up to a few $10\\%$) of the extragalactic radio background below $500 \\MHz$. The associated angular fluctuations, e.g. $\\delta T_l \\ga 260 (\\xi_e\\xi_B/5\\times 10^{-4}) (\\nu/100\\MHz)^{-3}\\K$ for multipoles $400\\la l\\la 2000$, dominate the radio sky for frequencies $\\nu \\la 10 \\GHz$ and angular scales $1\\arcmin\\la\\theta < 1\\de$ (after a modest removal of discrete sources), provided that $\\xi_e \\,\\xi_B\\ga 3\\times 10^{-4}$. The fluctuating signal is most pronounced for $\\nu \\la 500 \\MHz$, dominating the sky there even for $\\xi_e\\,\\xi_B = 5\\times 10^{-5}$. The signal will be easily observable by next generation telescopes such as LOFAR and SKA, and is marginally observable with present-day radio telescopes. The signal could also be identified through a cross-correlation with tracers of large-scale structure (such as \\gama-ray emission from intergalactic shocks), possibly even in existing $\\la 10\\GHz$ CMB anisotropy maps and high resolution $\\sim 1\\GHz$ radio surveys. Detection of the signal will provide the first identification of intergalactic shocks and of the warm-hot intergalactic medium (believed to contain most of the baryons in the low-redshift universe), and gauge the unknown strength of the intergalactic magnetic field. We analyze existing observations of the diffuse radio background below $500\\MHz$, and show that they are well fit by a simple, double-disk Galactic model, precluding a direct identification of the diffuse extragalactic radio background. Modelling the frequency-dependent anisotropy pattern observed at very low ($1$--$10\\MHz$) frequencies can be used to disentangle the distributions of Galactic cosmic-rays, ionized gas, and magnetic fields. Space missions such as the {\\it Astronomical Low Frequency Array} (ALFA) will thus provide an important insight into the structure and composition of our Galaxy. ",
        "introduction": "\\label{sec:Introduction} \\indent  The gravitational formation of structure in the universe inevitably produced strong, collisionless shocks in the intergalactic medium (IGM), owing to the convergence of large-scale flows. In these shocks, electrons are expected to be Fermi accelerated up to highly relativistic ($\\ga 10 \\TeV$) energies, limited by inverse-Compton cooling off cosmic microwave background (CMB) photons \\cite{LoebWaxman2000}. The resulting \\gama-ray emission is expected to trace the large-scale structure of the universe. Rich galaxy clusters, characterized by strong, high velocity accretion shocks, should be detected by future \\gama-ray missions as bright \\gama-ray sources \\cite{LoebWaxman2000, Totani00, WaxmanLoeb2000} in the form of accretion rings with bright spots at their intersections with galaxy filaments \\cite{Keshet03}. Recently, a possible association of $\\gamma$-ray radiation (as measured by the \\emph{Energetic Gamma-Ray Experiment Telescope}, EGRET) with the locations of Abell clusters was identified at a $3\\sigma$ confidence level \\cite[][but see also Reimer et al. 2003]{Scharf2002}. The integrated \\gama-ray background resulting from strong intergalactic shocks was calculated using hydrodynamical cosmological simulations \\cite{Keshet03,Miniati02}, at the level of $\\epsilon^2 dJ/d\\epsilon \\la 0.15 \\keV \\se^{-1} \\cm^{-2} \\sr^{-1}$. This signal constitutes $\\la 15\\%$ of the widely accepted estimates for the flux of the extragalactic \\gama-ray background \\cite[EGRB, see e.g.] [] {Sreekumar98,Strong03}, although a recent analysis \\cite{Keshet04a} suggests that the EGRB flux is lower than previously thought, at least by a factor of 2. The \\gama-ray background from weak merger shocks was estimated to be a factor of $>10$ weaker than the background from intergalactic accretion shocks \\cite{Gabici03}.  In addition to the inverse-Compton emission from intergalactic shocks, synchrotron radiation should also be emitted by the relativistic electrons, as they gyrate in the shock-induced magnetic fields. The resulting radio signature is expected to trace the structure of the universe at low redshifts ($z\\la1$), and to be dominated by rich, young galaxy clusters. Indeed, extended radio emission with no optical counterpart is observed in about $10\\%$ of the rich galaxy clusters \\cite[radio halos and radio relics, see ] [] {Giovannini99}, and in more than a third of the young, massive clusters \\cite[with X-ray luminosity $L_X>10^{45} \\erg \\se^{-1}$, see][]{Feretti03}. The radio emission has been identified as synchrotron radiation from relativistic electrons, but there are different models for the origin of the electrons and the magnetic fields involved \\cite[for a recent discussion, see][]{Bagchi03}. En\\ss{}lin et al. (1998) have used observations of nine radio relics to suggest an association between these sources and structure formation shocks, focusing on the possibility that the relics are revived fossil radio cocoons originating from nearby radio galaxies, re-energized by diffusive shock acceleration. It is interesting to note that recently, large-scale diffuse radio emission was discovered around a filament of galaxies, possibly tracing an accretion shock on this scale \\cite{Bagchi_etal02}.  Waxman \\& Loeb (2000) have proposed a simple model, which allows one to estimate both the radio and the \\gama-ray signatures of intergalactic shocks, produced by electrons accelerated from the inflowing plasma. Their model allows one to estimate the radio and the \\gama-ray backgrounds and their anisotropy characteristics, as well as the signature of individual clusters. The model uses dimensional analysis arguments to estimate the properties of the virialization accretion shock of a halo, as a function of redshift $z$ and halo mass $M$. Halo abundance estimates at different redshifts [such as the Press \\& Schechter (1974) halo mass function], may then be used to calculate various observables. The Waxman \\& Loeb model approximates the strong accretion shocks as being spherically symmetric, and neglects weak merger shocks. A fraction $\\xi_e \\simeq 0.05$ of the shock thermal energy is assumed to be carried by relativistic electrons, based on supernovae remnant (SNR) observations \\cite[for a discussion, see] [] {Keshet03}. Observations of $\\ga 0.1 \\muG$ magnetic fields in cluster halos \\cite{Kim89, Fusco-Femiano99, Rephaeli99} require that a fraction $\\xi_B\\simeq 0.01$ of the shock thermal energy be transferred into downstream magnetic fields. With these assumptions, Waxman \\& Loeb (2000) found that strong fluctuations in the predicted radio signal seriously contaminate CMB anisotropy measurements at sub-degree angular scales and frequencies below $10\\GHz$.  Identification of radio or \\gama-ray emission from intergalactic shocks holds a great promise for advancing current knowledge of shock formation in the IGM. It should provide the first direct evidence for such shocks, revealing the underlying large-scale cosmological flows. When combined with \\gama-ray detection, the radio signal will provide a direct measure of the unknown magnetic fields in the IGM, possibly shedding light on the processes leading to IGM magnetization. Emission from large-scale shocks traces the undetected warm-hot IGM at temperatures $10^5\\K \\la T\\la 10^7 \\K$ that is believed to contain most of the baryons in the low redshift universe \\cite[see e.g.][]{Davea01}. Moreover, the signal may be used to study non-thermal physical processes in the intergalactic environment, such as Fermi acceleration in low density shocks.  In this paper we explore the radio signature of intergalactic shocks and examine its observational consequences. Our results derive from a generalized version of the model of Waxman \\& Loeb (2000), adapted for a \\LCDM universe, modified to incorporate non-spherical accretion shocks, and calibrated using the global features of a simulated \\LCDM universe in a hydrodynamical cosmological simulation. We perform an analysis of the radio sky in order to assess the feasibility of observing the predicted signal; in this context, we review various known foreground and background radio signals, and analyze observations at low ($\\nu < 500\\MHz$) frequencies, where the diffuse extragalactic background is most pronounced with respect to the Galactic foreground. Our results have implications for existing high resolution radio telescopes (e.g. the Very Large Array\\footnote{see http://www.vla.nrao.edu/astro/guides/vlas}); for next generation ground-based radio telescopes such as the LOw Frequency Array (LOFAR\\footnote{see http://www.lofar.org}) and the Square Kilometer Array (SKA\\footnote{see http://www.skatelescope.org}); and for future space-borne telescopes such as the Astronomical Low Frequency Array (ALFA\\footnote{see http://sgra.jpl.nasa.gov/html\\_dj/ALFA.html}).  In \\S\\ref{sec:IGM_model} we calculate the radio signal from intergalactic shocks. We start by generalizing the model of Waxman \\& Loeb (2000, illustrated in Figure \\ref{fig:PS_illustration}), and adapting it for a \\LCDM universe. The non-uniformity of the thermal injection rate along the shock front is incorporated into the model by modifying the modelled shock front area. The free parameters of the model are calibrated using \\emph{global} features (such as the average baryon temperature) of a simulated \\LCDM universe, according to a previously analyzed hydrodynamical cosmological simulation \\cite{Springel2001}. The predictions of the calibrated model are then shown to agree with the results of the cosmological simulation, regarding the radio emission (see Figures \\ref{fig:calculated_spectrum}, \\ref{fig:calculated_correlation} and \\ref{fig:synch_Cl}) and its \\gama-ray counterpart \\cite[see][]{Keshet03}. We estimate the energy fractions $\\xi_e$ and $\\xi_B$ using SNR and cluster halo observations, and evaluate the uncertainty of these parameters. A qualitative agreement is shown to exist between observations of galaxy-cluster radio halos, and model predictions (see Figure \\ref{fig:clusters}).  In \\S\\ref{sec:feasibility} we examine the observational consequences of the radio signal for present and future radio telescopes. We assess the contribution of various foreground and background signals to the brightness of the radio sky (see Figure \\ref{fig:LFRB_sky}) and to the angular power spectrum (APS, see Figure \\ref{fig:Sky_Cl}). In particular, we examine the synchrotron foreground from our Galaxy, and the contamination from discrete radio sources. We also calculate the point-source removal threshold required so that the fluctuations in the intergalactic shock signal dominate the angular power spectrum on small angular scales. Possible methods by which the signal may be identified using present and future observations are discussed. Other extragalactic radio signals, namely bremsstrahlung emission from \\Lya clouds and 21 cm tomography, are also reviewed.  In \\S\\ref{sec:LFRB} we analyze the diffuse low frequency ($\\nu <500 \\MHz$) radio background (LFRB). After highlighting important observational features (see Figure \\ref{fig:LFRB_raw}), we present a model for the Galactic foreground that allows us to (i) try to separate between the Galactic foreground and the extragalactic background; (ii) examine if a simple Galactic model can account for the observed Galactic foreground; and (iii) demonstrate the importance of observations in the frequency range $1\\MHz \\la \\nu \\la 10\\MHz$, where absorption in our Galaxy is significant. We show that existing observational data is consistent with a simple double-disk Galactic model (see Figure \\ref{fig:LFRB_model}). The implied lack of direct evidence for a diffuse extragalactic radio background (DERB) is discussed, given earlier models for the LFRB and unsubstantiated claims for direct identifications of the extragalactic component. Finally, we show how future observations at very low ($1-10\\MHz$) frequencies, e.g. by the ALFA, may be used to disentangle the distributions of Galactic cosmic-rays, magnetic fields and ionized gas.  Finally, \\S\\ref{sec:discussion} summarizes our results and addresses their potential implications. We discuss the conditions under which emission from intergalactic shocks could be identified, taking into account various present and future telescope parameters, confusion with competing signals, and possible systematic errors in our model. Some consequences of a future positive detection of the signal are mentioned. ",
        "conclusions": "\\indent  We conclude that the design of next generation radio telescopes such as the LOFAR, the SKA and the ALFA, is more than sufficient for detection of the angular fluctuations introduced by intergalactic shocks, as calculated in \\S \\ref{sec:IGM_model} and confirmed by a cosmological simulation \\cite{Keshet04b}. Identification of the signal is limited by confusion with Galactic foreground and with discrete radio sources. Foreground fluctuations in the synchrotron emission of our Galaxy constrain a clear detection of the signal to sub-degree scales and to high Galactic latitudes. Confusion with discrete radio sources requires that the brightest sources be modelled and removed. With a feasible point source cut ($100\\mJy$ at $1.4 \\GHz$, or $3\\Jy$ at $150\\MHz$), the signal dominates over the competing signals at angular scales $10\\arcmin \\la \\theta \\la 1\\de$ (by roughly an order of magnitude according to the \\LCDM simulation), whereas detection on arcminute scales will require a more ambitious point source removal. The spectrum of emission from intergalactic shocks indicates that the signal is most pronounced at $\\sim 100\\MHz$ frequencies, planned to be covered by both the LOFAR and the SKA. We find that other extragalactic signals in the radio band, namely bremsstrahlung from \\Lya clouds and angular (but perhaps not spectral) 21 cm tomography, will be confusion limited by intergalactic shocks as well as by radio point sources.  Interestingly, detection of the fluctuating signal from intergalactic shocks is just possible with \\emph{present} high resolution radio telescopes, such as the VLA (at the maximal sensitivity configuration, see Figure \\ref{fig:Sky_Cl}). The calculated signal should be detectable at high latitudes, a few $100\\MHz$ frequencies and sub-degree scales. As noted by Waxman \\& Loeb (2000), the signal may have already been detected by CMB anisotropy studies, at $\\sim 10\\GHz$ frequencies and $\\theta <1\\de$ scales. In particular, the signal may have been detected in arcminute scales by low frequency CMB anisotropy experiments such as the Australian Telescope Compact Array (ATCA), Ryle, and the VLA. It is possible, that some of the noise removed from the corresponding maps \\cite[e.g.] [] {Subrahmanyan2000} is associated with emission from intergalactic shocks, and is thus correlated with tracers of large-scale structure.  The high resolution of present $\\sim \\GHz$ high resolution surveys raises the possibility that their maps already include the signal at an identifiable level. Such surveys are mostly available at low latitudes (say, $|b|<8\\de$), and thus exhibit stronger Galactic foreground with flatter angular power spectra than found in high latitudes. This implies significant foreground contamination, especially on small angular scales. Some low resolution studies have found an APS power index $\\beta\\simeq 2.4$ extending up to $l\\simeq 900$ \\cite{Giardino01, Tegmark00}, others finding $\\beta\\simeq 1.7$ extending up to $l\\simeq 6000$ \\cite{Tucci02}. In addition, detection of the signal becomes increasingly difficult for frequencies away from the optimal frequency, $\\sim 100\\MHz$. Nevertheless, large longitudinal fluctuations have been identified in the low latitude APS \\cite[e.g.] []{Baccigalupi01}, suggesting that the signal may surface in the 'quiet regions'. Analysis of such regions, including careful removal of discrete sources, may yield the desired signal, identifiable at $\\sim 10^\\prime$ scales by cross-correlating the maps with known tracers of large scale structure.     \\label{sec:discussion}  We have studied the radio signal produced by synchrotron emission from the strong intergalactic shocks associated with structure formation (see \\S \\ref{sec:IGM_model}). The analytical model of Waxman \\& Loeb (2000) was generalized by adapting it for a \\LCDM universe, incorporating spectral features and shock asymmetry into the model, and calibrating the free parameters of the model, as summarized in Table \\ref{tab:ModelParams}. The halo parameters $f_{acc}$, $f_T$, and $f_{r}$ were calibrated using a hydrodynamical cosmological simulation, by demanding that the model agrees with the simulation on the average baryon temperature, the average mass consumption rate by strong shocks, and the typical size of regions where the thermal energy is injected by shocks. After calibrating the model with these essentially global features, it yields radio (see Figures \\ref{fig:calculated_spectrum}, \\ref{fig:calculated_correlation} and \\ref{fig:synch_Cl}) and \\gama-ray \\cite[see] []{Keshet03} signals which are in good agreement with the signals extracted independently from the simulation. Although the parameter calibration scheme is inaccurate and the weak redshift dependence of the parameters needs yet to be determined by a more careful analysis, this agreement suggests that the calibration procedure is sensible. The localized nature of regions where most of the thermal energy is injected \\cite[e.g. at the intersections of X-ray cluster accretion shocks with galaxy filaments, channelling gas into clusters, see] [] {Keshet03}, enhances the synchrotron luminosity of a halo while leaving its inverse-Compton luminosity unchanged, because of the enhanced magnetic energy density. The radio luminosity-temperature relation according to the calibrated model was shown to be in qualitative agreement with observations of cluster radio halos (see Figure \\ref{fig:clusters}).  We have examined the observational consequences of the predicted radio signal, for present-day and for future radio telescopes (see \\S \\ref{sec:feasibility}). Our main results are illustrated in Figures \\ref{fig:LFRB_sky} and \\ref{fig:Sky_Cl} (for the calibrated model parameters discussed in \\S \\ref{subsec:model_LCDM_sim} and summarized in Table \\ref{tab:ModelParams}). Figure \\ref{fig:LFRB_sky} shows the contribution of various signals to the radio sky, suggesting that emission from intergalactic shocks contributes up to a few tens of percent of the extragalactic radio background below $500\\MHz$ (see also \\S \\ref{sec:LFRB}). Figure \\ref{fig:Sky_Cl} depicts the angular power spectrum of various signals on an angular scale of $\\sim0\\de.5$, along with the sensitivities and the angular resolutions of the telescopes. The figure demonstrates that the designs of next generation radio telescopes such as the LOFAR and the SKA, are more than sufficient for a detection of the angular fluctuations introduced by intergalactic shocks. In fact, even present-day high-resolution radio telescopes (such as the VLA in its maximum sensitivity configuration), are potentially sensitive to the predicted signal.  Foreground fluctuations in the synchrotron emission of our Galaxy limit a positive detection of the signal to sub-degree scales and to high Galactic latitudes. Confusion with discrete radio sources requires that the brightest sources be modelled and removed. With a feasible point source cut ($100\\mJy$ at $1.4 \\GHz$ or equivalently $3\\Jy$ at $150\\MHz$), the intergalactic signal dominates over all foreground and background signals, at angular scales $10\\arcmin \\la \\theta \\la 1\\de$ and frequencies $\\nu<10\\GHz$ (assuming $\\xi_e \\xi_B \\ga 3\\times 10^{-4}$, see \\S \\ref{subsec:model_efficiency} and the discussion below). For the above cut, the signal is of the same magnitude as the integrated emission from discrete sources at $\\sim1\\arcmin$ scales. Attempts to detect the signal at $\\la1\\arcmin$ scales are confusion limited by discrete sources, demanding increasingly more ambitious point source removal schemes. In addition to the labor involved, noise is always introduced when removing discrete sources, because of source model uncertainties. The spectrum of the signal suggests that it is most pronounced at frequencies around $\\sim 100\\MHz$, planned to be covered by both the LOFAR and the SKA.  The calculated level of angular fluctuations in the radio emission from intergalactic shocks is sensitive to uncertainties in the calibration of the model parameters. The logarithmic contribution to the variance (shown in Figure \\ref{fig:Sky_Cl}) scales roughly according to \\beq \\delta I_l(l;\\,\\xi_e,\\xi_B,f_{acc},f_T,\\widetilde{f}_{r}) = \\xi_e \\xi_B {f_{acc} f_T^2 \\over \\widetilde{f}_{r}^3} \\, \\delta I_l (\\widetilde{f}_{r}l; \\,1,1,1,1,1) \\fin \\eeq The radio signal calculated from the \\LCDM simulation \\cite{Keshet04b}, however, depends only on the energy fractions $\\xi_e$ and $\\xi_B$, through the combination $\\xi_e \\xi_B$. The results of the simulation, as shown in Figures \\ref{fig:calculated_spectrum}, \\ref{fig:calculated_correlation}, and \\ref{fig:synch_Cl}-\\ref{fig:Sky_Cl}, suggest that $\\delta I_l$ was in fact slightly \\emph{underestimated} by our parameter calibration scheme. This implies, for example, that intergalactic shocks introduce intensity fluctuations of magnitude $\\delta I_l \\ga 8\\times 10^{-19} (\\xi_e \\, \\xi_B / 5\\times 10^{-4}) (\\nu/100\\MHz)^{-1} \\IUnits\\!\\!$ on multipoles $400\\la l \\la 2000$, corresponding to temperature fluctuations $\\delta T_l \\ga 260 (\\xi_e \\, \\xi_B / 5\\times 10^{-4}) (\\nu/100\\MHz)^{-3} \\K$.  Our model implies that emission from intergalactic shocks on angular scales $\\sim 0\\de.5$ dominates the sky at $<500\\MHz$ frequencies if $\\xi_e \\xi_B\\ga 10^{-4}$, and is even dominant at frequencies as high as $10\\GHz$ if $\\xi_e \\xi_B\\ga 3 \\times 10^{-4}$. The simulation suggests that the signal is stronger than predicted by the (calibrated) model, in particular on small angular scales. Therefore, according to the simulation, emission from intergalactic shocks will dominate the sky at $<0\\de.5$ scales and $<500\\MHz$ frequencies, even if $\\xi_e \\xi_B \\simeq 5\\times 10^{-5}$. We have used observations of SNR shocks and of magnetic fields in the halos of galaxy clusters, in order to show that for strong intergalactic shocks $\\xi_e \\xi_B \\simeq 5 \\times 10^{-4}$, unlikely to be smaller than this value by more than a factor of $\\sim 4$. For this purpose, in \\S \\ref{subsec:model_efficiency} we have used dimensional analysis arguments to show that the physics of strong intergalactic shocks is essentially identical to the physics of strong SNR shocks, both of comparable velocities $v\\simeq 10^3\\km \\se^{-1}$, provided that an appropriate re-scaling of time is carried out \\cite[see also] [] {Keshet03}. We have assumed that strong shocks accelerate electrons to a power law distribution of index $p=2$ (equal energy per logarithmic interval of electron energy). Such a distribution is inferred from SNR observations and agrees with linear models for diffusive shock acceleration, although non-linear models suggest deviations from a pure power-law.  As noted by Waxman \\& Loeb (2000), emission from intergalactic shocks may have already been detected by CMB anisotropy studies at frequencies $\\la 10\\GHz$ and angular scales smaller than a degree. In particular, telescopes operating at frequencies of a few GHz such as ACSA, Ryle and the VLA, may have detected the signal at arcminute scales. High resolution $\\sim 1\\GHz$ radio surveys may have also detected the signal by now at an identifiable level. Such surveys are generally available at low Galactic latitudes, where contamination by Galactic foreground fluctuations is severe. Nevertheless, analysis of 'quiet regions' observed off the Galactic plane, including a careful removal of discrete sources, may yield the desired signal. It will probably be easiest to identify the signal at $1\\arcmin-10\\arcmin$ scales, by modelling discrete sources and cross correlating the maps with known tracers of large-scale structure. Detection of emission from intergalactic shocks is not unrealistic even in case the signal has been mildly overestimated. A signal $\\delta I_l$ lower than calculated in \\S\\ref{sec:IGM_model} by a factor of a few, may still be identified at $\\sim 1\\arcmin-10\\arcmin$ scales, if faint point sources are modelled and removed. If the signal was overestimated by $\\sim 2$ orders of magnitude, it may still be detectable by next generation radio telescopes, by means of cross-correlation with known tracers of large-scale structure such as galaxy counts, or, in the future, with \\gama-ray emission from intergalactic shocks.  As mentioned in \\S \\ref{sec:Introduction}, future detection of radio emission from intergalactic shocks will have important implications on our understanding of cosmology and astrophysics. Detection of the signal will provide the first identification of intergalactic shocks, revealing the underlying cosmological flows and providing a test for structure formation models. The signal, in particular when combined with \\gama-ray detection, will provide a measure of the unknown magnetic fields in the intergalactic medium. Although non-trivial for interpretation, such a measure of the magnetic field may provide insight into the unknown processes leading to IGM magnetization. The signal may also confirm the existence of the undetected warm-hot intergalactic medium, and provide a tracer for its distribution.  Synchrotron emission from intergalactic shocks is correlated with the large-scale structure of the low-redshift ($z<1$) universe, tracing young galaxy clusters and filaments. The signal could thus account for some observed features of the radio signature of galaxy clusters, namely radio halos and radio relics. Extended accretion shocks could contribute to radio halos \\cite[e.g. to the large radio halo of the Coma super-cluster, see][and the references therein]{Thierbach03}, whereas localized shocks (e.g. where galaxy filaments channel large amounts of gas into the cluster regions) may be responsible for some radio relics observed at the outskirts of clusters [such as the prototype relic found in the Coma super-cluster, 1253+275 \\cite{Ensslin98}, and the large-scale radio arcs observed in clusters A3667 \\cite{Rottgering97} and A3376 \\cite{Bagchi02}], and perhaps also for some of the anomalous features observed in radio halos \\cite[e.g. in the unrelaxed clusters described by] [] {Govoni04}.  Radio emission from intergalactic shocks is an important source of contamination for other radio signals, such as the low frequency CMB, Galactic synchrotron fluctuations on sub-degree scales, and competing extragalactic radio signals such as bremsstrahlung from \\Lya clouds, and 21 cm tomography (Zaldarriaga et al. 2003; Morales \\& Hewitt 2003; Loeb \\& Zaldarriaga 2003; Gnedin \\& Shaver 2003). For example, the integrated emission from intergalactic shocks introduces distortions in the low-frequency average brightness temperature, at the level of $\\delta T \\simeq 1.6 \\,(\\nu / 3 \\GHz)^{-3} \\mK$. Efforts to measure distortions in the CMB spectrum at low frequencies, e.g. in the next mission of the Absolute Radiometer for Cosmology, Astrophysics and Diffuse Emission (ARCADE\\footnote{See http://arcade.gsfc.nasa.gov/arcade/index.html}), may thus be sensitive to the signal at the low frequency bands \\cite[$\\sim 3\\GHz$, see] [] {Fixen04,Kogut04}. We have shown that radio emission from intergalactic shocks constitutes a significant fraction of the extragalactic radio background below $500\\MHz$. This component should thus be taken into account when evaluating the propagation of ultra-high energy photons with energies above $10^{19}\\eV$, because at these energies the effect of the radio background on the transparency of the universe is stronger than the effect of the CMB, and may thus be important in models for the origin of ultra-high energy cosmic-rays.  We have analyzed observations of the diffuse low frequency radio background below $500\\MHz$ (see \\S \\ref{sec:LFRB}), and highlighted the frequency-dependent anisotropy pattern observed in frequencies $1\\MHz \\la \\nu \\la 10\\MHz$ (see Figure \\ref{fig:LFRB_raw}). We presented a simple Galactic model, consisting of a thin disk, containing most of the ionized gas, and a thick disk. This model was shown to provide a good fit to the data (see Figure \\ref{fig:LFRB_model}), considering the observational uncertainties. The model thus enabled us to (i) assess the feasibility of directly inferring the flux of the diffuse extragalactic radio background from presently available observations; and (ii) demonstrate how an elaborate Galactic model may be constructed in a straightforward fashion, using the anisotropy pattern observed in very low frequencies.  The agreement of the Galactic model with observations, and the large uncertainties and poor resolution of present low frequency observations, preclude a direct identification of the diffuse extragalactic radio background. At best, an upper limit on the extragalactic component may be imposed, at very low ($\\la 30$ MHz) frequencies where free-free absorption in the Galaxy becomes important. Such a flux constraint, a factor of a few lower than the Galactic foreground, is itself higher than the calculated signal from discrete sources and from intergalactic shocks by a factor of a few. This implies that (i) the true extragalactic background in $\\nu\\simeq 10\\MHz$ frequencies is probably smaller than the Galactic foreground by a factor of a few; and (ii) emission from intergalactic shocks constitutes a significant fraction of the total extragalactic signal, measuring up to a few tens of percent.  We have reviewed previous models of the low frequency radio background, and demonstrated that previous claims for direct identification of the diffuse extragalactic component are based on insignificant, probably Galactic, spectral features. We pointed out that previous models fail to account for the combined spectral and angular data at very low ($1\\MHz<\\nu<10\\MHz$) frequencies, where absorption is non-negligible and the anisotropy pattern of the sky is highly frequency dependent. Modelling the combined data can be used to disentangle the distributions of Galactic cosmic-rays, ionized gas and magnetic fields, and lead to an elaborate three-dimensional model of the Galaxy. Our model is too simplistic, and the data uncertainties too large, to provide reliable estimates of the physical properties of our Galaxy, other than confirming its double-disk structure and the presence of a $\\sim 1\\kpc$ thick disk of cosmic ray electrons and magnetic fields, and imposing constraints on free-free absorption in the solar neighborhood. Future high resolution observations at very low frequencies, produced for example by the ALFA space mission, could thus provide valuable information regarding the structure and the composition of the Milky Way."
    },
    "0402/astro-ph0402489.txt": {
        "abstract": "{We have obtained deep spectra from 1640 to 10100\\AA\\ with the Space Telescope Imaging Spectrograph (STIS) of the Strontium Filament, a largely neutral emission nebulosity lying close to the very luminous star Eta Carinae and showing an uncommon spectrum. Over 600 emission lines, both permitted and forbidden, have been identified. The majority originates from neutral or singly-ionized iron group elements (Sc, Ti, V, Cr, Mn, Fe, Co, Ni). Sr is the only neutron capture element detected. The presence of Sr II, numerous strong \\ion{Ti}{ii} and \\ion{V}{ii} lines and the dominance of \\ion{Fe}{i} over \\ion{Fe}{ii} are notable discoveries. While emission lines of hydrogen, helium, and nitrogen are associable with other spatial structures at other velocities within the Homunculus, no emission lines from these elements correspond to the spatial structure or velocity of the \\ion{Sr}~Filament. Moreover, no identified \\ion{Sr}~Filament emission line requires an ionization or excitation energy above approximately 8 eV. Ionized gas extends spatially along the aperture, oriented along the polar axis of the Homunculus, and in velocity around the Strontium Filament. We suggest that the Strontium Filament is shielded from ultraviolet radiation at energies above ~8 eV, but is intensely irradiated by the central star at wavelengths longward of 1500\\AA. ",
        "introduction": "The Luminous Blue Variable star (LBV) Eta Carinae ($\\eta$ Car) experienced a major outburst in the 1840's with a secondary outburst in the 1890's \\citep[][ and references therein]{DH97}. Several solar masses of material were ejected, which is now directly seen as the expanding Homunculus \\citep{MDB98,SGH03} and the Little Homunculus \\citep{IGD03}. Although the central star provides 5x10$^6$ L$_{\\odot}$ at a characteristic temperature of $25,000\\degr $K, most of the gas in the Homunculus is neutral \\citep{DSG01}. The hollow bipolar lobes with an intervening equatorial skirt are seen primarily by dust-scattered radiation from $\\eta$ Car. STIS CCD spectra, recorded with a long aperture ($52\\arcsec \\times 0\\farcs2$), revealed a thin, interior skin in the light of [\\ion{Fe}{ii}] and [\\ion{Ni}{ii}], but no \\ion{H}{ii} nor \\ion{He}{i}, emission. \\citet{DSG01} used these lines, and the absorption H and K lines of \\ion{Ca}{ii}, to trace the inner, neutral surface of the Homunculus. Ground-based observations in the near-infrared \\citep{S02} revealed molecular hydrogen in the cool exterior of the Homunculus shell. Lines of \\ion{Fe}{ii}, H$\\alpha$ and [\\ion{Ni}{ii}] revealed an internal emission nebulosity, called the Little Homunculus \\citep{GI01,IGD03}. Several bright, ionized emission structures exist very close to $\\eta$ Car, known as the Weigelt Blobs B, C and D (Weigelt \\& Ebersberger 1986), are located within $0\\farcs1$ to $0\\farcs3$ from the central star.  \\citet{D96} identified the 5.52-year period in the high-excitation nebular and stellar emission lines of $\\eta$ Car and its surrounding nebulosity. In a long-term monitoring series of programs to understand this variability, the Weigelt blobs B and D, along with $\\eta$ Car, have been observed with HST/STIS at nearly annual intervals since 1998 \\citep{DIG99,GID01}. Over 2000 emission lines of the Weigelt blobs B and D were identified in the spectrum between 1700\\AA\\ and 10300\\AA\\ by \\citet{ZPHD01}. Changes between the spectroscopic minimum in 1998 and the broad maximum during 1999 and 2000 demonstrate that lines of higher ionization disappear during the spectroscopic minimum only to reappear as the system recovers. Many \\ion{H} Ly$\\alpha$-pumped \\ion{Fe}{ii} lines appear during the maximum \\citep{JH93,ZPHD01}, and disappear during the minimum. The \\ion{Fe}{ii} 2507, 2509 \\AA\\ lines are the most enhanced of these fluorescence lines, and they feed long-lived Fe II states involved in a closed radiative cycle showing stimulated emission \\citep{JL03}. \\citet{VGB02} used the CLOUDY model to predict the optical Fe II emission fluxes of the Weigelt B and D blobs during the spectroscopic minimum event of 1998.  During a preliminary test for the Homunculus mapping program in the 6400-7000 \\AA\\ region planned as a STIS GTO Key Project (HST proposal 8483), we noticed some very faint, narrow emission lines located 1.5\\arcsec\\ to the Northwest of $\\eta $ Car. The 1\\arcsec\\ long emission filament appeared not to be associated with any known structure in $\\eta$ Car. Yet the spatial and velocity structure was similar for these lines and indicated that they must originate from the same volume. \\citet{ZGH01} identified twenty of twenty one lines in the 6400-7000A\\ region, all originating from a structure moving at -100 km/sec. As the most spectacular discovery was the first identification of two [\\ion{Sr}{ii}] lines, the filament became known as the \\ion{Sr}~Filament. Peculiarly, no lines of hydrogen or helium were identified in the spectrum of this system. While lines of \\ion{Fe}{ii} were not identified, lines of \\ion{Fe}{i} were. However, from the identifications of this wavelength limited spectrum it was not possible to conclude whether these emission lines were due to a selective excitation mechanism or to different elemental abundances.  The limited spectrum of the \\ion{Sr}~Filament differed remarkably from similar spectra of other emission line nebular structures around $\\eta$ Car. This fact led to additional observations and line identifications in other wavelength regions. In the present paper we report on all HST observations obtained so far of the \\ion{Sr}~Filament and tabulate all the measured emission lines. Nearly 600 lines have been identified, and only a few strong lines remain unidentified. We also discuss the peculiarities found in the spectrum in terms of apparent enhancements and depletions in elemental abundances, as well as clear indications of special ionization and excitation conditions in the filament. ",
        "conclusions": "The \\ion{Sr}~Filament proves to be a very unusual emission nebulosity. Over 600 emission lines, mostly from neutral and singly-ionized iron-peak elements, have been identified. Yet no hydrogen, helium, nitrogen or oxygen emission lines, which characterize normal emission regions, have been detected. Several factors contribute to this unlikely emission-nebula spectrum: 1) the very massive star  system, while producing $5\\times10^6 L_{\\odot}$ with a characteristic temperature of $25,000 \\degr K$ \\citep{HDI01}, has a highly clumpy, extended but cool atmosphere \\citep{VBV03}; 2) much ionized gas shields this \\ion{Sr}~Filament from hard UV radiation capable of ionizing hydrogen, and other elements requiring photo-ionizing energies exceeding 13.6eV; 3) apparently the shielding may be sufficient to protect the \\ion{Sr}~Filament from radiation down to 6 or 7 eV as we see no direct evidence of line emission requiring photo-excitation with that energy; 4) an intense radiation field with energies less then 6 or 7 eV does bathe the Sr Filament, leading to a partially-ionized region; 5) the \\ion{Sr}{ii} modelling \\citep{BGI02} indicates a very high-density region as the electron density must be in the range of $10^7$ cm$^{-3}$.  We note that the ejecta surrounding $\\eta$ Car have a very non-uniform structure. While the overall Homunculus is a thin, hollow shell about ten percent thickness compared to the distance from the Central Source \\citep{S02,SGH03} the interior is likely a hot, low-density stellar wind. The thin surface interior to the shell is detected in [\\ion{Fe}{ii}], [\\ion{Ni}{ii}] emissions and \\ion{Ca}{ii} absorption \\citep{DSG01,S02}. In line of sight, ejecta absorptions of the iron-peak elements demonstrate a range of temperature and electron density that correlate with velocity (Gull, et al, \\apj\\ submitted). Interior to the Homunculus is the Little Homunculus, a miniature bilobed structure \\citep{IGD03}, which is seen in multiple emission lines of iron-peak elements and in the hydrogen Balmer lines. Between the bi-lobes of the Homunculus and the Little Homunculus is the skirt region, partially seen in emission lines, and in absorption lines. Close to $\\eta$ Car are several very intense emission knots, the Weigelt Blobs, seen strongly in many \\ion{Fe}{ii}, \\ion{Ni}{ii}, and \\ion{Cr}{ii} lines \\citep{ZPHD01}. Highly-excited emission lines seen in the Weigelt Blobs and in the Little Homunculus disappeared during the 1998.0 and 2003.5 minima, but then returned. Low-excitation lines maintain constancy in flux throughout the 5.52-year period. As the Sr Filament  emission lines are low-excitation only, this reinforces the concept that the \\ion{Sr}~Filament receives radiation with the harder photons filtered out. Based upon the spatial distribution of the blue-shifted velocities, the \\ion{Sr}~Filament is probably located in this equatorial skirt region. At a projected distance of the order of ten light-days, the \\ion{Sr}~Filament receives intense mid-UV, and longer wavelength, radiation from the Central Source. Likely it is the strong absorption by iron and other iron-peak elements in the ionized regions and just beyond the ionized regions that shield the \\ion{Sr}~Filament. We note that the ionization potential of Fe is 7.9 eV and that of Sr is 5.7 eV. As there is abundant \\ion{Fe}{i}  and little \\ion{Fe}{ii} in the \\ion{Sr}~Filament, this indicates that the strontium is singly-ionized, but protected from becoming doubly-ionized by an iron-shield. Moreover, many \\ion{Fe}{ii} absorptions are in the spectrum indicating the abundance of singly-ionized iron in the vicinity of the \\ion{Sr}~Filament. Shortward of 2500\\AA, much of the ultraviolet spectrum is chopped up, further protecting the neutral and singly-ionized species with ionization potentials above 4 or 5 electron volts. Indeed the question arises as to whether molecular species might reside in this region. Ground-based, near-IR observations of the Homunculus \\citep{S02} do not indicate molecular hydrogen at these velocities or spatial position, but the ground-based observations were accomplished with lower spatial resolution.  We have systematically obtained spectra of the \\ion{Sr}~Filament to characterize the spatial extent and the level of excitation through the identification of over 600 emission lines. We have also measured the fluxes of these lines in preparation for obtaining physical information of this neutral emission region. The first paper, based upon the \\ion{Sr}{ii} emission line ratios, has been published, characterizing the temperature and density of the \\ion{Sr}~Filament \\citep{BGI02}, other papers will follow discussing models of other iron-peak neutral and singly-ionized species. We hope to provide information on relative abundances of various ionic species, possibly elemental abundances, but modelling and possibly some laboratory work will first be necessary."
    },
    "0402/astro-ph0402050_arXiv.txt": {
        "abstract": "{ We find that the quiescent accretion disk of U Gem has a complicated structure. Along to the bright spot originating in the region of interaction between the stream and the disk particles, there are also explicit indications of spiral shocks. The Doppler map and the variations of the peak separation of the emission lines are indicative. }  \\addkeyword{Accretion, Accretion Disks} \\addkeyword{Line: Profiles} \\addkeyword{Shock Waves} \\addkeyword{Novae, Cataclysmic Variables} \\addkeyword{Stars: Individual: U Gem}  \\begin{document} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0402/astro-ph0402266_arXiv.txt": {
        "abstract": "Age derivation techniques for unresolved stellar populations at high redshifts are explored using the NUV spectrum of LBDS~53W091 ($z=1.55$) and LBDS~53W069 ($z=1.43$). The photometry and morphology of these galaxies -- which are weak radio sources -- suggest they are early-type systems, a feature that makes them ideal test beds for the analysis of their ages and metallicities with the use of population synthesis models. In the analysis that is based on {\\sl simple} stellar population models, we find a significant degeneracy between the derived ages and metallicities both in optical+NIR photometric and NUV spectroscopic analyses. This degeneracy is not so strong for LBDS~53W069. However even in this case the stellar age cannot be constrained better than to a range roughly encompassing one third of the age of the Universe at $z=1.43$ (90\\% confidence level). We have explored several independent population synthesis models and consistently found similar results. Broadband photometry straddling the rest-frame 4000\\AA\\ break is also subject to a strong age-metallicity degeneracy. The use of realistic chemical enrichment assumptions significantly helps in disentangling the degeneracy. Based on this method, we derive the average stellar age for both galaxies around $\\langle t_\\star\\rangle\\sim 3.6-3.8$~Gyr with better constraints on the youngest possible ages ($\\sim 3$~Gyr at the 90\\% confidence level). The comparison with simple stellar population models suggest sub-solar metallicities ($\\log Z/Z_\\odot =-0.2$). A composite model using chemical enrichment gives slightly higher metallicities in both galaxies ($\\log Z/Z_\\odot =-0.1$). Given that the stellar component in galaxies forms over times which are larger than a typical chemical enrichment timescale, we conclude that composite stellar populations must be used in all photo-spectroscopic analyses of galaxies. From the observational point of view, the most efficient (and feasible) way to set limits on unresolved stellar populations comprises a combination of Balmer absorption lines along with either low SNR rest frame NUV spectroscopy or accurate optical and NIR photometry. ",
        "introduction": "Estimating the age of the unresolved stellar populations observed in galaxies represents a major challenge in our understanding of galaxy formation. An ideal observation should allow us to infer the star formation history of galaxies from a set of various spectrophotometric observables. The light from a recent burst is dominated by the colour of OB stars, whereas the light from old stellar populations come predominantly from G- and K-type giants. Hence, as a zeroth order approximation, broadband colours track the stellar age. However, this stellar clock is not good enough because of the effect of metallicity on age estimates so that -- within the observational uncertainties -- the colours of galaxies, whose light is dominated by old and metal-poor stars, may be indistinguishable from galaxies mainly composed of young and metal-rich stars. Furthermore, this degeneracy cannot be simply overcome by a better spectral resolution as targeted spectral indices such as the Lick/IDS system (Worthey et al. 1994) suffer from an age-metallicity degeneracy similar to broadband photometry (Worthey 1994).  Accurate spectroscopic dating of stellar populations has been attempted over the past years (Stockton, Kellogg \\& Ridgway 1995; Dunlop et al. 1996; Spinrad et al. 1997; Yi et al. 2000) although a comprehensive analysis of the effect of metallicity has not been considered in detail until recently (Nolan et al. 2003; Ferreras, in preparation).  This paper focuses on LBDS~53W069 ($z=1.43$) and LBDS~53W091 ($z=1.55$), two faint radio galaxies from the Leiden-Berkeley Deep Survey (hereafter LBDS; Windhorst et al. 1984a; 1984b). The search for optical counterparts to faint radio emission is a technique which should allow us to spot old stellar populations at high redshifts (Kron, Koo \\& Windhorst 1985). However, the light from many of these objects is dominated not by starlight but by the active nucleus. An attempt was made to target old stellar populations by searching in the LBDS catalog for weak radio sources with faint NIR magnitudes ($K\\leq 18$) and red optical$-$IR colors ($R-K>5$). Among the reddest galaxies, LBDS~53W069 ($K=18.5$; $R-K=6.3$) and LBDS~53W091 ($K=18.7$; $R-K=5.8$) have been extensively studied (Dunlop et al. 1996; Spinrad et al. 1997; Dunlop 1999). Spinrad et al. (1997) presented the spectrum of 53W091 using Keck LRIS, which maps into its rest-frame NUV ($1960<\\lambda/$\\AA\\ $<3500$). Based on a comparison with their simple stellar populations (SSPs), they found a minimum age of $3.5$~Gyr, which imposed a significant constraint on cosmology. However, due to the substantial difference between SSP models (e.g., Yi et al. 2003) and the systematic effects in age derivation techniques, the subsequent analyses of Bruzual \\& Magris (1997), Heap et al. (1998) and of Yi et al. (2000) indicated significantly younger ages ($\\sim 1-2$~Gyr). The controversy has continued. Nolan et al. (2003) have recently performed a more detailed analysis exploring composite stellar populations (i.e. a mixture of metallicities) and found $\\sim 3$~Gyr roughly confirming their first age estimate. The discrepancy between these two age estimates can be translated into a formation redshift. By adopting a $\\Lambda$CDM cosmology with $\\Omega_m=0.27$ and $H_0=71$~km~s$^{-1}$~Mpc$^{-1}$ (Spergel et al. 2003), which is used in this paper hereafter, the literature suggests that the stellar component of LBDS~53W091 could have been formed at a redshift between $z_F\\sim 2.5$ (Yi et al. 2000) and $z_F\\simgt 4.7$ (Nolan et al. 2003).  The controversy over the actual age of an allegedly simple case such as LBDS~53W091 shows that it is imperative to make a robust estimate of the underlying uncertainties. We will explore in this paper the uncertainties inherent to any age estimate using simple fitting techniques to the observed NUV spectra of LBDS~53W069 and 53W091. Given the importance of these estimates to cosmology as well as to galaxy formation, we believe the case study presented in this paper is a timely and relevant exercise which should be borne in mind when extracting ages from the integrated properties of stellar populations.   \\begin{figure} \\includegraphics[width=3.4in]{lbds_f1.eps} \\caption{$\\chi^2$ map for a comparison of the observed SED of galaxies LBDS~53W091 ($z=1.55$; left) and LBDS~53W069 ($z=1.43$) with simple stellar populations from three independent population synthesis models (see text for details). In each panel the grey scales correspond to the $1$, $2$ and $3\\sigma$ confidence levels from centre to outside. The stars give the position of the minimum $\\chi^2$ (see Table~\\ref{tab:ages}).} \\label{fig:c2sed} \\end{figure} ",
        "conclusions": "The observed rest-frame NUV SEDs of high redshift weak radio galaxies have been claimed to be robust estimators for the ages of old stellar populations at high redshifts, which in turn allow us to set constraints on the age of the Universe and on cosmological parameters (Spinrad et al. 1997). However, in this paper we show that the combined effect of age and metallicity results in large error bars which are shown to be independent of the population synthesis model used. This problem persists even if we use the spectral energy distribution over a wide range of wavelengths instead of a set of broadband filters. Only at SNR=10 or greater can the data disentangle the degeneracy. Unfortunately, this corresponds to prohibitively long exposure times on a 10~m-class telescope.  A comparison of synthetic SEDs built from SSPs, with noise mimicking that of the observed data, shows that signal to noise ratios close to those used in the analysis of LBDS~53W091 (i.e. SNR$\\sim$ 3) are not high enough to yield age estimates with appreciable precision. Slightly higher SNR will help in reducing the error bar, but the age-metallicity degeneracy is still very strong unless one can achieve a SNR=10 in rest frame NUV spectroscopy or even higher in the optical spectral window. In principle, it may appear that the rest-frame NUV would be a desirable window to explore for these galaxies since it could pose stronger constraints on the allowed region of parameter space compared to a rest-frame optical SED. However, the weaker stellar continuum and the strong effect of dust in this spectral region along with our poorer knowledge of stellar emission in the NUV imply that it may be more feasible and useful to obtain high precision rest-frame optical photometry straddling the 4000\\AA\\  break. Old stellar populations are much brighter redward of the break, which results in higher SNRs.  Dating unresolved stellar populations from their spectral energy distribution over a wide range of wavelengths is strongly dependent on the overall shape of the continuum. A small error in the flux calibration will distort this shape, thus altering the estimated ages and metallicities. Hence, we want to emphasize that this type of studies requires SEDs with a very accurate flux calibration. As has been widely known for a decade, the use of Balmer absorption lines, in combination with broadband photometry, largely solves the problem, and helps us in understanding a possible systematic effect derived from uncertainties in the flux calibration of the SED. As we discussed in \\S 3, different concerns would lead observers to choose different Balmer indices. At $z=1.5$ indices such as H$\\beta$, H$\\gamma$ and H$\\delta$ all appear in the NIR spectral window, complicating an accurate ground-based spectroscopic measurement.  Hoping to break the infamous age-metallicity degeneracy from a theoretical point of view, we have explored a large set of chemically consistent $composite$ population models. Simply because no arbitrary metallicity is allowed in such a scheme, this approach helped us determine metallicities much better. We found a better constraint on the age estimates, giving a range of ages between $2.8$ and $4$~Gyr at the $3\\sigma$ confidence level for both galaxies, with a metallicity around solar with a $\\pm 0.5$~dex uncertainty. LBDS~53W069 seems to accept models with higher average metallicities. Our results -- involving different sets of population synthesis models and a detailed chemical enrichment scenario -- give similar results as the analysis of Nolan et al. (2003). Even though any model of chemical enrichment introduces further uncertainties in the modelling, galaxies should be considered composite models as the star formation takes place over times which are always longer than the characteristic chemical enrichment timescales. It is worth noticing that the best fits for a simple and a composite stellar population are marginally compatible (see figure~\\ref{fig:c2chem}). Hence, at the expense of adding further uncertainties to the modelling, we believe a proper mixture of stellar ages and metallicities should be considered in all photo-spectroscopic analyses of galaxies."
    },
    "0402/hep-ph0402279_arXiv.txt": {
        "abstract": "The existence of extra dimensions allows the possibility that the fundamental scale of gravity is at the TeV. If that is the case, gravity could dominate the interactions of ultra-high energy cosmic rays. In particular, the production of microscopic black holes by cosmogenic neutrinos has been estimated in a number of papers. We consider here gravity-mediated interactions at larger distances, where they can be calculated in the eikonal approximation. We show that for the expected flux of cosmogenic neutrinos these elastic processes give a stronger signal than black hole production in neutrino telescopes. Taking the bounds on the higher dimensional Planck mass $M_{D}$ $(D=4+n)$ from current air shower experiments, for $n=2\\ (6)$ elastic collisions could produce up to 118 (34) events per year at IceCube. On the other hand, the absence of any signal would imply a bound of $M_D\\gsim 5$ TeV. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0402/astro-ph0402338.txt": {
        "abstract": "{As stars close to the galactic centre have short orbital periods it has been possible to trace large fractions of their orbits in the recent years. Previously the data of the orbit of the star S2 have been fitted with Keplerian orbits corresponding to a massive black hole (MBH) with a mass of M$_{BH}$=3-4$\\times$10$^6$M$_{\\odot}$ implying an insignificant cusp mass. However, it has also been shown that the central black hole resides in a $\\sim$1'' diameter stellar cluster of a priori unknown mass.  In a spherical potential which is neither Keplerian nor harmonic, orbits will precess resulting in inclined rosetta shaped trajectories on the sky.  In this case, the assumption of non-Keplerian orbits is a more physical approach.  It is also the only approach through which cusp mass information can be obtained via stellar dynamics of the cusp members.  This paper presents the first exemplary modelling efforts in this direction. Using positional and radial data of star S2, we find that there could exist an unobserved extended mass component of several 10$^5$M$_{\\odot}$ forming a so-called 'cusp' centered on the black hole position. Considering only the fraction of the cusp mass M$_{S2_{apo}}$ within the apo-center of the S2 orbit we find as an upper limit that M$_{S2_{apo}}$/(M$_{BH}$ + M$_{S2_{apo}}$) $\\le$ 0.05. A large extended cusp mass, if present, is unlikely to be composed of sub-solar mass constituents, but could be explained rather well by a cluster of high M/L stellar remnants, which we find to form a stable configuration. ",
        "introduction": "Over the last decade, evidence has been found for the existence of massive black holes (MBH) in the centres of many nearby galaxies. With increasing observational data -- 12 MBHs candidates detected until 1995 (Kormendy \\& Richstone 1995), and more than 37 until 2001 (Kormendy 2001,  Ferrarese et al. 2001) -- it is argued that most galaxies harbor nuclei dominated by MBHs with masses that range between 10$^6$ and 10$^{9.5}$ solar masses.  Located at a distance of only $\\sim$8 kpc from the solar system (Reid 1993; Eisenhauer 2003), the Galactic Centre (GC) is the closest and therefore best object for investigating physical processes in the galactic nucleus of a typical spiral galaxy.  It offers a unique ``laboratory'' for studying stars and gas in the sphere of influence on a super-massive black hole, (e.g., Genzel, Hollenbach, $\\&$ Townes 1994; Morris $\\&$ Serabyn 1996; Mezger, Duschl, $\\&$ Zylka 1996; Melia $\\&$ Falcke 2001), with a degree of detail that cannot be accessed in any other galactic nucleus in the foreseeable future. With the now available high sensitivity and angular resolution, large ground-based telescopes offer the opportunity to obtain an unprecedented view of the Galactic centre.  Initially, with speckle imaging techniques and lately with adaptive optics techniques, high angular resolution images on the Galaxy's central cluster were obtained. This first set of observations was able to measure stellar motions on the plane of the sky, yielding estimates of the projected velocities (Eckart $\\&$ Genzel 1996; Ghez et al. 1998), projected accelerations (Ghez et al. 2000; Eckart et al. 2002a), and three-dimensional orbital motions (Eckart et al. 2002a, Sch\\\"odel et al. 2002, 2003; Ghez et al. 2003b), which each provided a successively stronger case for a super-massive black hole at the centre of the Milky Way and its association with the unusual radio source Sgr~A* (Lo et al. 1985).  More than a decade of high-resolution infrared observations of proper motions in the GC, with the ESO New Technology Telescope (NTT) and the ESO Very Large Telescope (VLT) (Eckart $\\&$ Genzel 1996; Eckart et al. \\cite{eckart02}; Sch\\\"odel et al. 2002, 2003), as well as with the Keck telescope (Ghez et al. 1998, 2000, 2003b), have revealed at least 6 stars that show substantial acceleration due to the super-massive black hole associated with Sgr A* and are on bound orbits around it. With a peri-centre of less than 0.6 mpc (15 mas) and an orbital period of $\\sim$ 15 years, S2 is the most striking case of these.  A series of observations with the NAOS/CONICA adaptive optics system/near-infrared camera at the ESO VLT unit telescope~4 that covered the peri-centre passage of the star S2 around Sgr~A* allowed Sch\\\"odel et al. (2002) to approximate a Keplerian orbit and to measure the enclosed dark mass down to a distance of $\\sim$ 0.6 mpc from Sgr~A*. With these observations, they could exclude a neutrino ball scenario (Munyaneza $\\&$ Viollier, 2002) as an alternative explanation for the dark mass concentration. They excluded as well a cluster of dark astrophysical objects (Maoz 1998) such as neutron stars, leaving a central super-massive black hole as the most probable explanation.   Using the Keck 10m telescope, Ghez et al. (2003a) measured a Keplerian orbit for the star S2 which agreed with the results of Sch\\\"odel et al. (2002). Additionally, Ghez et al. (2003a) reported the first detection of spectral absorption lines (both Br$\\gamma$ (2.1661 $\\mu$m) and He I (2.1126 $\\mu$m)), providing the first line-of-sight velocity measurements of star S2. These measurements resolved the ambiguity on the inclination of the S2 orbit indicating that its position was behind the black hole when it passed through its peri-centre. In addition, stellar rotational velocities suggest that S2 is an O8-B0 dwarf star and thus a massive ($\\sim$ 15~M$_{\\odot}$) young star ($\\le$ 10 Myrs).  From data taken with NAOS/CONICA and the new NIR integral field spectrometer SPIFFI at the ESO VLT, Eisenhauer et al. (2003) reported new astrometric observations and additional spectroscopic observations of the star S2, reducing the uncertainties on the orbital parameters. They also gave the most accurate primary distance measurement to the centre of the Milky Way of 7.94 $\\pm$ 0.42 kpc, which is in agreement with earlier determinations (see Reid 1993).  In this paper, we explore the possibility that there exists a compact, continuous mass distribution composed of several undetectable faint stars or perhaps some more exotic material in addition to the point mass of Sgr~A*. In this case, the orbit of S2 will not follow a simple Keplerian orbit, but will rather show peri-centre-shifts that result in rosetta shaped orbits.  This idea is motivated by the observation that the stellar density does not flatten out, but exhibits a steep peak towards the centre, a so-called cusp (Eckart et. al, 1995; Alexander , 1999; Genzel et al., 2003).  In contrast to earlier studies, the main approach of this work is that the mass-to-light ratio, $M/L$, is {\\it not} considered to be constant over the entire range of the GC stellar cluster. Indeed, there are indications that the stellar population varies with position and is not quite well mixed (Alexander, 1999). Furthermore, the exact composition of the cusp is still unknown and our current understanding of the stellar distribution in the GC is incomplete (Genzel et al. 2003). With current observations, low-mass stars ($K$$\\ge$ 21 mag) cannot be observed in these dense cusp regions and the true value of the M/L ratio is not known. Also, as pointed out by Baumgardt et al. (2003), dynamical evolution of a dense stellar cluster will result in a strong increase of M/L ratio by segregation of stellar evolution remnants to the centre. On the other hand, the approximation that the dynamics in the central region is Keplerian is directly related to the implicit assumption of Eckart et al. (2002a), Sch\\\"odel et al. (2002), and Ghez et al. (2003a) that the $M/L$ ratio at 2~$\\mu$m in the cusp is as low as in the outer stellar cluster ($M/L = 2M_{\\odot}/L_{\\odot}$).\\\\ Considering these circumstances, stars with short orbital periods, in particular the star S2, of which orbital data are best known, play a key role in exploring the gravitational potential. We show that the present observational data on S2 cannot discriminate between a \\emph{Keplerian} and a \\emph{non-Keplerian} potential. Subsequently, we study the influence of an extended distribution of dark mass near Sgr~A* taking into consideration the constraints set by the orbit of S2 as well as the limits set on the total enclosed mass at larger radii. In this study relativistic effects are neglected as they are second order corrections.  In \\textsection~2 we outline our modelling of the density distribution of the central region and show that the star counts near Sgr~A* can be described by a compact Plummer model core. The method used to compute orbits in non-Keplerian potentials is explained in\\textsection~3. In \\textsection~4 possible orbital models for S2 are presented. We discuss the implications of the limits on the cusp mass derived  from the \\emph{non-Keplerian} orbit modelling in \\textsection~5, and draw our conclusions in \\textsection~6. ",
        "conclusions": "In the previous sections it was demonstrated that the measured positions and line-of-sight velocities of S2 are not sufficiently well determined to exclude a Keplerian orbit. Assuming a simple Keplerian case, the central gravitational potential by definition is entirely dominated by the point mass associated with Sgr~A* and no dynamical constraints can be derived on any extended mass component due to e.g.\\ the surrounding stellar cluster. More revealing, however, is the more physical assumption of non-Keplerian orbits. The presence of a stellar cusp shows that -- at least to a certain extent -- there is some extended mass present near the black hole. This potential can be modelled by a central black hole plus an extended mass component. We assumed that the cusp has the same shape as the observed stellar number density counts which we fitted by a Plummer density model (see \\textsection{2}). We also took into account that the total enclosed mass at larger radii (1pc) would not exceed a value of $\\sim$~4.8~$\\times$~10$^6$~M$_{\\odot}$ (see section~3.1).  Having shown that the current data allow for an extended mass component higher than the one of 8540~M$_{\\odot}$ inferred from stellar number density counts which were scaled to match the velocity distribution at large radii (see \\textsection{2}), we will discuss its nature in sections \\textsection{5.1} and \\textsection{5.2}. The value of 8540~M$_{\\odot}$, representing the total mass of the 'inner cusp', is equivalent to a mass of 3100~M$_{\\odot}$ inside the core radius $R_1$ (15.5~mpc or 0.4''). It is based on the assumption that the ratio of the stellar number counts to total stellar mass does not vary with radius and environment in the Galactic Centre, i.e. that the mass-to-light ratio $M/L$(2$\\mu$m) is constant (see also Genzel et al., 2003). It is very likely that this assumption is unjustified. Genzel et al. (2003) find that the stellar population in the cusp differs to a certain degree from the population of the surrounding, large-scale cluster. Also, effects such as mass segregation and stellar collisions might work very effectively in the dense environment of the cusp and $M/L$ would show higher values towards the centre (Baumgardt et al. 2003). Therefore, we consider that the mass-to-light ratio, $M/L$ (2$\\mu$m) of the 'inner cusp', can be different from that of the outer stellar cluster. In the following, two hypotheses will be discussed: The existence of a cluster of faint low-mass stars, not yet detectable with the resolution and sensitivity of current instruments, and the existence of a cluster of heavy dark objects like stellar black holes and neutron stars.     \\subsection{Extending the K-band Luminosity Function to Faint Stars}  We use the $K$-band luminosity function (KLF) for K$\\leq$ 18 mag within a projected radius of $1.5''$ from Sgr~A* as it was determined by Genzel et al. (2003 - see their Fig.11). Note that the KLF is not corrected for extinction or for the distance modulus. The authors fitted the KLF with a power-law with a slope of $\\beta$=0.21 $\\pm$ 0.03, ($\\beta$ =dlogN/dK). Since this result is based on a number of roughly 60~stars within $1.5''$ of Sgr~A*, the observed stars themselves cannot account for a significant extended mass component. It is therefore necessary to extrapolate the KLF at the faint end. However, we assume that the observed stars trace the mass carrying population which can currently only be accessed theoretically by extrapolating the KLF towards its faint end.  For the extrapolation of the KLF two different slopes are considered, $\\beta$ = 0.21 and a steeper one, $\\beta$ = 0.35, which fits the Bulge KLF (Alexander \\& Sternberg, 1999) as well as a model of an old ($\\sim10$~Gyr) stellar cluster of solar metalicity (Zoccali et al. 2003). In the following we concentrate on the region within R$_1$ of the inner Plummer model component. The total number of stars in the inner cusp (R$\\leq 0.4''$ or 15.5~mpc) brighter than a given magnitude can be directly estimated from the (extrapolated) KLF.  The KLF gives the number of stars per surface area per magnitude. In order to calculate the total number of stars and the total stellar mass present within the spherical volume enclosed by the core radius of R$ = 0.4''$  of our central Plummer model component, we de-projected the KLF. Table~\\ref{clusters} lists the resulting numbers of stars, the total cluster mass, M$_{Cl}$, the average stellar mass in the cluster, M$_{aver}$, and the corresponding $M/L$ (2$\\mu$m) for clusters given by KLFs with slopes of $\\beta=0.21$ and $\\beta=0.35$ and for different cut-off magnitudes, between K$= 20$ (M$_{min}$ = 1M$_{\\odot}$) and $K$ = 28 (M$_{min}$ = 0.06M$_{\\odot}$). M$_{Cl}$ is calculated from the observed KLF using A$_K$=3 and the $M/L$~(2$\\mu$m) values listed in Table~\\ref{stars}.  The numbers listed in Table~\\ref{clusters} show that the mass of stars present within $R$ = 0.4'' deduced from a $\\beta=0.21$ slope KLF cannot be higher than 800 M$_{\\odot}$, even after extrapolation to the faintest magnitudes (see Table~\\ref{stars} for a list of K-magnitudes and0.1051 corresponding MS-stars at the GC).  Therefore, a cluster with a KLF of $\\beta=0.21$ cannot explain a mass of 3100 M$_{\\odot}$ within R$ = 0.4''$, estimated from direct number density counts with an $M/L$(2$\\mu$ m) = 2~M$_{\\odot}$/L$_{\\odot}$ (see \\textsection~2).   $M/L$~(2$\\mu$m) converges to a value 2~M$_{\\odot}$/L$_{\\odot}$ for $\\beta$ = 0.35, and a magnitude limit between 26 and 27, in this case we find the value of 3100 M$_{\\odot}$ required by our mass distribution modelling in \\textsection{2}. Lower magnitude limits would increase $M/L$~(2$\\mu$m) further, and the mass of the inner cusp calculated from our Plummer model would be underestimated.  Stellar types corresponding to the required M/L~(2$\\mu$m) ratios are listed in Table~\\ref{stars}.  Table~\\ref{clusters} also shows that for a main-sequence stellar population M/L(2$\\mu$m) does not exceed $4.0$~M$_{\\odot}$/L$_{\\odot}$, even if we consider very faint/low-mass stars ( $K$ = 28) and a steep faint-end KLF ($\\beta$ = 0.35). This value corresponds to a cusp representing only 0.2$\\%$ of the total mass (see Table~\\ref{T5}). We conlude that the amount of extended mass allowed inside the core radius of the 'inner cusp' would be far too large to be explained by a stellar population of MS stars, and therefore would require another type of mass carriers. We show in \\textsection~5.3 that such a configuration is possible.   \\begin{table} \\caption{Extinction free mass-to-light ratios, $M/L$(2$\\mu$m), for different stellar types in luminosity class V at the distance of the Galactic Centre.} \\begin{tabular}{||l||l||l||l||l||} \\hline M/L(2$\\mu$m)[M$_{\\odot}$/L$_{\\odot}$] & K[mag] & Mass[M$_{\\odot}$] & Spectral types \\\\ \\hline 1.29   & 16.425 & 1.60 & F0 \\\\ 1.93   & 17.005 & 1.40 & F5 \\\\ 2.73   & 17.695 & 1.05 & G0 \\\\ 5.02   & 18.665 & 0.79 & K0 \\\\ 5.36   & 18.915 & 0.67 & K5 \\\\ 11.77  & 20.065 & 0.51 & M0 \\\\ 9.23   & 20.065 & 0.4  & M2 \\\\ 16.66  & 20.915 & 0.33 & M3 \\\\ 38.49  & 22.315 & 0.21 & M5 \\\\ 87.55  & 23.815 & 0.12 & M7 \\\\ 166.43 & 25.265 & 0.06 & M8 \\\\ \\hline \\end{tabular} \\label{stars} \\end{table}  \\begin{table*} \\caption{For each 'observed' K-magnitude limit (not corrected for extinction) we give the numbers of stars $N$, cluster mass M$_{Cl}$, average stellar masses M$_{aver}$, and mass-to-luminosity ratio, $M/L$(2$\\mu$m), within the central $0.4''$ for hypothetical Plummer model type clusters with KLFs extrapolated by two power-law slopes of $\\beta=0.21$ and $\\beta=0.35$.  M$_{Cl}$ is the mass inside the core radius R$_{1}$ of the innermost Plummer model component. The derived $M/L$ values include an extinction correction for A$_K$=3.} \\begin{tabular}{||l||llll||llll||} \\hline & & & $\\beta$= 0.21 & & & &  $\\beta$=0.35 &  \\\\ \\hline Magnitude limit& N & M$_{Cl}$[M$_{\\odot}$] & M$_{aver}$[M$_{\\odot}$] & M/L(2$\\mu$m)[M$_{\\odot}$/L$_{\\odot}$] & N & M$_{Cl}$[M$_{\\odot}$] & M$_{aver}$[M$_{\\odot}$] & M/L(2$\\mu$m)[M$_{\\odot}$/L$_{\\odot}$]\\\\ K $\\leq$ 28 & 3500 &800 & 0.9 & 1.0 & 44500 & 5200 & 0.1 & 3.7 \\\\ K $\\leq$ 27 & 2200 & 710& 1.4 & 0.9 & 20000 & 3700 & 0.15 & 2.8 \\\\ K $\\leq$ 26 & 1350 & 620& 2.3 & 0.8 & 9000 & 2400 & 0.3 & 1.8 \\\\ K $\\leq$ 25 & 820  & 525& 3.7 & 0.7& 4100 & 1500 & 0.8 & 1.1  \\\\ K $\\leq$ 24 & 510  & 450& 6.0 & 0.6& 1920 & 1000  & 1.6 & 0.8 \\\\ K $\\leq$ 23 & 310  & 390& 10  & 0.5& 950 & 700 & 3.3 & 0.5  \\\\ K $\\leq$ 22 & 190  & 325& 16  & 0.4& 510 & 480 & 6.0 & 0.4 \\\\ K $\\leq$ 21 & 120  & 235& 26  & 0.4& 320 & 380 & 9.7 & 0.3 \\\\ K $\\leq$ 20 & 70   & 193 & 42  & 0.3& 230 & 270 & 11.2 & 0.2 \\\\  \\hline \\end{tabular} \\label{clusters} \\end{table*}  \\begin{table} \\caption{Mass-to-light ratios $M/L$(2$\\mu$m), for the different \\emph{offset} position fitting orbits obtained from our model calculations.  The errors on M/L are model dependent. From scaling to the enclosed mass estimate in Fig.~\\ref{fig2} we estimate an error of 5\\%. Columns 3, 4, and 5 give the cusp mass inside the orbit of S2 for core radii of respectively 20.2mpc, 15.5mpc and 13.2mpc.} \\begin{tabular}{||@{}l@{}||l||l||l||l||@{}l@{}||} \\hline M$_{tot}$         &  $f$ & M$_{S2_{apo}}$&M$_{S2_{apo}}$& M$_{S2_{apo}}$&  M/L(2$\\mu$m)\\\\ $[10^6 M_{\\odot}]$&      & R$_1$=        & R$_1$=       & R$_1$=        & [M$_{\\odot}$/L$_{\\odot}$] \\\\ & & 20.2mpc & 15.5mpc  & 13.2mpc &  \\\\ \\hline 3.65 & 0.0 & 0.001 & 0.001 & 0.001  & 2.00\\\\ 3.65 & 0.05 & 0.011 & 0.021 & 0.031 & 42.74   \\\\ 3.65 & 0.10 & 0.022 & 0.044 & 0.062 & 85.48  \\\\ 4.10 & 0.05 & 0.013 & 0.024 & 0.033 & 48.01 \\\\ 4.10 & 0.10 & 0.025 & 0.049 & 0.071 & 96.02  \\\\ 4.10 & 0.15 & 0.038 & 0.071 & 0.104 & 144.03 \\\\ 4.10 & 0.20 & 0.052 & 0.098 & 0.140 & 192.04  \\\\ 4.45 & 0.20 & 0.049 & 0.105 & 0.155 & 227.7 \\\\ 4.80 & 0.20 & 0.061 & 0.108 & 0.157 & 224.82 \\\\ 4.80 & 0.25 & 0.069 & 0.135 & 0.190  & 269.32  \\\\ 4.80 & 0.30 & 0.089 & 0.164 & 0.237  & 337.24 \\\\ \\hline \\noalign{\\smallskip} \\end{tabular} \\label{T5} \\end{table}  \\subsection{Stability of a Cluster of Low-Mass Stars}  For a multi-mass stellar distribution, high mass stellar remnants (stellar black holes and/or neutron stars) are expected to migrate to the centre as a consequence of dynamical friction.  One would expect that, within a Hubble time, these compact objects show a higher concentration toward the centre than the lighter ones (Morris et al., 1993; Miralda-Escud$\\acute e$ $\\&$ Gould, 2000), which should be transferred by that mechanisms to orbits at greater distances from the centre of the cluster. This argues against the existence of a cluster of low-mass stars in the inner cusp.  On the other hand, the possibility of a cluster of low-mass stars cannot be excluded and we are far from understanding the properties of the stars in the cusp. There is for example the unexplained presence of massive young stars, e.g. MS O/B-type stars close to the black hole (Genzel et al., 1997; Eckart, Ott \\& Genzel, 1999; Figer et al., 2000; Gezari et al., 2002; Ghez et al., 2003). These stars have not had enough time to achieve energy equipartition with the fainter older stellar population. They are hence dynamically un-relaxed. Also, there are indications for a radial anisotropy of the stars in the cusp which might be un-relaxed (Sch\\\"odel et al., 2003), in spite of the expected short relaxation time in this dense environment. Because of this general lack of theoretical understanding of the cluster near Sgr~A*, we consider that a cluster of faint/low mass stars should not be ruled out entirely from the possible interpretations of the inner cusp.   \\subsection{Is the Cusp Dominated by Dark and Massive Objects?}  In this section, we consider orbital fits with errors of $\\le$1$\\sigma$ that result in high cusp masses yet are in good agreement with the enclosed mass measurements. The evaluation presented in \\textsection~5.1 shows that such a heavy cusp is unlikely to consist of stars only. Here we study whether such a cusp could consist of heavier mass carriers like stellar black holes or neutron stars.  \\subsubsection{Presence of stellar remnants in the center due to dynamical friction} When the compact remnants of massive stars are themselves significantly more massive than the normal field stars in the Galactic Centre, as would be the case of black hole remnants, than they are susceptible to inward migration as a consequence of dynamical friction. The resulting mass segregation can lead to a pronounced concentration of compact objects within the central stellar core within a Hubble time (Morris, 1993). $N$-Body simulations of globular clusters showed that the combination of stellar evolution (production of stellar mass black holes, neutron stars, white dwarfs, and of binaries including such objects) and stellar dynamics will almost certainly lead to a strong increase of $M/L$ in the central parts of the nuclear star cluster. Black holes and neutron stars sink to the centre and may coalesce (G\\\"urkan et al. 2003). Such highly detailed self-consistent simulations of the dynamical episodes are, however, not yet possible on the scale of the Galaxy. $N$-Body simulations (even with cutting-edge special purposes computers like GRAPE-6 (Makino 2001)) cannot follow the evolution of a galactic nucleus over a Hubble time if relaxation is appreciable (Freitag \\& Benz, 2002).   Under certain assumptions and for some range of initial parameters, less realistic but more efficient Monte Carlo numerical simulations of the evolution of the GC (Freitag \\& Benz, 2002) showed that the SBHs sink to the center on a short timescale of a few gigayears, settle into a centrally concentrated distribution and dominate the stellar mass there. Similarly, Murphy Cohn $\\&$ Durisen (1991, see their figure 8b) showed that densities higher than 10$^9$ M$_{\\odot}$~pc$^{-3}$ could reside in regions as close as few mpc from the central SBH. Other studies by Morris (1993) and Miralda-Escud\\'e $\\&$ Gould (2000) estimate that 10$^4$ - 10$^5$ SBHs, due to dynamical friction, would have accumulated at distances of less than about 1 pc from the center. These latter works can only account for about 10$^2$ - 10$^4~$~M$_{\\odot}$ within separations of about 20 mpc from the centre which represents only 1$\\%$ to 10$\\%$ of our upper mass limit derived from the orbit of S2. However, the central density of the SBH cluster depends on various uncertain quantities: the SBH mass function, the stellar IMF and formation rate, the remnant progenitor masses, and the dynamical age of the GC. Morris (1993) argued that within a wide range of assumptions about the IMF, and about the minimum stellar mass capable of producing a black hole remnant, the total mass of remnants concentrated into the inner few tenth of a parsec, would be about 0.4-5 10$^6$ M$_{\\odot}$. If the black hole remnants were to achieve equipartition with the field stars, they would form an inner core with a radius of a size as small as 50~mpc. These would coalesce or form a quasi-stable cluster of stellar mass black holes. Even if a catastrophic merger of stellar remnants did occur at some point at the age of the galactic nucleus, the continuous influx of massive stellar remnants would ensure that a concentrated population of them is present within the stellar core. These findings apply, however, in the absence of an initial central black hole. To our knowledge, a similar study of a cluster with already a pre-existing super-massive black hole, as in the centre of our Galaxy, is not yet available.  \\subsubsection{In situ formation of stellar remnants} In the case of the Galactic Centre, complex dynamical episodes have taken place. A numerous variety of young early-type, bright and massive stars exist at distances of 10~mpc - 400~mpc from the centre. There exist a dozen of bright O/B main sequence stars within about 40 mpc of Sgr~A*, these are fast moving S-stars (mass $\\sim$ 20~M$_{\\odot}$) of which the star S2 is an example. There exists also $\\sim$~30 more massive (30~M$_{\\odot}$~$\\lesssim$~mass~$\\lesssim$~100~M$_{\\odot}$) very bright early-type stars, the so-called He stars exhibiting He/HI emission lines. The existance of these two types of stars at these regions from the center is still not understood. These could have formed there or would have migrated there due to different in-spiraling processes (Genzel et al. 2003; Ghez et al. 2003; Gerhard et al. 2001, Krabbe et al. 1995). If such stars were always present in the center of the Galaxy, due to stellar evolution, their remnants would contribute very efficiently to the formation of a dense cluster also at the center. Here, we assume that these stars would end their lives in the formation of neutron stars or stellar black holes with masses between 1.5~M$_{\\odot}$ and 10~M$_{\\odot}$ (on average 7~M$_{\\odot}$), we also consider their lifetimes to be $\\lesssim$~10$^7$ yrs. If after a single lifetime about 40 of such stars form stellar remnants, throughout the age of the Galaxy ($\\sim$ 10$^{10}$ yrs), it is possible to account for about 10$^5$~M$_{\\odot}$ needed to explain our upper mass limit of the 'inner cusp'.  Considering the two above described scenarios - while these do not represent the complete history of the galactic nucleus - it is, however, a fair conclusion that \\textit{strong} to \\textit{moderate} dynamically caused $M/L$ variations prevail at the GC.  Our attempts in this work is not to explain the formation of a high density cluster of stellar remnants at the GC. We would rather like to analyse such a configuration if existant. In the following we investigate if such a hypothetically high $M/L$ configuration of stellar remnants can form a stable configuration.\\\\  \\subsubsection{Stability of a cluster of stellar remnants} Rauch $\\&$ Tremaine (1996) studied the configuration of a central massive black hole plus an extended mass distribution $M$ of radius $R$ consisting of objects with mass $m$ in terms of its $non-resonant$ relaxation time $t_{rel}^{nr}$. Under the assumption that $M$ $<<$ $M_{BH}$, Rauch $\\&$ Tremaine (1996) derive how $t_{rel}^{nr}$ depends on $M_{BH}$, $M$, and the orbital time scale $t_{orb}$ at the outer edge of the cluster. If (a) the stellar orbits have random orientations and moderate eccentricities, (b) the density of stars is approximately uniform within $R$, and (c) $M_{cluster}$ $/$ $M_{BH}$ = 10$^{-2}$ - 10$^{-5}$ then  Rauch $\\&$ Tremaine (1996) find that  \\begin{equation} t^{nr}_{rel} \\sim \\frac{M_{BH}^2}{m^2N ln\\Lambda} t_{orb}, \\label{eq:6} \\end{equation}  where $ln \\Lambda$ is the Coulomb logarithm $\\sim$~13 in this case. This situation should - to first order - be applicable to the Black Hole/cusp scenario at the Galactic centre.  Condition (a) is probably fulfilled with the possible exception that the stars in the cusp might have fairly high eccentricities (Sch\\\"odel et al. 2003).  Assuming a Plummer model as a cusp description fulfills condition (b). We consider the case where a peri-centre-shift of the order of 1$^{\\circ}$ is induced. As an example, we choose the case of a cusp mass of 15\\% with a total mass of 4.1~$\\times~10^6$~M$_{\\odot}$. This cusp mass is well below the upper limit derived in section 4.1. In the case of a cusp mass of 15\\%, $M/M_{BH}$ is of the order of 10$^{-2}$, if we restrict ourselves to the region within the core radius $R_{1}$.  This is close to what is required by condition (c).  Here, $M$$_{BH}$ = 3.5 $\\times$ 10$^6$ and $m_{sr}$ is the average mass of the stellar remnants. A value of $m_{sr}$ = 5 M$_{\\odot}$ is roughly consistent with a composite cluster made of neutron stars ($m$ $\\sim$ 1.5~M$_{\\odot}$) and stellar black holes (5~M$_{\\odot}$ $<$ m $<$ 25~M$_{\\odot}$). We then find that $t_{rel}^{nr}$ $\\sim$ 10$^6$ $\\times$ t$_{orb}$, if we assume that most of the mass inside the core radius is present in the form of stellar remnants of average mass $m_{sr}$. Considering that the core radius of  15-20~mpc will define t$_{orb}$ we find $t_{rel}^{nr}$ to be about $2\\times 10^7$ years. \\\\ \\\\ We can also investigate whether such a configuration is stable by estimating how many stellar black holes evaporate. Integrating a Maxwell distribution function for the velocities above the escape velocities gives the percentage of stars not bound to the BH. For the velocitiy dispersion, following Alexander et al. (2003), we can write: \\begin{equation} v_{escape}=  \\sqrt{2}v_{circular} = \\sqrt{2(1 + \\alpha)}~\\sigma . \\label{eq:7} \\end{equation}  For a value of $\\alpha$ = 2, about 0.03\\% of the present stellar black holes will be evaporated after each relaxation time.  For a steeper cusp with $\\alpha$=3 we find that 0.006\\% of the stellar black holes will evaporate on that time scale.  For $t_{rel}^{nr}$ of a few 10$^7$ years and $\\alpha$ between 2 and 3 about 50\\% of the stellar black holes will have evaporated after 25 to 250 $t_{rel}^{nr}$ corresponding to about $5\\times10^8$ to $5\\times10^{9}$ years.  Given the fact that there will also be an influx of mass i.e. stellar black holes or neutron stars from outside the cusp one can consider such a configuration stable over a significant fraction of the Milky Way's age.  Given equation (\\ref{eq:6}), for a fixed $t_{orb}$, the relaxation time decreases linearly with $N$, i.e with the cusp mass. Thus, cusp masses well exceeding 2~$\\times~10^5$~M$_{\\odot}$ would not form a stable configuration compared to the age of the Milky Way."
    },
    "0402/astro-ph0402500_arXiv.txt": {
        "abstract": "In this paper we present a new, essentially empirical, model for the relation between the mass of a dark matter halo/subhalo and the luminosity of a galaxy hosted in it.  To estimate this, we replace the assumption of linearity between light and mass fluctuations with the assumption of monotonicity between galaxy light and halo or subhalo mass. We are enabled to proceed with this less restrictive ansatz by the availability of new, very high resolution dark matter simulations and more detailed and comprehensive global galactic luminosity functions.  We find that the relation between halo/subhalo mass and hosted  galaxy luminosity, is fairly well fit by a double power law. That between halo mass and group luminosity has a shallower slope for an intermediate mass region, and is fairly well fit by a two branch function, with both branches double power laws. Both relations asymptote to $L\\propto M^4$ at low $M$,  while at high mass the former follows $L\\propto M^{0.28}$ and the latter  $L\\propto M^{0.9}$.  In addition to the mass-luminosity relation, we also derive results for the occupation number, luminosity function of cluster galaxies, group luminosity function and multiplicity function. Then, using a prescription for the mass function of haloes in under/overdense regions and some further assumptions on the form of the mass density distribution function, we further derive results for biasing between mass and light and mass and galaxy number, light distribution function and the void probability distribution.  Our results for the most part seem to match well with observations and previous expectations. We feel this is a potentially powerful way of modelling the relation between halo mass and galaxy luminosity, since the main inputs are readily testable against dark matter simulation results and galaxy surveys, and the outputs are free from the uncertainties of physically modelling galaxy formation. ",
        "introduction": "In recent years, N-body numerical simulations have given us a good understanding of dark matter structure for standard cosmological scenarios, while large scale observational surveys have done the same for the distribution of galaxies. In this way, we are now developing a good picture of how mass and luminosity are distributed in the universe. However, it is still not well known how to connect the two pictures. While it is well established that dark matter haloes are the hosts of the observed galaxies, it is still poorly understood how the former are related to the latter. Further, the picture is complicated by the fact that what is usually taken as a halo in simulations would often host multiple galaxies, especially for higher masses. To analyse the issue fully, it is necessary to look at the halo substructure, since each subhalo can host a galaxy.  Establishing such a link between halo mass and galaxy luminosity would be important because, first of all, it would allow us to have a direct connection between theory and observation, dark matter haloes and galaxies. Further, it could also shed some light into the theory of galaxy formation.  There are several ways in which this problem can be studied. The more direct ones involve either numerical simulations including gas dynamics \\citep{whs,ytj,pjf,nfco,Berlindetal}, or semi-analytical models of galaxy formation \\citep{kns,gbf,kcda,kcdb,bba,bbb,sd,sls,wsb,bfb,Berlindetal}, but, while they explicitly give the properties of galaxies located in a given halo, they have the added difficulty that many of the mechanisms involved in galaxy formation are poorly understood, and difficult to compute. Their complexity could also mask any fundamental relations that might be present between halo and galaxy properties.  More indirect approaches have also been studied. The halo occupation distribution (HOD) model \\citep{seljak,benson,bws,zt,BerlindWeinberg,Berlindetal,mp} is based on the probability $P(N|M)$ that a halo of mass M is host to N galaxies.  By specifying the $P(N|M)$ function, along with some form for the distribution of dark matter and galaxies within each halo, it is then possible to relate different statistical indicators of the dark matter and galaxy distributions, such as correlation functions, to each other. This fully specifies the bias between the galaxy and the underlying matter distributions.  A recent paper by \\citet{Kravtsovetal}, has done a detailed study of results from simulations and related them to the HOD model, and has concluded that the form of $P_s(N_s|\\mu)$ for the subhaloes is approximately universal, where $\\mu$ is the subhalo mass scaled to an appropriate minimum mass.  This paper also give results for the relation between galaxy absolute magnitude and halo circular velocity.  Other work \\citep{BoschYangMo,YangMoBosch,Moetal} has taken this approach one step further by studying not only the number of galaxies associated with each halo, but also their luminosity, by building the conditional luminosity function, $\\Phi(L|M){\\rm d}L$. This gives the number of galaxies with luminosities in the range $L\\pm{\\rm d}L/2$ contained in a halo of mass $M$. While this work directly relates the halo mass to the galaxy luminosity, it does so only to the average values, lacking the full statistical treatment which is analysed in the HOD models.  Other authors have used a slightly different method. Instead of trying to specify the number of galaxies in each halo, they treat the halo as a whole and identify it with a galaxy group. Then, by comparing the group luminosity function with the halo mass function, they obtain the luminosity associated with each halo \\citep{PeacockSmith,MarinoniHudson}, and also develop ways to estimate the number of galaxies hosted in a halo, thus coming back partly to the $P(N|M)$ estimate of the HOD models.  In the present paper, we follow a new and conceptually clear approach based on one simplified and testable hypothesis: there is a one to one, monotonic correspondence between halo/subhalo mass and resident galaxy luminosity. We might call this the empirical (rather than the semi-analytical) approach because there is no attempt to physically model the galaxy formation process. Instead we take from observations the galaxy luminosity distribution and match it with the theoretical halo/subhalo distribution. This has the additional advantage of naturally giving a lower mass threshold for haloes that host luminous galaxies, as the luminosity decreases sharply with mass for less massive haloes. It also implicitly gives rise to galaxy systems, if one identifies a system of a massive halo and its subhaloes with the central galaxy and its satellites in groups and clusters. We find that the single assumption is very powerful and allows us to compute, and compare to observations many quantities from bias to the void distribution function to the spatial correlation function.  This paper is organized as follows: in section 2 we present our model for the subhalo mass distribution in a parent halo, and build the global subhalo mass function. In section 3, we derive the relation between mass and luminosity, as well as some other functions such as the luminosity function of cluster galaxies, the group luminosity function and the multiplicity function. In section 4 we study how to apply the relation we obtain to get the light density and the number density of galaxies as a function of mass density, and also obtain results for the distribution function of light density and the void probability function. Finally, we conclude in section 5.  Throughout we have used a concordance cosmological model, with $\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$, $h=0.7$ and $\\sigma_8=0.9$ \\citep{triangle,wmap}. ",
        "conclusions": "In this paper we present a new model for relating halo mass to hosted galaxy luminosity, based on the dark matter substructure.  We feel this is a potentially powerful way to approach this problem, since it is based on two main inputs, the halo/subhalo distribution and the galaxy luminosity function, which can be well tested and adjusted to results from simulations and surveys, respectively.  Additionally, the model requires only one further assumption: that there is a one to one and monotonic relation between imbedded halo/subhalo mass and galaxy luminosity. This is more explicit but in general agreement with the general assumptions made in past studies and general assumptions made in similar work of the same subject. It is a much less restrictive assumption than the still sometimes utilized ansatz of a linear ``bias'' between galaxy numbers and dark matter density.  We have shown how, starting with a prescription we describe for the subhalo mass distribution in a parent halo, it is possible to obtain a relation for the luminosity of a hosted galaxy, as well as the group luminosity when the system of a halo and its subhaloes is identified with groups or clusters of galaxies. The subsequent results appear to match well with general assumptions of the behaviour of such a relation, as well as results for the mass and luminosity at scales of dwarf galaxies, $L^*$ galaxies and massive clusters.  From this model, it is then possible to derive many quantities that can be directly compared with further simulation or observational results. We have shown four examples of this, namely the occupation number, the luminosity function of cluster galaxies, the group luminosity function and the multiplicity function.  Using further assumptions on how to calculate the mass function for regions of different average densities, and on the shape of the dark matter distribution function, we have also obtained a relation between mass and light and number densities for different smoothing lengths close to what is expected from previous bias studies. We also obtain the distribution function of light density, and the void probability function. The latter is a powerful additional test, since it probes a highly non-linear regime, and our results seem to match well with previous observations.  The major difficulty with the model as it is presented here has to do with the identification of the mass of the subhaloes. To be able to apply the monotonic correspondence to the galaxy luminosity, the original mass in the subhaloes has to be taken into account. However, the mass distribution we use is measured for the stripped mass of the subhaloes in their parent halo. In order to build the model, we took the approximation of taking an average for the stripping factor. This is, however, a not wholly satisfactory approach, since in general the stripping history will be highly variable from subhalo to subhalo, and also in different parent haloes. Ideally, we would like to use the maximum circular velocity instead of the mass to identify the subhaloes, since this quantity should less sensitive to stripping for the relatively massive subhaloes which can host galaxies. Unfortunately, we do not have at present a good distribution for subhalo abundance as a function of the maximum circular velocity, although we wish to study this further in the future. Nonetheless, although somewhat crude, the approximation we took is reasonable, and since this work deals with statistical averages for most of the quantities, it does fit somewhat well into the whole structure. More importantly, the results we obtain seem for the most part to match well with observational data. In future work which takes into account the statistical variation, however, this factor will certainly be of great importance."
    },
    "0402/astro-ph0402446_arXiv.txt": {
        "abstract": "We have performed high resolution 3D simulations of the Local Bubble (with 1.25 pc finest resolution) in a \\emph{realistic background ISM}, jointly with the dynamical evolution of the neighbouring Loop I superbubble. We can reproduce (i) the size of the bubbles (in contrast to similarity solutions), (ii) the interaction shell with Loop I, discovered with ROSAT, (iii) predict the merging of the two bubbles in about 3 Myr, when the interaction shell starts to fragment, and, (iv) the generation of blobs like the Local Cloud as a consequence of a dynamical instability. ",
        "introduction": "The Local Bubble (LB) is a cavity, elongated towards the Galactic North Pole, with an average extension of about 100 pc, copiously radiating in soft X-rays. There exist a number of discrepancies between observations and modelling that have all been outlined in a recent panel discussion (Breitschwerdt \\& Cox 2004). Most models proceed from a multi-supernova origin of the LB (cf.\\ Bergh\\\"ofer \\& Breitschwerdt 2002 for possible supernova (SN) progenitors), in which the LB is the result of successive SN explosions, although part of the soft X-ray emission could be of heliospheric origin, generated by charge exchange reactions between solar wind ions and heliospheric plasma. ",
        "conclusions": "We use a 3D Godunov type parallelized AMR hydrocode to track small scale structures down to 1.25 pc, where necessary (for details see Avillez 2000). According to Bergh\\\"ofer \\& Breitschwerdt (2002) 20 stars of Pleiades subgroup B1 with masses between 10 and 20 ${\\rm M}_\\odot$ were moving through the LB within the last 14 Myr, exploding after their main sequence life time, $\\tau_{\\rm ms}$, along a path crossing $x=175$, $y=400$ pc (see Fig.~1), thus generating the Local Cavity into which the LB expands. The Galactic SN rate has been used for the setup of {\\em other SNe} in the background disk. \\begin{figure}[t] \\centering \\vspace*{1.in} Left panel: breitsch$\\_$fig1a \\hspace*{0.5cm} Right panel: breitsch$\\_$fig1b \\vspace*{1.in}  \\caption{\\emph{Left:} Temperature map (Galactic plane cut) of a 3D LB simulation at present (i.e.\\ 14.4 Myr after first explosion) with the LB centered at (175, 400) pc and Loop I 200 pc to the right. \\emph{Right:} Same, but at $t=17$ Myrs, showing fragmentation of the interaction shell and formation of cloudlets by hydrodynamic instabilities. \\label{fig1}} \\end{figure} First we derive analytic similarity solutions, taking an initial mass function for Galactic OB associations with powerlaw index $\\Gamma=-1.1$, and using $\\tau_{\\rm ms} = 3\\,10^7 \\, (m/[10 {\\rm M}_\\odot])^{-\\alpha}$ yr (Stothers 1972), with $\\alpha = 1.6$, for stars within the mass range $7\\, {\\rm M}_\\odot \\leq m \\leq 30 \\, {\\rm M}_\\odot$. The mechanical energy input rate (see Bergh\\\"ofer \\& Breitschwerdt 2002) then is $L_{\\rm SB} = L_0 \\, t_7^{\\delta}$, where $L_0 = 4.085 \\times 10^{37} \\,{\\rm erg/s}$, $\\delta = -(\\Gamma/\\alpha + 1) = -0.3125$ and $t_7 = t/10^7$ yr. Modifying suitably the similarity solutions of McCray \\& Kafatos (1987), we obtain $R_b = A t^\\mu = 251 \\left(2 \\times 10^{-24} {\\rm g}/{\\rm cm}^3 /\\rho_0\\right)^{1/5} t_7^{\\mu} \\, {\\rm pc}$ with $\\mu = (2-\\Gamma/\\alpha)/5=0.5375$. Note, that our value of $\\mu$ is between the canonical value of $0.4$ for a Sedov-type solution and $0.6$ for a wind type solution, due to the declining energy input rate with time. If we now try to match a present day average radius of the LB of even 146 pc at time $t_{\\rm dyn}=14.4$ Myr, we need a constant ambient density of $n_0\\simeq 40 \\, {\\rm cm}^{-3}$, a value way above the average ISM density. Comparison with our excellently matching numerical results shows, however, that this discrepancy must be due to mass loading, and turbulence in a SN disturbed \\emph{background} medium."
    },
    "0402/astro-ph0402670_arXiv.txt": {
        "abstract": "We investigate the effects that starspots have on the light curves of eclipsing binaries and in particular how they may affect the accurate measurement of eclipse timings. Concentrating on systems containing a low-mass main-sequence star and a white dwarf, we find that if starspots exhibit the Wilson depression they can alter the times of primary eclipse ingress and egress by several seconds for typical binary parameters and starspot depressions. In addition, we find that the effect on the eclipse ingress/egress times becomes more profound for lower orbital inclinations. We show how it is possible, in principle, to determine estimates of both the binary inclination and depth of the Wilson depression from light curve analysis  The effect of depressed starspots on the O--C diagrams of eclipsing systems is also investigated. It is found that the presence of starspots will introduce a `jitter' in the O--C residuals and can cause spurious orbital period changes to be observed. Despite this, we show that the period can still be accurately determined even for heavily spotted systems. ",
        "introduction": "\\label{sec:wilsonintro}  The study of eclipsing binary stars provides our best source of accurate stellar parameters, such as masses and radii, which are essential to theories of stellar structure and evolution. One can also use eclipsing systems to study, in great detail, the orbital period evolution of binary stars -- which in turn can be used to test theories of binary star evolution.  Measurements of the orbital period can be made by timing the eclipse ingress and egress, and then determining a standard time marker ({\\rm e.g.} the time of mid-eclipse) from which a linear ephemeris can be calculated. The method of searching for orbital period changes then amounts to comparing observed time markers with those predicted from the ephemeris using a standard O--C analysis, and then looking for any systematic differences. These measurements can be especially accurate for close binary systems where a compact object such as a white dwarf is eclipsed, as these provide sharp eclipse transitions. Furthermore, with the advent of ultra-fast CCD cameras such as {\\sc ultracam} (see \\citealt{dhillon01b}; \\citealt{dhillon04}), timings of eclipse ingress/egress events in these objects have been measured to an accuracy of $\\sim$0.1 seconds.  In many of the short-period binaries containing a low-mass main-sequence star (the secondary star) eclipsing a compact object (the primary star), the secondary star is either known to show high levels of magnetic activity ({\\rm e.g.} the pre-cataclysmic variable V471 Tau -- see \\citealt{ramseyer95}) or has all the necessary prerequisites for a magnetically active star (rapid rotation coupled with a convective envelope). TiO studies by \\citet{oneal98} have shown that the spot coverage on active stars can exceed 50 per cent of the stellar surface. Combine this with the fact that observations of sunspots show that they are depressed below the surrounding photosphere by 100's of kilometres (the so-called `Wilson depression') then we may well expect the surfaces of such stars to appear heavily pitted. Therefore, it is plausible that the appearance of a starspot on the limb of the secondary star as it occults the primary could then affect the time of eclipse ingress or egress and hence introduce spurious orbital period changes.  In this paper we show how the presence of starspots and the Wilson depression can influence eclipse light curves. We start in Section~\\ref{sec:model} with a description of the model that was used to investigate the effect, and present simulated eclipse light curves in Section~\\ref{sec:results} which demonstrate its magnitude. In Section~\\ref{sec:observe} we investigate the consequences that depressed starspots may have on O--C analysis.  Finally, in Section~\\ref{sec:conclusions} we discuss the implications that our results have. ",
        "conclusions": "\\label{sec:conclusions}  In Section~\\ref{sec:results} we showed that depressed starspots can delay ingress or advance egress in eclipsing main-sequence white-dwarf binaries by several seconds, and that the magnitude of this effect becomes stronger for lower inclinations. We have also shown that it is, in principle at least, possible to determine both the inclination of the system and the depth of the Wilson depression from light curve analysis.  Depressed starspots also cause a `jitter' in the residuals of O--C diagrams, which can also result in the false detection of spurious orbital period changes. Such changes due to starspots would be distinguishable from other mechanisms that cause period changes, and it should still be possible to determine the orbital period accurately.  Observational evidence for such effects are of interest. First, it would confirm that starspots also show the Wilson depression exhibited by sunspots -- providing further tests of the solar-stellar connection.  Second, it is not entirely clear whether the giant starspots that are found in Doppler images of rapidly rotating stars are actually monolithic, or groups of far smaller, individually unresolved, starspots. The effects described here will only be produced by relatively large starspots, and observations of these effects provide one of the few ways in which it would be possible to determine the monolithic nature (or otherwise) of starspots.  The best candidates in which to see the effects of Wilson--depressed starspots are eclipsing detached white--dwarf/late M--dwarf binaries, where the Applegate mechanism will be weak. However, even where the Applegate mechanism is strong, Wilson--depressed starspots would still cause variations in the eclipse width to be seen, which would otherwise remain constant under simple period changes.  Finally, \\cite{kalimeris02} have shown that intensity variations due to starspots (not taking into account Wilson depressions) can introduce disturbances of up to $\\sim$0.01 days in the O--C residuals of contact binaries. Given the rapid evolutionary timescales of spots (of the order of days) seen in recent Doppler images of the contact binary AE Phe (\\citealt{barnes04}), this may lead to explaining some of the observed jitter in the O--C curves of these objects. Since, in this work, the effects of a Wilson depression seem to result in a scatter of only a few seconds in the O--C residuals, it seems unlikely that the Wilson depression will be a significant source of jitter for contact binaries. However, given the different geometries involved and procedures for determining standard time-markers, a proper assessment of whether this latter statement holds true is beyond the scope of this paper."
    },
    "0402/astro-ph0402350_arXiv.txt": {
        "abstract": "We present the results of linear analysis and two-dimensional local magnetohydrodynamic (MHD) simulations of the Parker instability, including the effects of cosmic rays (CRs), in magnetized gas disks (galactic disks). As an unperturbed state for both the linear analysis and the MHD simulations, we adopted an equilibrium model of a magnetized two-temperature layered disk with constant gravitational acceleration parallel to the normal of the disk. The disk comprises a thermal gas, cosmic rays and a magnetic field perpendicular to the gravitational accelerartion. Cosmic ray diffusion along the magnetic field is considered; cross field-line diffusion is supposed to be small and is ignored. We investigated two cases in our simulations. In the mechanical perturbation case we add a velocity perturbation parallel to the magnetic field lines, while in the explosional perturbation case we add cosmic ray energy into a sphere where the cosmic rays are injected. Linear analysis shows that the growth rate of the Parker instability becomes smaller if the coupling between the CR and the gas is stronger (i.e., the CR diffusion coefficient is smaller). Our MHD simulations of the mechanical perturbation confirm this result. We show that the falling matter is impeded by the CR pressure gradient, this causes a decrease in the growth rate. In the explosional perturbation case, the growth of the magnetic loop is faster when the coupling is stronger in the early stage. However, in the later stage the behavior of the growth rate becomes similar to the mechanical perturbation case. ",
        "introduction": "It has been suggested that magnetic fields may play important roles for active phenomena in space, for example in astrophysical jets \\citep[e.g.,][]{shibata85,matsumoto96}, solar activities \\citep[e.g.,][]{priest82,yokoyama01}, and active galaxies \\citep[e.g.,][]{kuwabara00}. With such active phenomena, if a gas layer is supported by the horizontal magnetic fields against gravity, then the Parker instability may appear and can play an important role. Magnetic fields are also thought to play an important role in accretion disks \\citep[e.g.,][]{stella84,kato86} and galactic disks. For example, in galactic disks, the interstellar matter (ISM) is aggregated and grows to giant cloud complexes in spiral arms of galaxies via the Parker instability \\citep[e.g.,][]{mouschovias,mouschovias74,elmegreen82a,elmegreen82b, elmegreen86}. On the other hand, cosmic rays (CRs) may also play an essential role in the dynamics of the ISM, since it is recognized that the energy density of cosmic rays is of the same order as that of the magnetic field and turbulent gas motions \\citep{parker69,ginzburg76,ferriere01}. The importance of the effects of cosmic rays has been acknowledged, and a discussion concerning CRs was also given in the original work on the Parker instability \\citep{parker66}. For studying the dynamics of CRs, there are several approaches. The particle-particle approach treats the plasma and the CRs as particles; the fluid-particle approach treats the plasma as a fluid and the CRs as particles; and the fluid-fluid approach treats the plasma and the CRs as fluids. The hydrodynamical approach to the Parker instability and Parker-Jeans instability without CRs has been done by nonlinear calculation \\citep[e.g.,][]{matsumoto88,shibata89,chou00,kim02}. On the other hand, in spite of the suggestions of many astrophysical applications, there are very few papers on the evolution of the Parker instability with the effects of cosmic rays \\citep{hanasz97,hanasz00}. \\citet{hanasz00} carried out calculations on the Parker instability induced by cosmic ray injection from a supernova under the thin-flux-tube approximation, and suggested that the instability grows on shorter timescales, with the values of the diffusion coefficient decreasing. As the diffusion coefficient decreases, the diffusion speed of the CRs decreases. Since the region where the CR energy is injected keep it for a long time, the dynamics is dominated by the interactions near the injection region.   In this paper, we present the results of a linear analysis and MHD simulation of the Parker instability with the effects of cosmic rays, starting from an equilibrium two temperature layered disk. We adopt the hydrodynamic approach for cosmic ray propagation \\citep{drury81,ko92}. The paper is organized as follows. In \\S\\,2, we present our physical assumptions and the equilibrium model. Linear stability analysis of the system is given in \\S\\,3, and the numerical results are described in \\S\\,4. \\S\\,5 provides a summary and discussion. ",
        "conclusions": "Using linear analysis and a time-dependent non-linear calculation, we studied the Parker instability (or magnetic buoyancy instability) with the effect of cosmic rays. Several works on the linear analysis of the Parker instability with the effect of CRs have been published \\citep[e.g.][]{hanasz97.1,ryu03}. In \\citet{hanasz97} the CR energy equation (including diffusion) was not solved. In \\citet{ryu03} the effect of rotation is not included, and only two cases of CR pressure were described (one was without CR $\\beta=0$, and the other was with the same unperturbed CR and gas pressures $\\beta=1$). Since \\citet{ryu03} showed that the effect of cross-field-line diffusion of CRs is negligible in the context of ISM, we neglected the effect of cross-field-line diffusion in our analysis for simplicity.   In the linear analysis, the growth rate becomes larger as the CR diffusion coefficient along the field line $\\kappa_{\\|}$ increases, and is saturated at large $\\kappa_{\\|}$. This result is consistent with the result by \\citet{ryu03}. The growth rate also becomes larger when the initial ratio of the CR pressure to gas pressure $\\beta$ increases, and is saturated at large $\\beta$. This is consistent with \\citet{ryu03}, except for some slight differences. In our results, the maximum growth rate of the normal Parker case ($\\beta=0$) is almost half of the fiducial case ($\\beta=1$), and the critical wave number, over which the instability is stabilized, of the normal Parker case is about $0.7$ times of the fiducial case. In \\citet{ryu03}, the maximum growth rate of the normal Parker case is less than half of the $\\beta=1$ case, and the critical wave number of the normal Parker case is about half of the $\\beta=1$ case. The differences perhaps come from how the normalization was taken. In fact, we succeeded in producing their results by taking the same scale height under the same equilibrium condition. The scale height of the disk, $H=(1+\\alpha+\\beta)C_{\\rm s0}^2/(\\gamma_{\\rm g}g_z)$, changes with the values of $\\alpha$ and $\\beta$, and it takes the value $H=2C_{\\rm s0}^2/(\\gamma_{\\rm g}g_z)$ in the normal Parker case ($\\alpha=1$, $\\beta=0$), and $H=3C_{\\rm s0}^2/(\\gamma_{\\rm g}g_z)$ in the fiducial case ($\\alpha=\\beta=1$). We allowed the change of the scale height because we preferred not to change the gravitational acceleration. This is the reason why we took the unit of length as $H_0=C_{\\rm s0}^2/(\\gamma_{\\rm g}g_z)$. The effect of the rotation stabilizing the Parker instability is similar to the case without CRs. The $k_{x{\\rm max}}$ increases as the $\\Omega$ increases. Our result for the effect of rotation is consistent with that by \\citet{hanasz97.1}, except for the difference in the region of small wave number. The growth rate with rotation is small near $k_x=0$. It increases linearly with the wave number in \\citet{hanasz97.1}. However, in our result the growth rate increases faster than linear, and this is also observed in the normal Parker case \\citep[see][chap.~17]{kato}.  In the MHD simulation, we showed that the growth rate of the instability becomes smaller when the diffusion coefficient $\\kappa_{\\|}$ becomes smaller, which agrees well with the result of the linear analysis. At the end of the linear growth, the morphology of the magnetic loop developed from the initial perturbation is more or less the same in the three models ($\\kappa_{\\|}=200$, $\\kappa_{\\|}=40$, and $\\kappa_{\\|}=10$) studied in the mechanical perturbation case. However, in the non-linear stage, the magnetic loop in different models develops into different morphologies. From the distribution of CR pressure, density, and velocity along a magnetic field line at the end of linear growth, we found several characteristics. In the case of small diffusion coefficient (e.g., $\\kappa_{\\|}=10$, i.e., the coupling between the CRs and the gas is strong), the CR pressure distribution is rather non-uniform. CRs tend to accumulate near the foot point of the magnetic loop, and the CR pressure gradient force towards the top of the loop becomes larger. The falling motion of matter is then impeded by the CR pressure gradient force, and the growth rate of the Parker instability decreases. In the case of a large diffusion coefficient (e.g., $\\kappa_{\\|}=200$, $\\kappa_{\\|}=40$), the falling speed of matter along a magnetic field line exceeds the speed of sound, and a shock is formed near the foot point of the magnetic loop. Moreover, the CR pressure distribution along a magnetic field line in the cases of large diffusion coefficient (e.g., $\\kappa_{\\|}=200$) reminds us of the profile of CR pressure in cosmic-ray-modified shocks \\citep[e.g.,][]{drury81,ko97}. The linear growth rate in the simulations agrees well with that in the linear analysis. We also found that the speed along the disk is saturated at the initial Alfv\\'en speed. This result agrees with that in the normal Parker instability \\citep[i.e., without CRs,][]{matsumoto88}. The unperturbed state has the scale height $H=(C_{\\rm s}^2+\\beta C_{\\rm s}^2+V_{\\rm A}^2/2)/g_z$. When the Parker instability takes place, the gas falls down along the magnetic field lines. The CR pressure tends to distribute uniformly along a magnetic field line (at least in the case of large diffusion coefficient), and its contribution to the scale height disappears. Thus the scale height along the magnetic field lines settles to $H'=C_{\\rm s}^2/g_z$ at later times. The released gravitational energy in the form of kinetic energy per unit mass is estimated as $V_{\\rm A}^2/2$. Hence, we obtain the same results as the normal Parker case, even when the effect of CRs is included.  The explosional perturbation case has been studied by \\citet{hanasz00}. They stated that the smaller the diffusion coefficient the larger the growth rate of the instability. This trend is the opposite of what we found from linear analysis and simulation in the mechanical perturbation case. We thus computed the explosional perturbation case for a longer time. Our result showed that the growth rate is larger in the smaller diffusion coefficient model only in the early stage. The growth rate becomes smaller when compare with that of the large diffusion coefficient model in the later stage. The growth of instability is suspended by the CR pressure gradient force interfering with the falling motion of the matter in the small $\\kappa_{\\|}$ model, while the magnetic loop can grow up to larger scales in the large $\\kappa_{\\|}$ model."
    },
    "0402/astro-ph0402399_arXiv.txt": {
        "abstract": " ",
        "introduction": "The CMB experiments of the recent years provided a solid background for understanding the structure and early evolution of the Universe.  The comparison of the results of independent experiments has led to satisfactory mutual agreement. Excellent statistical agreement was found by comparing measurements of the angular power spectrum of the anisotropy detected by different experiments. Moreover, the maps of the BOOMERanG \\cite{Bern1}, \\cite{Bern2}, \\cite{Nett2}, \\cite{Ruhl}, MAXIMA and ARCHEOPS \\cite{benoit} experiments have been recently compared \"pixel to pixel\" to the high quality maps from the WMAP satellite and excellent agreement was found \\cite{Bern3}, \\cite{abroe}.  In the present article we will study the WMAP map \\cite{Ben03} for the region coinciding with the BOOMERanG one, to study the distortion found in 150 GHz Boomerang maps. The latter showed ellipticity for hot and cold anisotropy areas (spots) about 2.2 (2.5 for several degree areas) invariant to the temperature threshold within its certain interval \\cite{ellipse}. The effect was detected both for small areas, i.e. containing from several to 10 pixels, but also for larger ones, with over 100 pixels. The latter spots are larger than the scale of the horizon at the last scattering surface. The biases of the estimator and of the noise are negligible for the larger areas \\cite{ellipse1}. If this ellipticity effect, first detected for the COBE-DMR data \\cite{GT}, is due to the geodesic mixing \\cite{GK1}, then it might indicate a non precisely zero curvature of the Universe, and hence be related with the origin of the low CMB multipoles anomaly detected by WMAP (see e.g. \\cite{Aur03,Efs03,Ell03,Lum03}). The study below of the ellipticity in the WMAP maps, on the same region observed by BOOMERanG, acts also as an additional comparison of the results of the two experiments.  The geometry of the excursion sets of random fields including the the ellipticity for Gaussian fluctuations have been studied before, and predict average ellipticity  1.4, with decrease towards higher thresholds (see \\cite{Adler,Bond}).  Measuring the ellipticity in the CMB maps actually implies the estimation of the Kolmogorov-Sinai (KS) entropy of the dynamical process which might lead to that effect (the geodesic mixing or whatever). This is based on the essential fact that, KS-entropy being local (in time) characteristics of the dynamical system, enable to determine the properties of the evolution of the system. With the CMB we have an analogous problem: having the maps, i.e. the parameters of the moving photon beams of various temperature at present epoch, we aim to recover their former history. ",
        "conclusions": "The analysis of WMAP maps in the same region observed by BOOMERanG detects ellipticity of anisotropies of the same average value (around 2), as  found for BOOMERanG, even though there is difference in the maps at least due to the absence of low multipoles in the BOOMERanG data. Ellipticity for large scale areas had been found also in COBE-DMR maps \\cite{GT}. The WMAP data confirm the effect for scales both smaller and larger than the horizon at the last scattering surface. This suggests that the effect is not due to physical effects at the last scattering surface, and can arise after, while the photons are moving freely in the Universe. As a large-scale effect this can be related to the WMAP low multipole anomaly, since the geodesics mixing and the low multipoles are both related to the diameter of the Universe - the first one via hyperbolic geometry, the second one, via boundary conditions - and possibly even with vacuum's relevant modes' contribution to the dark energy \\cite{Lambda}."
    },
    "0402/astro-ph0402166_arXiv.txt": {
        "abstract": "We present an overview of our ongoing systematic search for wide (sub)stellar companions around the stars known to host rad-vel planets. By using a relatively large field of view and going very deep our survey can find all directly detectable stellar and brown dwarf companions (m$>$40\\,$M_{Jup}$) within a 1000\\,AU orbit. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0402/astro-ph0402216_arXiv.txt": {
        "abstract": "{Recent XMM-Newton observations of clusters of galaxies have provided detailed information on the distribution of heavy elements in the central regions of clusters with cooling cores providing strong evidence that most of these metals come from recent SN type Ia. In this paper we compile information on the cumulative mass profiles of iron, the most important metallicity tracer. We find that long enrichment times ($\\ge 5$ Gyr) are necessary to produce the central abundance peaks. Classical cooling flows, a strongly convective intracluster medium, and a complete metal mixing by cluster mergers would destroy the observed abundance peaks too rapidly. Thus the observations set strong constraints on cluster evolution models requiring that the cooling cores in clusters are preserved over very long times. We further conclude from the observations that the innermost part of the intracluster medium is most probably dominated by gas originating predominantly from stellar mass loss of the cD galaxy.}  \\authorrunning{B\\\"ohringer et al.} \\titlerunning{Implications of Central Metal Abundance Peaks in Galaxy Clusters} ",
        "introduction": "Since the discovery of the iron emission line in the thermal spectra of the intracluster Medium (ICM) in clusters of galaxies (Mitchell et al. 1976), the abundances of the heavy elements in the ICM received much attention and provided insights into the star formation history in the cluster volumina. The large amount of iron implied that the cluster galaxies must have lost a large fraction of the iron produced during their star formation histories to the ICM. The total iron mass in the ICM is in fact larger than the total iron mass in all cluster galaxies. Attempts to model the production of the observed iron masses on the basis of the observed present day stellar population have shown that such models have to make maximizing assumptions like a top heavy IMF in early epochs of star formation to increase the number of historic SN II (Arnaud et al. 1992, Elbaz et al.  1993) or an increased rate of type Ia supernovae in the past (Renzini et al. 1993).  In this paper we concentrate on the implications of the heavy element abundances in the central regions. With the ASCA satellite observatory which provided the first possibility of spatially resolved spectroscopy of the cluster ICM it was found that some clusters with cD galaxies had a strong increase in the metal abundance towards the center (Fukazawa et al. 1996, 2000, Matsumoto et al. 1996, Matsushita et al. 1998, Ikebe et al. 1999). With similar results from BeppoSAX DeGrandi and Molendi (2001)  and DeGrandi et al. (2003) showed that these iron abundance peaks are a signature of so-called cooling flow clusters, cluster with a very peaked X-ray surface brightness profile and thus a high central gas density, which we term cooling core clusters in this paper. Again with observational data from ASCA, Finoguenov et al. (2000) could derive abundance profiles of Fe and Si and could show that the central regions in those clusters with metal abundance peaks where strongly enriched in iron while the bulk volume of the clusters outside the central region had an iron-to-silicon ratio close to the yields of SN type II. This implies that the central excess seen in cooling core clusters is most probably due to enrichment by SN type Ia in the cD galaxies.  Now with the improved capabilities of the XMM-Newton satellite observatory as well as with the Chandra observatory, more detailed information on the origin, the metal source distribution, and the ICM transport history can be obtained. In this paper we compile information on the central element abundance profiles from four nearby clusters observed with XMM and interpret the results. We discuss the origin of the X-ray emitting gas in terms of the enrichment by supernovae type Ia and whether it should be considered to be the interstellar medium of the central galaxy or the cluster ICM. We calculate the time it takes to produce the Fe abundance peak and discuss the implications of the implied large enrichment times for the cluster evolution history.  Throughout the paper we use a Hubble constant of $h_{70} = 1$ where $h_{70} = H_0 / 70$ km s$^{-1}$ Mpc$^{-1}$. The other cosmological parameters play not significant role at the relevant distances. The Virgo distance is assumed independently of $H_0$ as $17$ Mpc. ",
        "conclusions": "The large metal abundance excesses in cooling core regions of clusters of galaxies indicate that large enrichment times are necessary to create them. For the four examples studied here, the Virgo, Perseus, Centaurus, and Abell 1795 clusters, we find enrichment times of about 4 - 10 Gyrs for the central 50 kpc radius region. Thus we conclude that the dense cooling cores in clusters are very persistent phenomena. These studies can easily be expanded to a larger sample of clusters to test if these conclusions hold in general. But since very similar abundance gradients are observed for most cooling core clusters we expect that this will be the case. A more detailed study of the abundance ratios of several tracer elements will also provide more detailed information on the enrichment history by the different SN types as well as on the possible secular evolution of the SNIa."
    },
    "0402/astro-ph0402093_arXiv.txt": {
        "abstract": "{This paper will review a new technique of detecting companion stars in LMXBs and X-ray transients in outburst using the Bowen fluorescence lines at $\\lambda\\lambda$4634-4640. These lines are very efficiently reprocessed in the atmospheres of the companion stars, and thereby provide estimates of the $K_2$ velocities and mass functions. The method has been applied to Sco X-1, X1822-371 and GX339-4 which, in the latter case, provides the first dynamical evidence for the presence of an accreting black hole. Preliminary results from a VLT campaign on V801 Ara, V926 Sco and XTE J1814-338 are also presented.}    \\addkeyword{accretion} \\addkeyword{accretion discs} \\addkeyword{X-rays: binaries} \\addkeyword{X-ray: stars}   \\begin{document} ",
        "introduction": "\\label{sec:intro}  The Galaxy is populated with $\\sim$50 known {\\bf persistent}, bright Low Mass X-ray Binaries (LMXBs) whose optical emission is triggered by X-ray reprocessing in the gas surrounding the compact object, mainly the accretion disc. The companion star is $\\sim~10^6$ times fainter than the optical disc and hence completely undetectable. This has hampered dynamical studies of LMXBs which have been restricted so far to radial velocity studies of X-ray transients in {\\bf quiescence} (e.g. Charles \\& Coe 2003). In several cases, the quiescent companion spectrum is just too faint for current instrumentation (e.g. GX339-4, N Oph 93) or the target is contaminated by a bright line-of-sight star (e.g. Aql X-1, 4U 2129+47). Dynamical studies and mass determination of compact stars in LMXBs can yield new black hole discoveries and, more importantly, probe for the existence of ``massive'' neutron stars. The latter would rule out soft equations of state and give further support that LMXBs are indeed the progenitors of millisecond pulsars, spun up by accretion.  \\begin{figure}[!t] \\includegraphics[width=\\columnwidth, height=5cm]{jcasares_fig1.ps} \\caption{Trailed spectrogram of the Bowen blend and \\ion{He}{II} $\\lambda$4686 line showing the Doppler shift of the narrow \\ion{C}{III} and \\ion{N}{III} components. From Steeghs \\& Casares (2002).} \\label{fig:trail} \\end{figure} ",
        "conclusions": ""
    },
    "0402/astro-ph0402436_arXiv.txt": {
        "abstract": "The Gemini Deep Deep Survey (GDDS) is an ultra-deep ($K<20.6$ mag, $I<24.5$ mag) redshift survey targeting galaxies in the ``redshift desert'' between $1<z<2$. The primary goal of the survey is to constrain the space density at high redshift of evolved high-mass galaxies. We obtained 309 spectra in four widely-separated 30 arcmin$^2$ fields using the Gemini North telescope and the Gemini Multi-Object Spectrograph (GMOS). The spectra define a one-in-two sparse sample of the reddest and most luminous galaxies near the $I-K$\\, vs. $I$\\, color-magnitude track mapped out by passively evolving galaxies in the redshift interval $0.8<z<1.8$.  This sample is augmented by a one-in-seven sparse sample of the remaining high-redshift galaxy population.  The GMOS spectrograph was operating in a Nod \\& Shuffle mode which enabled us to remove sky contamination with high precision, even for typical exposures times of 20--30 hours per field.  The resulting spectra are the deepest ever obtained.  In this paper we present our sample of 309 spectra, along with redshifts, identifications of spectral features, and photometry. This makes the GDDS the largest and most complete infrared-selected survey probing the redshift desert.  The 7-band ($VRIzJHK_s$) photometry is taken from the Las Campanas Infrared Survey.  The infrared selection means that the GDDS is observing not only star-forming galaxies, as in most high-redshift galaxy surveys, but also quiescent evolved galaxies.  In our sample, we have obtained 225 secure redshifts, 167 of which are in the redshift interval $0.8 < z < 2$. About 25\\% of these show clear spectral signatures of evolved (pure old, or old + intermediate-age) stellar populations, while 35\\% of show features consistent with either a pure intermediate-age or a young + intermediate-age stellar population. About 29\\% of the galaxies in the GDDS at $0.8 < z < 2$ are young starbursts with strong interstellar lines. A few galaxies show very strong post-starburst signatures. Another 55 objects have less secure redshifts, 31 of which lie in the redshift interval $0.8 < z < 2$.  The median redshift of the whole GDDS sample is $z=1.1$.  Spectroscopic completeness varies from a low of $\\sim70\\%$ for red galaxies to $>90\\%$ for blue galaxies. In this paper we also present, together with the data and catalogs, a summary of the criteria for selecting the GDDS fields, the rationale behind our mask designs, an analysis of the completeness of the survey, and a description of the data reduction procedures used. All data from the GDDS are publicly available. ",
        "introduction": "\\label{sec:introduction}   The Gemini Deep Deep Survey (GDDS) is an infrared-selected ultra-deep spectroscopic survey probing the redshift range $0.8<z<1.8$. It is designed to target galaxies of all colors at high redshift with an emphasis on the reddest population.  The survey is designed with the following scientific goals in mind: (1) Measurement of the space density and luminosity function of massive early-type galaxies at high redshift. (2) Construction of the volume-averaged stellar mass function in at least three mass bins and two redshift bins over the target redshift range. (3) Measurement of the luminosity-weighted ages and recent star-formation histories of $\\sim 50$ evolved galaxies at $z>1$ \\citep[cf.][]{dun96}. The over-arching goal of the survey is to use these sets of observations to test hierarchical models for the formation of early-type galaxies.  Many studies (see \\citet{ell01} for a recent review) have probed the evolving space density of early-type systems and it is now clear that the number density of early-types does not evolve rapidly out to $z=1$, as once predicted by matter-dominated models (see \\citealt {ell01} for a review).  However, $\\Lambda$-dominated cosmologies push back the formation epoch of most early-type systems out to at least $z=1$ even in a hierarchical picture.  Alternative theories for the origin of early-type galaxies (e.g. high-z monolithic collapse {\\em vs.} hierarchical formation from mergers) now start to become readily distinguishable at exactly the redshift ($z=1$) where spectroscopy from the ground becomes problematic \\citep{kau98}.  Our focus on the redshift range $0.8<z<1.8$ is motivated by two additional considerations. Firstly, the star-formation histories of individual galaxies in this redshift range have been very poorly explored.  We do not even know whether most red objects in this range are old and quiescent, or very young and active and heavily reddened by dust. Distinguishing between these two possibilities requires high signal-to-noise in the continuum so that the characteristic photospheric features of evolved stars become evident, but most work in this redshift range has focused on emission lines. Secondly, and irrespective of model predictions, this redshift range appears to correspond to the peak epoch of galaxy assembly inferred by integrating under the `Madau/Lilly plot', an observationally-defined diagram quantifying the volume-averaged star-formation history of the Universe as a function of redshift \\citep{mad96,lil96,ste99}. The high-redshift tail of this plot is subject to large and uncertain dust and surface-brightness corrections, and remains rather poorly determined, and recent observations have pushed back the peak of star-formation, showing a broad maximum in redshift (for a summary of the observational situation, see Figure 2 in Nagamine et al. 2003). However, the integral under the Madau/Lilly plot is simply the total mass assembled in stars per unit volume, so by measuring this quantity directly in the GDDS we can undertake a basic consistency check of the overall picture inferred from global star-formation history and luminosity density diagrams.  Spectroscopy of galaxies in the redshift range we seek to probe suffers from technical challenges brought on by the lack of strong spectral features at visible wavelengths. The redshift range $1<z<2$ has come to be known as the ``redshift desert'', in reference to the paucity of optical redshifts known over this interval. Fortunately, it is now becoming clear that redshifts and diagnostic spectra {\\em can} be obtained using optical spectrographs in this redshift range, by focusing on rest-frame UV metallic absorption features. Good progress is now being made in obtaining redshifts for UV-selected samples in the redshift desert using the blue-sensitive LRIS-B spectrograph on the Keck telescope \\citep{ste03,erb03}. However, UV-selected surveys are biased in favor of high star-formation rate galaxies, and the passive red galaxies with high mass-to-light ratios that are missed by UV-selection could well dominate the high-$z$ galaxy mass budget, motivating deep $K$-selected surveys such as the VLT K20 survey \\citep{cim03}, and ultra-deep small area surveys such as FIRES \\citep{fra03}. We refer the reader to \\citet{cim04} for an excellent summary of recent results obtained from infrared-selected surveys probing high-redshift galaxy evolution.  The GDDS is designed to build upon these results.  Because their rest-UV continuum is so weak, determining the redshifts of passive red galaxies at $z\\sim 1.5$ with 8m-class telescopes presently requires extreme measures. Ultra-deep ($>10$ hour) integration times and Poisson-limited spectroscopy are required in order to probe samples of red galaxies with zero residual star-formation and no emission lines.  This poses a severe problem, because MOS spectroscopy with 8m-class telescopes is generally not Poisson-limited unless exposure times are short (less than a few hours). The main contributors to the noise budget are imperfect sky subtraction and fringe removal. At optical wavelengths, both of these problems are most severe redward of 7000\\AA, where most of the light from evolved high redshift stellar populations is expected to peak. To mitigate against these effects, the Gemini Deep Deep Survey team has implemented a Nod \\& Shuffle sky-subtraction mode \\citep{gla01,cui94, bla95} on the Gemini Multi-Object Spectrograph \\citep{mur03,hoo03}. This technique is somewhat similar to beam-switching in the infrared, and allows sky subtraction and fringe removal to be undertaken with an order of magnitude greater precision than is possible with conventional spectroscopy.  In order to undertake an unbiased inventory of the high-redshift galaxy population, the GDDS is $K$-band selected to a sufficient depth ($K = 20.6$ mag) to reach $L^\\star$ throughout the $1 < z < 2$ regime. IR-selected samples that do not reach to this $K$-band limit are limited primarily to the $z < 1$ epoch, while samples that go substantially deeper outrun the capability of ground-based spectroscopic follow-up for the reddest objects, and once again become biased. The standard definition for an `Extremely Red Galaxy', or ERG, is $I-K\\ga 4$, a threshold which roughly corresponds to the expected color of an evolved dust-free early-type galaxy seen at $z\\sim 1$. As will be shown below, the effective limit for obtaining absorption-line redshifts for red galaxies with weak UV continua is about $I=25$ mag with an 8m telescope. (Our GDDS spectroscopy --- the deepest ever undertaken --- has a magnitude limit of $I=24.5$ mag). Therefore at present it is only just possible to obtain a nearly complete census of redshifts for all evolved red objects in a $K\\sim21$ imaging survey. Our strategy with the GDDS is to go deep enough to allow redshifts to be obtained for $L_\\star$ galaxies irrespective of star-formation history at $z\\sim 1.5$, while simultaneously covering enough area to minimize the effects of cosmic variance.  Our sampling strategy (based on photometric redshift pre-selection to eliminate low-$z$ contamination) is different from that adopted by most other redshift surveys. In terms of existing surveys, the K20 survey \\citep{cim03} is probably the closest benchmark comparison to the GDDS, although the experimental designs are very different, making the K20 and GDDS surveys quite complementary.  The K20 survey has about twice as many redshifts as the GDDS, but because K20 survey does not preferentially select against low-redshift objects, most of these are at $z<1$. The GDDS has between two and three times as many redshifts as K20 in the interval $1.2<z<2$ (the precise number depending on the minimum acceptable redshift confidence class), and has a higher median redshift ($z\\sim 1.1$ vs. $z\\sim0.7$).  The GDDS also goes about 0.6 mag deeper in $K$ and has over twice the area (121 square arcmin in four widely-separated site-lines in the GDDS vs. 52 square arcmin in two widely-separated site-lines in K20).  A plan for this paper follows. In \\S2 we describe our experimental design, with a particular focus on how our targets were selected from the {\\em Las Campanas Infrared Survey}.  In \\S3 we outline our observing procedure, but defer the details of the Nod \\& Shuffle mode that are not specific to the GDDS to an Appendix. (Nod \\& Shuffle on the Gemini Multi-Object Spectrograph was implemented for use on the GDDS but is now a common-user mode. Since many observers may wish to use the mode themselves in contexts unrelated to faint galaxy observations, the Appendices to this paper will act as a stand-alone reference to using the mode on Gemini).  In Section 4 we summarize the data obtained from the GDDS, both graphically and as a series of tables. Composite spectra obtained by co-adding similar spectra are presented in Section 5. We used these composite as templates to obtain redshifts in the GDDS, but others may wish to apply them to their own work for other purposes\\footnote{It should be born in mind that these composites are constructed from galaxies covering a wide range of redshift and time. Most analyses on composites would require a more restricted range, e.g.  paper II (Savaglio et al. 2004).}.  Some implications from the data obtained are discussed and our conclusions given in Section 6. The major results from the GDDS will be presented in three companion papers\\footnote{Savaglio et al. 2004 [paper II] presents measurements of column densities and metallicities of star-forming galaxies in our Sample. Glazebrook et al. 2004 [paper III] presents the mass function from the GDDS. McCarthy et al. 2004 [paper IV] will present an analysis of the stellar populations in the reddest galaxies in our sample.}.  Appendix A describes the operation of the Nod \\& Shuffle mode on the Gemini Multi-Object Spectrograph \\citep{hoo03} (GMOS) in the context of the GDDS. A more general description of the implementation of the mode will be given in Murowinski et al. (in preparation). Appendix B describes how the two-dimensional data from the GDDS were reduced. Appendix C describes how the final one-dimensional spectra were extracted from the two-dimensional data.  The catalogs presented in this paper, as well as reduced spectra for all galaxies in the GDDS, are available in electronic form as a digital supplement to this article. All software described in this paper, as well as auxiliary data, are publicly available (in both raw and fully reduced form) from the central GDDS web site located at http://www.ociw.edu/lcirs/gdds.html.  Throughout this paper we adopt a cosmology with $H_0$=70 km/s/Mpc, $\\Omega_M=0.3$, and $\\Omega_\\Lambda=0.7$. ",
        "conclusions": "\\label{sec:summary}  The Gemini Deep Deep Survey was undertaken in order to explore galaxy evolution near the peak epoch of galaxy building. The survey probes a color-selected sample in a manner that minimizes the strong star-formation rate selection biases inherent in most high-redshift galaxy surveys. It is designed to bridge the gap between landmark surveys of highly complete samples at $z<1$ \\citep{cow94,lil95,ell96} and UV-selected surveys at higher redshift \\citep{ste96,sha01,ste03}. The signal-to-noise ratios of the spectra in the GDDS are sufficient to distinguish old stellar populations (dominated by F-type stars) from post-starburst systems (dominated by A-type stars) and reddened starbursts with their flat spectra and strong interstellar lines. In this paper we have described the motivation for the survey, the choice of fields, the experimental design underlying our choice of targets, and our data reduction process. We have presented final catalogs of redshifts and photometry and an analysis of their statistical completeness. Spectra for individual objects are available as an electronic supplement to this paper. Further information on the GDDS and the data reduction software used in the project are available on the World Wide Web at \\url{http://www.ociw.edu/lcirs/gdds.html}."
    },
    "0402/astro-ph0402108_arXiv.txt": {
        "abstract": "{ We present the X-ray spectrum of TW\\,Hya observed at high and intermediate spectral resolution with the Reflection Grating Spectrometer (RGS) and the European Photon Imaging Camera (EPIC) onboard the {\\em XMM-Newton} satellite. TW\\,Hya is the first classical T Tauri star for which simultaneous X-ray data with both high spectral resolution and high sensitivity were obtained, thus allowing to probe the X-ray emission properties of stars in the early pre-main sequence phase.  Despite  TW\\,Hya's high X-ray luminosity in excess of $10^{30}$\\,erg/s its X-ray spectrum is dominated by emission lines from  rather cool plasma (T $\\approx$ 3 MK), and only little emission measure is present at high temperatures (T $\\approx$ 10 MK). We determine photon fluxes for the emission lines in the high resolution spectrum, confirming the earlier result from {\\em Chandra} that the predominant emission is from neon and oxygen, with comparatively weak iron lines. Further, the line ratios of He-like triplets of nitrogen, oxygen and neon require  densities of $n_{\\rm e} \\sim 10^{13}\\,{\\rm cm^{-3}}$, about two orders of magnitude higher than for any other star observed so far at high spectral resolution. Finally, we find that nearly all metals are underabundant with respect to solar abundances, while the abundances of nitrogen and neon are enhanced.  The high plasma density, the (comparatively) low temperature, and peculiar chemical abundances in the X-ray emitting region on TW\\,Hya are untypical for stellar coronae.  An alternative X-ray production mechanism is therefore called for and a natural explanation is an accretion column depleted of grain forming elements. The metal depletion could be either due to the original molecular cloud that formed TW\\,Hya or due to a settling of dust in the circumstellar disk of TW\\,Hya. ",
        "introduction": "\\label{sect:intro}  T Tauri stars (TTS) are late-type pre-main sequence (PMS) stars, usually found in or near star formation regions. The two subclasses of TTS, termed `classical' and `weak-line', are distinguished by the strength of their H$_\\alpha$ emission. This empirical distinction most likely also reflects different evolutionary states: Most of the `classical' TTS (cTTS), i.e., those with strong H$_\\alpha$ emission, show signatures of circumstellar matter (e.g. IR and UV excesses) accreting onto the central star, while in contrast, `weak-line' TTS (wTTS) show no signs for the presence of disks which presumably have been lost. A detailed description of early stellar evolution and the observable stellar properties related with the different phases can be found in \\citey{Feigelson99.1}.  X-ray emission from TTS was first detected with the {\\em Einstein} satellite (e.g. \\cite{Feigelson81.1}, \\cite{Walter81.1}, \\cite{Montmerle83.1}). It was soon recognized that X-ray observations provide an efficient means of discovering wTTS which can easily escape other identification methods because of their lack of conspicuous features in the optical or near-infrared. Extensive X-ray studies of a variety of star forming regions were undertaken with {\\em ROSAT}. In particular the spatial completeness of the {\\em ROSAT} All-Sky Survey has resulted in a large number of newly identified PMS stars (e.g. \\cite{Neuhaeuser95.1}, \\cite{Alcala95.1}, \\cite{Wichmann96.1}).  The origin of the X-ray emission from wTTS is commonly interpreted in terms of scaled-up solar-type magnetic activity, but the situation is less clear for their accreting counterparts, the cTTS. While nothing should prevent the formation of magnetic fields and ensuing coronae in cTTS, X-rays could alternatively also be produced above the hot spot where a magnetically funnelled accretion flow impacts on the surface of the star (\\cite{Lamzin99.1}). A distinction between a corona or an accretion funnel as the site of X-ray emission in cTTS requires a precise assessment of the physical conditions in the X-ray emitting region.  With the lower resolution X-ray data available from {\\it Einstein-} and {\\em ROSAT}-observations of both cTTS and wTTS it was impossible to adequately characterize the emission site, although some differences between the two types of TTS were indicated by statistical comparisons of their emission level and variability (see e.g. \\cite{Neuhaeuser95.1}, \\cite{Stelzer00.1}, \\cite{Stelzer01.1}).  With the advent of high-resolution X-ray spectroscopy with the grating spectrometers onboard the {\\em XMM-Newton} and {\\em Chandra} observatories this situation has fundamentally changed. The reflection gratings onboard {\\em XMM-Newton} and the transmission gratings onboard {\\em Chandra} have spectral resolutions sufficiently large to resolve individual emission lines in the soft X-ray range. These emission lines and in particular the analysis of their relative strengths provide unprecedented tools for the plasma diagnostics of the outer envelopes of stars and allow an assessment of the temperature structure, density, and chemical abundances in the X-ray emitting region.  The high signal-to-noise ratio (SNR) required for a meaningful analysis of high-resolution X-ray spectra clearly favors the study of the nearest star forming regions.  The nearby TTS TW\\,Hya (located at a distance of only $\\sim 57$\\,pc; \\cite{Wichmann98.1}) was first thought to be an isolated cTTS (\\cite{Rucinski83.1}), but later five more stars were found in the same region and shown to be members of a physical cluster of young stars dominated by TW Hya. The whole group is now known as the TW~Hydra association (TWA; \\cite{Kastner97.1}), which is currently known to host about $20$ young stars (youth being probed by Li absorption), most of which -- in contrast to TW\\,Hya itself -- are wTTS.  The first -- and so far only -- example of a high-resolution X-ray spectrum of an accreting PMS star was a {\\em Chandra}/HETGS spectrum of TW\\,Hya presented by \\citey{Kastner02.1}.  Interestingly, flux ratios of the OVII triplett lines indicated densities much higher than found in any stellar corona before, and consequently \\citey{Kastner02.1} called into question the solar-stellar analogy for TW\\,Hya and cTTS in general.  About a year after the {\\em Chandra} observation TW\\,Hya was observed by {\\em XMM-Newton}, resulting in a higher SNR grating spectrum with the {\\em XMM-Newton} Reflection Grating Spectrometer (RGS) as well as a simultaneously measured medium-resolution CCD spectrum with high sensitivity over a broad energy range ($E = 0.3-10$\\,keV) from the European Photon Imaging Camera (EPIC). The purpose of this paper is to present the results of this 30\\,ksec {\\em XMM-Newton} observation and properly discuss their implications in the context of X-ray emission from TTS. ",
        "conclusions": "\\label{sect:discussion}  So far TW\\,Hya is the only cTTS, i.e. a PMS star with strong H$_\\alpha$ emission, for which a high-resolution {\\em XMM-Newton} X-ray spectrum is available. Since in young stars H$_\\alpha$ emission is related to accretion TW\\,Hya presents a unique test case for X-ray emission in this early phase of stellar evolution. Furthermore, it is one of the few stars observed with both {\\em XMM-Newton} and {\\em Chandra} at high spectral resolution. In the following we discuss the properties of the X-ray emission of  TW\\,Hya, and we present a physical scenario explaining it.   \\subsection{Variability}\\label{subsect:variability}  During the {\\em Chandra} observation of TW\\,Hya discussed by \\citey{Kastner02.1} a sharp rise in intensity followed by a nearly linear decay occurred, that the authors interpreted as a flare. We find little evidence for variability in the {\\em XMM-Newton} observation presented here (acquired $\\sim$\\,1\\,yr after the {\\em Chandra} data). In any case the lightcurve does not carry the clear signatures of a flare, and no change in spectral hardness is observed. Similarly \\citey{Kastner02.1} found no changes in the spectral parameters during the variable phase of the {\\em Chandra} observation, very untypical for stellar (flare) activity. The flare-like event in the {\\em Chandra} observation has a decay time of $\\sim 10$\\,ksec. Assuming a flare loop at $T \\sim 3$\\,MK (the dominating temperature derived from the X-ray spectrum) that decays through radiative cooling the density should be $\\sim 10^{10}\\,{\\rm cm^{-3}}$. This is in sharp contrast to the result from the analysis of line ratios (see Sect.~\\ref{subsect:triplets}), according to which the densities are higher by two orders of magnitude.  We directly compare the emission levels by computing X-ray fluxes for the time-averaged spectrum. From the {\\em Chandra} HETGS spectrum \\citey{Kastner02.1} derived an X-ray flux of $2.5\\times10^{-3}\\,{\\rm photons\\,cm^{-2}\\,s^{-1}}$ for the $0.45-6.0$\\,keV band. Integrating the RGS spectrum in the energy band which comes nearest to the spectral range of the HETGS, i.e. $5.5 - 27.5$\\,\\AA\\ or $0.45-2.25$\\,keV, we arrive at a similar value of $f_{\\rm x,RGS\\,1} \\approx f_{\\rm x,RGS\\,2} = 2.3\\times10^{-3}\\,{\\rm photons\\,cm^{-2}\\,s^{-1}}$ corresponding to $L_{\\rm x,RGS} = 1.4\\times10^{30}\\,{\\rm erg/s}$ at a distance of $57$\\,pc. This suggests that TW\\,Hya exhibits only moderate long-term variability.  Based on the small amplitude of long-term variations discussed above, comparing fluxes measured for individual emission lines allows to check the calibration of the instruments on both satellites. We observe a trend towards slightly higher fluxes in the RGS as compared to the HETGS (Fig.~\\ref{fig:fluxcomp}). However, the difference is a factor two at most. \\begin{figure} \\begin{center} \\resizebox{9cm}{!}{\\includegraphics{./fig6.eps}} \\caption{Comparison of photon fluxes for individual emission lines in the X-ray spectrum of TW\\,Hya measured with the {\\em XMM-Newton} RGS (see Table~\\ref{tab:rgs_lines}) and the {\\em Chandra} HETGS-MEG (see \\protect\\cite{Kastner02.1}). The {\\em XMM-Newton} and {\\em Chandra} observations are separated by about one year. All lines are constant within about a factor of two.} \\label{fig:fluxcomp} \\end{center} \\end{figure}     \\subsection{Density in the X-ray Emitting Region}\\label{subsect:triplets}  \\begin{figure*} \\begin{center} \\parbox{18cm}{ \\parbox{9cm}{\\resizebox{9cm}{!}{\\includegraphics{./fig7a.eps}}} \\parbox{9cm}{\\resizebox{9cm}{!}{\\includegraphics{./fig7b.eps}}} } \\caption{He-like triplets of Ne\\,IX and O\\,VII measured with the RGS together with best fit Lorentzians. The continuum close to each triplet is approximated by a straight horizontal line. The intensity enhancement to the right of the Ne\\,IX triplet in the RGS spectrum is due to Fe\\,XVII. But the major contamination of the Ne\\,IX triplet is expected to come from a Fe\\,XIX line which is unresolved from the triplet intercombination line (see text in Sect.~\\ref{subsect:triplets}).} \\label{fig:triplets} \\end{center} \\end{figure*}  The high-resolution X-ray spectrum {obtained with the RGS} allows us to measure flux ratios between individual emission lines and use them for plasma diagnostics. The triplets of He-like ions are of particular importance, because ratios between the intensity of resonance $r$, intercombination $i$, and forbidden $f$ line are sensitive to (electron) temperature (through $G=(f+i)/r$) and density (through $R=f/i$). While the $G$-ratio is of limited value for temperature diagnostics, the $R$-ratio constitutes a unique tool to probe the electron density in stellar coronae. Although the triplets from silicon (Si\\,XIII) to nitrogen (N\\,VI) fall in the spectral range of the RGS, in TW\\,Hya only three elements -- nitrogen, oxygen, and neon -- show sufficient emission in the respective transitions of their He-like ions to warrant analysis.  In Fig.~\\ref{fig:triplets} we zoom in on the spectral region around the oxygen and neon triplets of TW\\,Hya overlaying the data with our line fit result. For the modeling of both triplets we used three Lorentzians representing the $r$, $i$, and $f$ lines. Some caution is in order when interpreting the spectrum near the position of the Ne\\,IX triplet: the triplet is not fully resolved in the RGS, and is known to form a blend with a number of iron lines, predominantly a Fe\\,XIX line at 13.53\\,\\AA~ that contaminates the Ne intercombination line.  However, we performed fits of the neon triplet including iron and find undistinguishable results.  Since we find little hints for the presence of Fe\\,XIX (and generally any iron) in the RGS spectrum (e.g. the lines at 13.80\\,\\AA~ and 14.66\\,\\AA~ are not seen) we assume that iron contamination plays only a minor role and can be neglected.  The remarkable feature of the He-like triplets of TW\\,Hya is the weakness of the forbidden lines. In the LETGS and RGS spectra of all late-type stars observed so far the forbidden lines are always stronger than the intercombination line (in oxygen and neon) and in many cases the measurements are consistent with the so-called low-density limit (e.g. \\cite{Ness01.1}, \\cite{Ness02.1}, \\cite{Stelzer02.1}). This is definitely not the case for TW Hya. We note in particular that with TW\\,Hya being rather cool (spectral type K7; \\cite{Rucinski83.1}) the low intensity of the forbidden line cannot be attributed to UV radiation at the wavelengths corresponding to the $f \\longrightarrow i$ transitions (\\cite{Costa00.1}). Instead it is a direct indicator of high densities. The formally derived $f/i$-value for {\\rm O\\,VII} is $R_{\\rm O\\,VII} < 0.16$, i.e., only an upper limit can be given due to the weakness of the {\\rm O\\,VII} $f$-line. Comparing our measurements to recent calculations by \\citey{Porquet01.1} without any radiation field we find a lower limit on the density of $n_{\\rm e,O\\,VII} > 1\\times10^{12}\\,{\\rm cm^{-3}}$. This is corroborated by the low flux of the nitrogen forbidden line which is only marginally detected (cf., Table \\ref{tab:rgs_lines}). The $R$-ratio for neon ($R_{\\rm Ne\\,IX} = 0.47 \\pm 0.11$) indicates a density of $n_{\\rm e, Ne\\,IX} \\sim 1\\times10^{13}\\,{\\rm cm^{-3}}$. Inclusion of the Fe\\,XIX blend in the fit of the neon triplet would lower the strength of the $i$ line, and lead to a somewhat higher value for $R_{\\rm Ne\\,IX}$ ($0.62 \\pm 0.12$), and subsequently lower density. However the numbers are compatible with the ones cited above within the statistical uncertainties.  The densities derived from the oxygen and neon triplets are at least two orders of magnitude above typical coronal densities. Our density measurements thus fully confirm the results derived by \\citey{Kastner02.1} from their HETGS spectrum, and we conclude in particular that the high densities cannot be a result of the flaring reported by \\citey{Kastner02.1} in their {\\em Chandra} data.    \\subsection{Elemental abundances}\\label{subsect:abund}  The {\\em XMM-Newton} RGS spectrum as well as our spectral fits to the EPIC spectrum show the X-ray emission of TW\\,Hya  clearly dominated by neon emission in agreement with \\citey{Kastner02.1}. Our results specifically suggest a neon abundance significantly enhanced with respect to the solar value ($Ne \\sim 4\\,Ne_\\odot$), while iron and oxygen, the elements with the largest number of lines in the sensitive region of the {\\em XMM-Newton} and {\\em Chandra} instruments, are underabundant with respect to solar values, again in agreement with \\citey{Kastner02.1}. For the elements silicon, magnesium, and nitrogen \\citey{Kastner02.1} claim solar abundances, while carbon lines were not covered by the HETGS spectrum. Interestingly, the lines of He-like and H-like silicon are clearly visible in the HETGS spectrum of TW\\,Hya (cf., Fig.~2 in \\cite{Kastner02.1}), while the same authors note ``a curious lack of Mg features, relative to the [i.e. their] model.''  While the RGS spectra are not sufficiently sensitive in the relevant spectral region, the EPIC spectrum clearly suggests subsolar abundances also for magnesium and silicon.  Let us now focus on the line emission from neon, oxygen, nitrogen, and carbon as seen by the RGS. Our spectral analysis both with RGS and EPIC shows that very little low temperature material (with $\\log{T}\\,{\\rm [K]} < 6.3$) can be present. \\citey{Ness02.1} report photon flux ratios for $Ly_\\alpha/r$ of 0.59 (oxygen) and $0.94$ (nitrogen) for the low-temperature F-star Procyon, and $1.87$ (oxygen) and $3.26$ (nitrogen) for the higher-temperature K-star $\\epsilon$\\,Eri, which must be compared to the values of $1.75$ (oxygen) and $4.07$ (nitrogen) measured for TW\\,Hya despite its overall low temperature. Although DEM distributions are typically not unique because of the ill-defined nature of the inversion problem, we want to emphasize that the peculiar chemical abundances of TW Hya are not caused by our choice of the EM distribution. For this we examined the energy flux ratio between the $Ly_\\alpha$ lines of C\\,VI and N\\,VII as a function of temperature and find that this line ratio can not be explained by any single temperature or combination of temperatures under the assumption of solar abundances for both elements (Fig.~\\ref{fig:temp_ratios2}\\,a). Therefore nitrogen must be overabundant relative to carbon (and also oxygen) with respect to the solar abundance ratio, again in accordance with our fit results. We then investigated the flux ratio between the $r$ line of neon and the oxygen $Ly_\\alpha$ line in the same way (Fig.~\\ref{fig:temp_ratios2}\\,b), and again the observed flux ratio can not be reproduced by any single temperature or combination of temperatures if the abundances were solar, implying that neon must be overabundant compared to oxygen with respect to solar abundances, as found in our fit results. Finally, inspecting the flux ratio between the C\\,VI $Ly_\\alpha$ line and the O\\,VII $r$ line we find perfect consistency assuming the C/O ratio being solar. Since our EPIC results suggest that silicon, magnesium, iron, and oxygen all have similar  abundance (with the possible exception of magnesium) we thus conclude that the elemental abundances of these elements as well as that of carbon are reduced with respect to solar values by about a factor of $3-4$. Nitrogen is enhanced with respect to these elements by a factor of $3$ as evidenced by the {\\em XMM-Newton} RGS spectrum, and neon is enhanced by a factor of $10$ as evidenced by both the {\\em XMM-Newton} RGS and EPIC spectra. \\begin{figure*} \\begin{center} \\parbox{18cm}{ \\parbox{6cm}{\\resizebox{6cm}{!}{\\includegraphics{./fig8a.ps}}} \\parbox{6cm}{\\resizebox{6cm}{!}{\\includegraphics{./fig8b.ps}}} \\parbox{6cm}{\\resizebox{6cm}{!}{\\includegraphics{./fig8c.ps}}} } \\caption{Temperature dependence of various line ratios measured in the RGS spectrum of TW\\,Hya. Solid lines are CHIANTI model calculations (\\protect\\cite{Dere97.1}), shaded regions denote the $1\\,\\sigma$ range observed with the RGS. (a) - N\\,VII $Ly_\\alpha$ / C\\,VI $Ly_\\alpha$; (b) - Ne\\,IX $r$ / O\\,VIII $Ly_\\alpha$; (c) - C\\,VI $Ly_\\alpha$ / O\\,VII $r$. The plots demonstrate that the C/O abundance ratio is close to solar, while the N/O and Ne/O abundance ratio can not possibly be solar.} \\label{fig:temp_ratios2} \\end{center} \\end{figure*}    \\subsection{Interpretation}\\label{subsect:inter}  The three properties distinguishing the X-ray spectrum of TW\\,Hya from that of ``normal'' late-type stars are (1) the absence of a hot component with large emission measure, (2) the high density of the X-ray emitting region, and (3) the peculiar elemental composition, and in particular the low abundance of iron. The first two properties can be naturally explained by attributing the X-ray emission to an accretion shock, the third by the specific environment of TW\\,Hya.  Employing the strong shock formula $$ T = \\frac {3 \\mu m_{\\rm H} V_{\\rm pre}^2} {16 k} $$ with $\\mu$ denoting the mean molecular weight, $m_{\\rm H}$ the hydrogen atom mass, and $k$ Boltzmann's constant, we can convert the measured temperature $\\lg{T}\\,{\\rm [K]} = 6.45$ to an infall preshock velocity $V_{\\rm pre}$ of about $350/\\sqrt{(\\mu)}$\\,km/s. The preshock particle density $n_{\\rm pre}$ should be one fourth of the measured post shock density $n_{\\rm post}$, i.e., $2.5 \\times 10^{12}\\,{\\rm cm^{-3}}$. In order to estimate $L_{\\rm post}$, the thickness of the postshock cooling layer, we equate the infalling kinetic energy flux with the radiated energy flux (\\cite{Lamzin95.1}) through $$ \\frac {\\mu m_{\\rm H} n_{\\rm pre} V_{\\rm pre}^3} {2} = n_{\\rm post}^2 \\Lambda_{\\rm mean} L_{\\rm post}.$$ Here $\\Lambda_{\\rm mean}$ denotes the total mean radiative post shock cooling rate per unit emission measure. Denoting $\\Lambda_{\\rm mean}$ in units of 10$^{-23}\\,{\\rm erg\\,cm^3\\,s^{-1}}$, we derive $$ L_{\\rm post} = \\frac {9 \\times 10^{7}} {\\sqrt{\\mu}} {(\\Lambda_{\\rm mean,23})}^{-1}\\ \\ [{\\rm cm}].$$ For an isothermal cosmic abundance plasma at $\\lg{T}\\,{\\rm [K]}=6.45$ $\\Lambda_{\\rm mean,23}$ is generally close to unity; if we assume  $\\Lambda_{\\rm mean,23} = 0.3$ because of the depleted elemental abundances and $\\mu = 1.5$ the post-shock thickness is $L_{\\rm post} \\approx 2500$\\,km with admittedly considerable uncertainty. Using  the observed O\\,VIII energy flux of $3.2 \\times 10^{-13}\\,{\\rm erg/cm^2/s}$ and an assumed O\\,VIII $Ly_{\\alpha}$ cooling function of $9.6 \\times 10^{-25}\\,{\\rm erg\\,cm^3\\,s^{-1}}$ (assuming $\\lg{T}\\,{\\rm [K]}=6.45$ and an oxygen depletion of $0.3$) we compute an overall volume emission measure of $EM = 1.3\\times10^{53}\\,{\\rm cm^{-3}}$ at the distance of $57$\\,pc. Since $n_{\\rm post} \\approx 10^{13}\\,{\\rm cm^{-3}}$, the emitting volume $V_{\\rm emit}$ must be $V_{\\rm emit} = 1.3\\times10^{27}\\,{\\rm cm^{3}}$. The shock area $A_{\\rm shock}$ then becomes $A_{\\rm shock}= V_{\\rm emit}/L_{\\rm post} = 5.1\\times10^{18}\\,{\\rm cm^2}$, which would constitute only a very tiny fraction of the visible stellar surface ($< 0.01$\\,\\%). Clearly, these estimates are by necessity very rough, but in qualitative agreement with previous observations in other wavelengths: \\citey{Muzerolle02.1} have shown that the shape of the optical/UV spectral energy distribution of TW\\,Hya can be explained by an accretion shock which fills $\\sim 0.3$\\,\\% of the stellar surface, and small accretion shocks are also found on other T Tauri stars (\\cite{Calvet98.1}). Computing the mass accretion rate in the funnel flow through $M_{\\rm acc} = 1/4 A_{\\rm shock} n_{\\rm post} \\mu m_H V_{\\rm pre}$ we find $M_{\\rm acc} \\sim 1 \\times 10^{-11}\\,{\\rm M_\\odot/yrs}$, which is entirely plausible and within one order of magnitude of the mass accretion rate of TW\\,Hya derived from UV data (\\cite{Muzerolle02.1}).  What kind of material is flowing through this funnel onto TW Hya ? We have shown that this material is in general metal-depleted in particular with respect to grain-forming elements such as Mg, Si, Fe, C and O. Overabundances of neon are found in stellar coronae of RS~CVn systems (\\cite{Audard03.1}), and nitrogen overabundances in evolved stars (\\cite{Schmitt02.1}). The low abundances especially of iron and carbon call for a different explanation for these chemical peculiarities. TW\\,Hya is rather young and formed ``just recently'' out of a molecular cloud. Models of cloud chemistry at high densities ($n > 10^7\\,{\\rm cm^{-3}}$) predict that almost all metals condense into grains except nitrogen (\\cite{Charnley97.1}). If the remaining depleted gas is then decoupled from the grains in an accretion disk around a young star (as is expected in planetary formation models), the accreted and shock-heated gas is expected to be metal-deficient as well. Indeed, \\citey{Weinberger02.1} found evidence for the presence for condensed silicates in the disk of TW\\,Hya, and \\citey{Herczeg02.1} speculated about a deficiency of Si in the accretion flow of TW\\,Hya  based on the examination of its far-UV spectrum. {\\em XMM-Newton} has completed this picture of TW\\,Hya as an accreting PMS star, and for the first time provided a consistent view of X-ray emission from a stellar accretion shock by observation.       To summarize, the marked differences between the X-ray properties of TW\\,Hya and active main-sequence dwarf stars suggest either (i) strong evolutionary effects on coronal structures, or (ii) an entirely different origin for the X-ray emission of TW\\,Hya. X-ray emission from a shock at the bottom of an accretion column provides a plausible mechanism for the latter possibility. The derived shock temperatures, gas densities, filling factors and mass accretion rates as well as the obvious lack of grain-forming elements in the X-ray emitting region provide strong indications in favor of this scenario.  A high-resolution X-ray spectrum of another member of the TW\\,Hya association, HD\\,98800, was recently obtained with {\\em Chandra} (CXC Press Release 03-04). The object is a quadruple star with a circumbinary disk. In terms of H$_\\alpha$ emission HD\\,98800 is a wTTS. Therefore it does not come as a surprise that its X-ray properties are very different from those of TW\\,Hya, and clearly reminiscent for typical stellar activity. Similarly, the high-resolution X-ray observation of PZ\\,Tel, a $\\sim 20$\\,Myr old post-TTS in the Tucanae association, does not show properties drammatically different from older late-type stars (Argiroffi et~al., in prep.).  We conclude that TW\\,Hya stands out as a presently unique example for a star with strong evidence for accretion signatures in its X-ray spectrum. Future high spectral resolution and high sensitivity X-ray observations will have to show whether soft emission, high densities, and metal-deficiency are general characteristics of cTTS."
    },
    "0402/astro-ph0402422_arXiv.txt": {
        "abstract": "We present a complete census of RR Lyrae stars in a halo field of the Andromeda galaxy.  These deep observations, taken as part of a program to measure the star formation history in the halo, spanned a period of 41 days with sampling on a variety of time scales, enabling the identification of short and long period variables.  Although the long period variables cannot be fully characterized within the time span of this program, the enormous advance in sensitivity provided by the Advanced Camera for Surveys on the Hubble Space Telescope allows accurate characterization of the RR Lyrae population in this field. We find 29 RRab stars with a mean period of 0.594 days, 25 RRc stars with a mean period of 0.316 days, and 1 RRd star with a fundamental period of 0.473 days and a first overtone period of 0.353 days.  These 55 RR Lyrae stars imply a specific frequency $S_{RR}\\approx 5.6$, which is large given the high mean metallicity of the halo, but not surprising given that these stars arise from the old, metal-poor tail of the distribution. This old population in the Andromeda halo cannot be clearly placed into one of the Oosterhoff types: the ratio of RRc/RRabc stars is within the range seen in Oosterhoff II globular clusters, the mean RRab period is in the gap between Oosterhoff types, and the mean RRc period is in the range seen in Oosterhoff I globular clusters.  The periods of these RR Lyraes suggest a mean metallicity of [Fe/H]$\\approx-1.6$, while their brightness implies a distance modulus to Andromeda of 24.5$\\pm 0.1$, in good agreement with the Cepheid distance. ",
        "introduction": "The textbook picture of a spiral galaxy halo comes from that of our own Milky Way, which is old and metal-poor (VandenBerg 2000; Ryan \\& Norris 1991).  However, the stellar population of the Andromeda (M31; NGC224) halo offers a striking contrast to this picture, with its wide range in metallicity (Durrell, Harris, \\& Pritchet 2001) and age (Brown et al.\\ 2003; Brown 2003).  These spreads in metallicity and age can significantly affect the variable star population.  For example, the characteristics of RR Lyraes in Galactic globular clusters place these clusters into two distinct Oosterhoff types (Oosterhoff 1939), while the RR Lyraes in Local Group dwarf spheroidals (dSphs) place these galaxies in the gap between the two Oosterhoff types (Siegel \\& Majewski 2000; Dall'Ora et al.\\ 2003; Pritzl et al.\\ 2002, 2004).  Compared to the M31 halo, dSphs have an even broader age range (van den Bergh 1999) and are generally more metal-poor (Mateo 1998), but the old ($> 10$~Gyr) component capable of producing RR Lyrae stars might be similar in each case.  The specific frequency of RR Lyraes in the M31 halo has been the subject of some debate.  In a field 40 arcmin from the \\linebreak  {\\small \\noindent $^1$Based on observations made with the NASA/ESA Hubble Space Telescope, obtained at the Space Telescope Science Institute, which is operated by AURA, Inc., under NASA contract NAS 5-26555. These observations are associated with proposal 9453.}  \\noindent nucleus on the southeast minor axis, Pritchet \\& van den Bergh (1987) found 30 RR Lyraes, and with an estimated completeness of 25\\%, determined that the frequency per unit luminosity was very high (about half of that in variable-rich M3).  More recently, Dolphin et al.\\ (2003) found only 24 RR Lyraes in a larger field that included the Pritchet \\& van den Bergh (1987) field, and with their estimated completeness of 24\\%, claimed that the frequency of RR Lyrae was $\\sim$15 times smaller.  To support their claim, they estimated that the deep color magnitude diagram (CMD) of Brown et al.\\ (2003) contained only 10 RR Lyraes, but as we shall show here, this was a severe underestimate.  We have observed a field along the southeast minor axis of the M31 halo, 51 arcmin from the nucleus, using the Advanced Camera for Surveys (ACS; Ford et al. 1998) on the Hubble Space Telescope (HST). The primary goal of this program was to investigate the halo star formation history, by constructing a deep CMD in the F606W (broad $V$) and F814W ($I$) bandpasses, reaching $V\\approx 30.7$~mag on the main sequence (Brown et al.\\ 2003). However, the 250 individual exposures are scattered with variable time sampling over a 41 day period, and thus provide excellent time series photometry for the variable star population in the M31 halo; in particular, the completeness for RR Lyraes in our field is approximately 100\\%.  In this paper, we present a survey of the RR Lyraes and the other bright variables in our field. ",
        "conclusions": "We have presented a complete survey of the RR Lyrae stars in an M31 halo field, 51 arcmin from the nucleus.  We find 29 RRab stars with a mean period of 0.594 days, 25 RRc stars with a mean period of 0.316 days, and 1 RRd star with a fundamental period of 0.473 days and a first overtone period of 0.353 days.  The RR Lyrae population of the M31 halo cannot be clearly placed into either Oosterhoff type, and is distinct from the Milky Way cluster and halo field populations.  In a broad sense, the Local Group dSphs share the intermediate Oosterhoff status found the M31 halo, but the characteristics of the M31 RR Lyraes ($<$$P_{ab}$$>$, $<$$P_{c}$$>$, $N_c / N_{abc}$) are distinct from the dSphs, suggesting that the M31 halo is not comprised of dissolved globular clusters like those in the Milky Way or dissolved Local Group dSphs.  The specific frequency of RR Lyraes ($S_{RR} = 5.6$) is very high for a mean halo metallicity of [Fe/H]$=-0.8$, but within the normal range when considered as a component of the old, metal-poor halo population.  The mean metallicity of the RR Lyrae population is indeed much lower than that of the halo, with a mean [Fe/H]$=-1.77$ for the RRab stars and a mean [Fe/H]$=-1.43$ for the RRc stars.  The distance to M31 determined from the RR Lyrae luminosity is $(m-M)_o = 24.5 \\pm 0.1$~mag, in good agreement with the Cepheid distance."
    },
    "0402/astro-ph0402614_arXiv.txt": {
        "abstract": "{ We present high resolution ($\\lambda$/$\\Delta \\lambda$~=~49\\,000) \\'echelle spectra of the intermediate mass, pre-main sequence stars \\object{BF Ori}, \\object{SV Cep}, \\object{WW Wul} and \\object{XY Per}. The spectra cover the range 3800~--~5900~\\AA\\ and monitor the stars on time scales of months and days. All spectra show a large number of Balmer and metallic lines with variable blueshifted and redshifted absorption features superimposed to the photospheric stellar spectra. Synthetic Kurucz models are used to estimate rotational velocities, effective temperatures and gravities of the stars. The best photospheric models are subtracted from each observed spectrum to determine the variable absorption features due to the circumstellar gas; those features are characterized in terms of their velocity, $v$, dispersion velocity, $\\Delta v$, and residual absorption, $R_{\\rm max}$. The absorption components detected in each spectrum can be grouped by their similar radial velocities and are interpreted as the signature of the dynamical evolution of gaseous clumps with, in most cases, solar-like chemical composition. This infalling and outflowing gas has similar properties to the circumstellar gas observed in UX~Ori, emphasizing the need for detailed theoretical models, probably in the framework of the magnetospheric accretion scenario, to understand the complex environment in Herbig Ae (HAe) stars. WW Vul is unusual because, in addition to infalling and outflowing gas with properties similar to those observed in the other stars, it shows also transient absorption features in metallic lines with no obvious counterparts in the hydrogen lines. This could, in principle, suggest the presence of CS gas clouds with enhanced metallicity around WW~Vul. The existence of such a metal-rich gas component, however, needs to be confirmed by further observations and a more quantitative analysis. ",
        "introduction": "Observations reveal that the dynamics of the circumstellar (CS) gaseous disks around intermediate mass, young main sequence (MS) and pre-main sequence (PMS) stars is extremely complex. Variable absorption components detected in many lines of different elements and ions constitute good examples of such complexity. The kinematics and intensity strength of the absorption components contain relevant information on the physical properties of the gas. Further, their detailed characterization and analysis provide clues and constraints on plausible formation mechanisms as well as on theoretical scenarios describing the structure and nature of the CS gaseous disks.  The presence of metal-rich planetesimals in the young MS  $\\beta$~Pic system has been inferred both observationally and theoretically in a long series of papers \\citep[e.g.][and references therein]{lagrange2000}. Summarizing, there are two main arguments on which this inference is based. Firstly, dust causing the far-IR excess and also seen in the $\\beta$~Pic disk images may be second generation material continuously replenished by collisions of large solid bodies or slow evaporation. Secondly, transient spectral line absorptions, usually redshifted, of different chemical species in a wide range of ionization states can be modelled in terms of the evaporation of km-sized bodies on star-grazing orbits. Star-grazing planetesimals have also been suggested to exist in the 51~Oph system, a star with an uncertain evolutionary status \\citep{roberge2002,vandenancker2001}; also, blueshifted absorption in excited fine structure lines of \\ion{C}{ii}$^*$ at 1037~\\AA\\ and \\ion{N}{ii}$^*$ at 1085 and 1086~\\AA\\ have been used to infer the presence of $\\sim$1~m bodies in the Vega-type binary MS system $\\sigma$~Her \\citep{chen2003}.  Absorption features similar to those observed in $\\beta$~Pic have been observed towards many HAe stars (see Natta et al. 2000 for the description of these stars as a subgroup of HAeBe stars), particularly in the UXOR-subclass \\citep[e.g.][]{grady1996} and, by analogy, they have been interpreted as indicative of the presence of large solid bodies in the CS disks  around these PMS stars \\citep[e.g.][and references therein]{grady2000}. In principle, this interpretation is not in conflic with the accepted time scale for the formation of planetesimals of $\\sim$10$^4$~yr \\citep{beckwith2000}, which suggests that planetesimals should already be present during the PMS phase of stars ($\\sim$1-10~Myr). This explanation for the variable absorption features observed in HAe stars is, however, controversial and, in fact, \\citet{grinin1994} pointed out other alternatives, more concretely dissipation of dust clouds and the simultaneous infall of cool gas onto the star. \\citet{natta2000} have analyzed the chemical composition of a strong redshifted event in UX~Ori and shown it to have a solar-like composition. Instead of the planetesimal origin for the transient components, those authors suggested gas accretion from a CS disk. This result is supported by \\citet[from now on Paper~I]{mora2002} who presented the analysis of a large series of high resolution optical spectra of UX~Ori. Many variable absorption events in hydrogen and metallic lines were attributed to the dynamical evolution of gaseous clumps with non metal-rich, roughly solar chemical compositions. In addition, \\citet{beust2001} have shown that the  $\\beta$~Pic infalling planetesimal model would not produce detectable absorptions in typical PMS HAe CS conditions. We also note that dust disks around HAe stars are primordial and can be explained in the context of irradiated PMS CS disk models \\citep{natta2001}.  In this paper we present high resolution spectra of the HAe stars BF~Ori, SV~Cep, WW~Vul and XY~Per and perform an analysis similar to that carried out for UX~Ori \\citepalias{mora2002}. The spectra show very active and complex CS gas in these objects; many transient absorption features in hydrogen and metallic lines are detected, indicating similar properties of the gas around these stars to those of UX~Ori CS gas. In addition, the spectra of WW Vul show metallic features without obvious hydrogen counterparts; in this sense, this star presents a peculiar behaviour. The layout of the paper is as follows: Section~\\ref{observations} presents a brief description of the observations. Section~\\ref{analysis} presents the results and an analysis of the photospheric spectra and the CS contribution. Section~\\ref{discussion} presents a short discussion on the kinematics and strength of the variable features, and on the metallic events detected in WW~Vul. Finally, Sect.~\\ref{conclussions} gives some concluding remarks. ",
        "conclusions": "\\label{conclussions}  We have analyzed optical high resolution spectra of the HAe stars BF Ori, SV Cep, WW Vul and XY Per. These spectra monitor the stars on time scales of months and days and, as in the case of the previously studied UX Ori \\citepalias{mora2002}, they provide observational constraints, which should be considered for any realistic scenario of the gaseous circumstellar disks around intermediate-mass PMS stars. Our results and conclusions can be summarized as follows:  \\begin{enumerate}  \\item The gaseous circumstellar environment of these stars is very complex and active. The spectra always show circumstellar line absorptions with remarkable variations in their strength and dynamical properties.  \\item Variable absorption features are, in most  cases, detected simultaneously in hydrogen and in many metallic lines with similar velocities. In each case, there are several  kinematic components in each line, both blue-shifted and red-shifted with respect to the systemic velocity, denoting the simultaneous presence of infalling and outflowing gas. We attribute the variable features detected in both Balmer and metallic lines to gaseous clumps of solar-like composition, evolving dynamically in the circumstellar disks of these objects. In this respect, the disks around the stars studied in this paper are similar to the UX Ori disk. Following the conclusions of \\citetalias{mora2002} we suggest that these clumps and their dynamical evolution should be investigated in the context of detailed magnetospheric accretion models, similar to those of T Tauri stars.  \\item The star WW Vul is peculiar and behaves differently from the other stars studied in this paper and also from UX~Ori. It is the only star that shows, in addition to events seen both in metallic and hydrogen lines, similar to those observed in the other stars, also transient absorption components in metallic lines that do not apparently have any obvious counterpart in the hydrogen lines. This result, taken at its face-value, would indicate the presence of a metal-rich gas component in the environment of WW~Vul, possibly related to the evaporation of solid bodies. However, any such conclussion is premature. We think that a series of optical spectra with  better time resolution (hours) and longer monitoring (up to around seven days), spectra in the far UV range - to analyze Lyman and metallic resonance lines - and detailed NLTE models of different CS gas environments are essential for further progress and to provide clues on the origin of these apparently metal-rich  events, in terms of their appearance/disapperance statistics, dynamics, metallicity  and nature.  \\end{enumerate}"
    },
    "0402/astro-ph0402087_arXiv.txt": {
        "abstract": "We examine the spectrum of diffuse emission detected in the 17\\arcmin\\ by 17\\arcmin\\ field around \\sgrastar\\ during 625~ks of \\chandra\\ observations. The spectrum exhibits He-like and H-like lines from Si, S, Ar, Ca, and Fe, that are consistent with originating in a two-temperature plasma, as well as a prominent low-ionization Fe K-$\\alpha$ line. The cooler, $kT \\approx 0.8$~keV plasma differs in surface brightness across the image between $(0.2 - 1.8)\\times10^{-13}$~\\ergcmsarcmin\\ (observed, 2--8~keV). This soft plasma is probably heated by supernovae, along with a small contribution from the winds of massive Wolf-Rayet and O stars. The radiative cooling rate of the soft plasma within the inner 20 pc of the Galaxy could be balanced by 1\\% of the kinetic energy of one supernova every $3\\times10^5$ y. The hotter, $kT \\approx 8$ keV component is more spatially uniform, with a surface brightness of $(1.5-2.6)\\times10^{-13}$~\\ergcmsarcmin\\ (observed; 2--8)~keV. The intensity of the hard plasma is correlated with that of the soft, but they are probably only indirectly related, because neither supernova remnants nor WR/O stars are observed to produce thermal plasma hotter than $kT \\approx 3$ keV. Moreover, a $kT \\approx 8$ keV plasma would be too hot to be bound to the Galactic center, and therefore would form a slow wind or fountain of plasma. The energy required to sustain such a freely-expanding plasma within the inner 20 pc of the Galaxy is $\\sim 10^{40}$~\\ergsec. This corresponds to the entire kinetic energy of one supernova every 3000 y, which is unreasonably high. However, alternative explanations for the $kT \\approx 8$ keV diffuse emission are equally unsatisfying. The hard X-rays are unlikely to result from undetected point sources, because no known population of stellar object is numerous enough to the observed surface brightness. There is also no evidence that non-thermal mechanisms for producing the hard emission are operating, as the expected shifts in the line energies and ratios from their collisional equilibrium values are not observed. We are left to conclude that either there is a significant shortcoming in our understanding of the mechanisms that heat the interstellar medium, or that a population of faint ($< 10^{31}$ erg s$^{-1}$), hard X-ray sources that are a factor of 10 more numerous than CVs remains to be discovered. ",
        "introduction": "Bright, diffuse X-ray and $\\gamma$-ray emission has been observed all along the Galactic plane, but is particularly bright toward the Galactic center \\citep[e.g.,][]{wor82,koy86b,yam96,yam97,ski97,vm98}. The origin of this emission is uncertain. Unlike the cosmic X-ray background, the Galactic ridge emission has not yet been resolved into discrete point sources. On the one hand, \\asca\\ observations revealed degree-scale differences in the surface brightness of the diffuse emission that could only be produced by Poisson fluctuations in the numbers of undetected point sources if they have a luminosity of $\\sim 10^{33}$~\\ergsec\\ \\citep{yam96}. On the other hand, \\chandra\\ observations clearly demonstrate that there are not enough discrete sources with $L_{\\rm X} \\gtrsim 10^{31}$~\\ergsec\\ to account for more than 10\\% the Galactic ridge X-ray emission \\citep{ebi01}. However, the strengths of Fe lines observed from the diffuse emission are similar to those observed  from Galactic X-ray point sources, which suggests that they could be one and the same (Wang, Gotthelf, \\& Lang 2002)\\nocite{wgl02}. Furthermore, discrete yet extended sources could produce the spatial variations in the diffuse emission. Several classes of extended features have recently been identified, including: regions of bright iron fluorescence that are ascribed to molecular clouds being illuminated by X-rays from a bright point source \\citep{smp93,mur00} or bombarded by low-energy cosmic-ray electrons \\citep{ylw02}; arcminute-scale features with hard spectra that resemble supernova shocks \\citep{bam03}; and X-ray counterparts to radio features that are thought to be magnetic filaments \\citep[][Lu, Wang, \\& Lang 2003]{sak03}\\nocite{lwl03}. \\chandra\\ and \\xmm\\ observations are only beginning to establish how much flux faint point sources and discrete extended features contribute to the diffuse Galactic X-ray emission.  If the Galactic ridge X-ray emission is truly diffuse, then it could be produced either by hot, $T \\ga 10^7$~K plasma or by cosmic-rays interacting with neutral material in the interstellar medium (ISM). The spectrum of the diffuse emission is one of the most useful diagnostics of its origin. \\asca\\ observations in the 0.5--10~keV band reveal lines from H-like and He-like ions of Mg, Si, S, and Fe, which indicates that the diffuse emission cannot originate from a plasma with a single temperature \\citep{yam96}. As a result, several authors have modeled the diffuse emission as originating from two plasma components \\citep[e.g.,][]{koy96,kan97,tan00}: one with $kT_{\\rm cool} \\approx 0.8$~keV (which we refer to here as the ``cool'' or ``soft'' component), and a second with $kT_{\\rm hot} \\approx 8$~keV (which we refer to as ``hot'' or ``hard''). The soft plasma is thought to be produced by supernova shock-waves \\citep{kan97}, which are the largest source of energy for heating the ISM \\citep{schl02}.  The origin of the $kT \\approx 8$~keV component of the Galactic ridge emission is far less certain. The temperature of the putative hot plasma is too high for it to be bound to the Galactic disk, so that the energy required to sustain it could be equivalent to the release of kinetic energy from one supernova occurring every 30~years \\citep[e.g.,][]{kan97,yam97}. However, supernovae are not observed to produce thermal plasma with $kT \\gtrsim 3$~keV, and there is no known alternative source in the Galactic disk for that much energy. Therefore, several alternative sources of the hot plasma have been proposed. One possibility is that the $kT \\approx 8$ keV plasma is heated by magnetic reconnection in the ISM, and subsequently confined to the Galactic plane by a large-scale toroidal field \\citep{tan99}. It is also possible that the hot diffuse X-ray emission is a low-energy extension of the emission with a power-law spectrum observed above 10~keV, which suggests that it may result from a non-thermal mechanism \\citep[see][but see Lebrun \\etal\\ 2004]{ski97,yam97,vm98}. Among the proposed mechanisms are charge-exchange interactions between cosmic ray ions and interstellar matter (Tanaka, Miyaji, \\& Hasinger 1999; Valinia 2000)\\nocite{tmh99,val00}, bremsstrahlung radiation from cosmic-ray electrons or protons \\citep{vm98}, and quasi-thermal emission from plasma that is continuously accelerated by supernova shocks propagating in the $kT \\approx 0.8$~keV component of the ISM \\citep{dog02,mas02}. These non-thermal processes should produce line emission with energies and flux ratios that differ significantly from those expected from a plasma in thermal equilibrium. Unfortunately, previous \\asca\\ observations were unable to determine the ionization state of the plasma unambiguously, because the spectrum was contaminated by an instrumental Fe line between 6--7~keV.  Clearly, further study of the diffuse X-ray emission is important for understanding stellar life cycles, magnetic fields, and cosmic ray production in the Galaxy. In this paper, we study the spectral properties of the diffuse X-ray emission from a 17\\arcmin\\ by 17\\arcmin\\ field around \\sgrastar\\ that has been observed for over 600~ks with \\chandra. These observations have several advantages over previous ones with \\asca\\ \\citep{koy96} and \\bepposax\\ \\citep{sm99}: (1) the long integration provides sufficiently large signal-to-noise ratio to study the spectrum of the diffuse emission from arcminute-scale sub-regions of the field, (2) the 0\\farcs5 angular resolution allows us to resolve the truly diffuse emission from filamentary features and point sources, and (3) the relative lack of instrumental lines, particularly between 6--7~keV, allows us to measure the ionization state of the diffuse emission with greater confidence. The layout of the paper is as follows. In Section~2.1 we present images that provide an overview of the diffuse emission from the field. In Section~2.2, we examine how the spectrum of the diffuse emission differs across the field. In Section~2.3, we compare the spectra of the point sources and diffuse emission, and place upper limits on the contribution of undetected point sources to the diffuse emission. In Section~3.1, we derive the properties of the putative plasma responsible for the diffuse emission. These are used in Sections~3.2 and 3.3 to examine the likely origins of the diffuse emission. In Section~3.4, we discuss the number of undetected point sources that may be present in the field. Finally, in Section~4, we list the contributions of point sources, extended features, and diffuse emission to the X-ray luminosity of the Galactic center. ",
        "conclusions": "Using a 600~ks exposure with the ACIS-I aboard \\chandra, we have studied the spectrum of diffuse X-ray emission from several regions within a projected distance of 20 pc of \\sgrastar. The spectrum of the diffuse emission  exhibits He-like and H-like lines from Si, S, Ar, Ca, and Fe, as well as a prominent low-ionization Fe line. If the spectrum is modeled as originating from diffuse plasma, two components with temperatures of 0.8~keV and 8~keV are required, along with line emission from low-ionization Fe at 6.4~keV. The energies and flux ratios of the lines from both temperature components are consistent with emission from plasmas in collisional ionization equilibrium.  In Table~\\ref{tab:summary}, we provide a summary of the origins of the X-ray emission between 2--8 keV from the inner 20 pc of the Galaxy. By far the largest contribution to the luminosity of the Galactic center is from diffuse emission. In comparison, detected point sources contribute only 10\\% to the luminosity of the Galactic center, while discrete filamentary features contribute less than 5\\% of the total luminosity of the inner 20~pc of the Galaxy. These results are potentially useful for understanding the origin of diffuse X-ray emission from distant galaxies with quiescent central black holes. However, it is important to note that these observations of the Galactic center are strongly affected by interstellar absorption with a column density of at least $6\\times 10^{22}$ cm$^{-2}$. Therefore, the cool emission with $kT \\la 0.5$ keV that produces most of the 0.5--8.0 keV flux from distant galaxies is obscured at the Galactic center. At the same time, \\chandra\\ observations of other galaxies are not sensitive to the $kT \\approx 8$ keV plasma that dominates the flux we observe from the Galactic center, because this hard emission has a much lower surface brightness than the $kT \\la 1$ keV emission where the \\chandra\\ effective area is largest (0.5--3 keV), and it is difficult to resolve from bright X-ray binaries. Thus, these observations of the Galactic center provide a unique view of the hottest components of the ISM of galaxies.  The properties of the soft, $kT \\approx 0.8$~keV plasma component of the diffuse emission vary significantly across the image, both in temperature between 0.7 and 0.9~keV, and in surface brightness between $2\\times10^{-14}$ and $1.7\\times10^{-13}$~\\ergcmsarcmin. The variation in the properties of the soft plasma suggest that it is relatively young, because differential rotation at the Galactic center should destroy any coherent features within $< 10^6$ y. Supernovae probably supply most of this energy, although the winds from WR and early O stars could also contribute.  Within the inner 20 pc of the Galaxy, the $\\approx 3\\times10^{36}$ erg s$^{-1}$ lost by the plasma through radiative cooling could be replaced by 1\\% of the kinetic energy of one supernova occurring every $\\sim 10^5$ y. The inner 20 pc of our Galaxy contains about 0.1\\% of its total mass, so assuming that one supernova occurs every 100 y in the Galaxy, this rate is roughly consistent with that expected near the Galactic center.  The hard component of the diffuse emission is more spatially uniform than the soft, but the intensities of the two components are still correlated. Although this might suggest a common origin for the two plasma components, supernovae and massive stars are not usually observed to produce plasma with $kT \\gtrsim 3$ keV. This hard emission is distributed throughout the Galactic plane, so it is not likely to be associated with an outburst from \\sgrastar. Instead, the hard emission could result from a $kT\\approx 8$~keV plasma that is heated indirectly by massive stars and supernova remnants, which, for example, could drive reconnection in the magnetic fields near the Galactic center \\citep{tan99}. However, a 8~keV thermal plasma would freely expand away from the Galactic center, and would require $\\approx 10^{40}$ erg s$^{-1}$ to sustain. This is equivalent to the entire kinetic energy of one supernova every 3000 years, which is a much larger rate than usually assumed for supernova. Supernova are the most energetic source of heat for the ISM, so if the hard diffuse emission is produced by a $kT \\approx 8$ keV plasma, it would imply that there is a significant shortcoming in our understanding of heating mechanisms for the ISM.  Alternative explanations for the hard diffuse emission that were intended to lessen the energy required are equally unsatisfying. The suggestion that the hard diffuse emission originates from undetected stellar X-ray sources is unlikely because there is no known class of source that are numerous enough, bright enough, and hot enough to produce the observed flux of $kT \\approx 8$ keV diffuse emission. Likewise, if the hard diffuse emission originates from non-thermal processes, such as the shocks that accelerate cosmic rays, the energies and ratios of the intensities of the line emission should deviate measurably from the values expected for a plasma in thermal equilibrium \\citep[e.g.,][]{mas02}. These deviations are not observed in our \\chandra\\ observations, which presents a challenge to the current non-thermal models.  Further observations should clarify the nature of the diffuse X-ray emission from the Galactic center. X-ray missions with higher spectral resolution, such as ASTRO-E 2, will be able to better-constrain the properties of the putative diffuse plasma by resolving the individual transitions of He-like and H-like Fe, and possibly measuring the velocity dispersions of the Fe ions themselves. Alternatively, a future hard X-ray survey, such as EXIST, could identify a heretofore unknown population of numerous, faint, hard X-ray sources that may be responsible for producing the $kT \\approx 8$ keV diffuse emission."
    },
    "0402/hep-ph0402098_arXiv.txt": {
        "abstract": "We perform calculations of the dependence of nuclear magnetic moments on quark masses and obtain limits on the variation of $(m_q/\\Lambda_{QCD})$ {}from recent measurements of hydrogen hyperfine (21 cm) and molecular rotational transitions in quasar absorption systems, atomic clock experiments with hyperfine transitions in H, Rb, Cs, Yb$^+$, Hg$^+$ and optical transition in Hg$^+$. Experiments  with Cd$^+$,  deuterium/hydrogen, molecular SF$_6$ and Zeeman transitions in $^3$He/Xe are also discussed. ",
        "introduction": "Interest in the temporal and spatial variation of  major constants of physics has been recently revived by astronomical data which seem to suggest a variation of the electromagnetic constant $\\alpha=e^2/\\hbar c$   at the $10^{-5}$ level for the time scale 10 billion years, see \\cite{alpha} (a discussion of other limits can be found in the review \\cite{uzan} and references therein). However, an independent experimental confirmation is needed.  The hypothetical unification of all interactions implies that variation of the electromagnetic interaction constant $\\alpha$ should be accompanied by the variation of masses and the strong interaction constant. Specific predictions need a model. For example, the grand unification model discussed in Ref.~\\cite{Langacker:2001td} predicts that the quantum chromodynamic (QCD) scale  $\\Lambda_{QCD}$ (defined as the position of the Landau pole in the logarithm for the running strong coupling constant) is modified as follows: $\\delta \\Lambda_{QCD} / \\Lambda_{QCD} \\approx 34 \\, \\delta \\alpha / \\alpha$. The variation of quark and electron masses in this model is  given by $\\delta m / m \\sim 70 \\, \\delta \\alpha / \\alpha $. This gives an estimate for the variation of the dimensionless ratio \\begin{equation} \\label{mQCD} {\\delta(m/ \\Lambda_{QCD}) \\over(m/\\Lambda_{QCD})} \\sim \\, 35 {\\delta \\alpha \\over \\alpha} \\end{equation} This result is strongly model-dependent (for example, the coefficient may be an order of magnitude smaller and even of opposite sign \\cite{dent}). However, the large coefficients in these expressions are generic for grand unification models, in which modifications come from high energy scales: they appear because the running strong coupling constant and Higgs constants (related to mass) run faster than $\\alpha$. This means that if these models are correct the variation of masses and the strong interaction scale may be easier to detect than the variation of $\\alpha$.  One can only measure the variation of dimensionless quantities and therefore we want to extract from the measurements the variation of the dimensionless ratio $m_q/\\Lambda_{QCD}$ -- where $m_q$ is the quark mass (with the dependence on the renormalization point removed). A number of limits on the variation of $m_q/\\Lambda_{QCD}$ have been obtained recently from consideration of Big Bang Nucleosynthesis, quasar absorption spectra and the Oklo natural nuclear reactor, which was active about 1.8 billion years ago \\cite{FS,oliv,dmitriev,FS1} (see also \\cite{Murphy1,Cowie,Oklo,c12,savage}). Below we consider the limits which follow from  quasar absorption radio spectra and laboratory atomic clock comparisons. Laboratory limits with a time base of the order one year are especially sensitive to oscillatory variations of fundamental constants. A  number of relevant measurements have been performed already and even larger numbers have been started or are planned. The increase in precision is happening very fast.  It has been pointed out by Karshenboim \\cite{Karschenboim} that measurements of ratios of hyperfine structure intervals in different atoms are sensitive to any variation of nuclear magnetic moments. {}First rough estimates of the dependence of nuclear magnetic moments on $m_q/\\Lambda_{QCD}$ and limits on the variation of this ratio with time were obtained in Ref.~\\cite{FS}. Using H, Cs and Hg$^+$ measurements \\cite{prestage,Cs}, we obtained a limit on the variation of $m_q/\\Lambda_{QCD}$ of about $5 \\cdot 10^{-13}$ per year. Below we calculate the dependence of nuclear magnetic moments on  $m_q/\\Lambda_{QCD}$ and  obtain the limits {}from recent atomic clock experiments with hyperfine transitions in H, Rb, Cs,Yb$^+$,Hg$^+$ and the optical transition in Hg$^+$. It is convenient to assume that the strong interaction scale, $\\Lambda_{QCD}$, does not vary, so we will speak about the variation of masses (this means that we measure masses in units of $\\Lambda_{QCD}$). We shall restore the explicit appearance of $\\Lambda_{QCD}$ in the final answers.  The hyperfine structure constant can be presented in the following {}form \\begin{equation}\\label{A} A=const \\times [\\frac{m_e e^4}{\\hbar ^2}] [ \\alpha ^2 F_{rel}( Z \\alpha)] [\\mu \\frac{m_e}{m_p}] \\end{equation} The factor in the first bracket is an atomic unit of energy. The second ``electromagnetic'' bracket determines the dependence on $\\alpha$. An approximate expression for the relativistic correction factor (Casimir {}factor) for an s-wave electron is the following \\begin{equation}\\label{F} F_{rel}= \\frac{3}{\\gamma (4 \\gamma^2 -1)} \\, , \\end{equation} where $\\gamma=\\sqrt{1-(Z \\alpha)^2}$ and Z is the nuclear charge. Variation of $\\alpha$ leads to the following variation of $F_{rel}$ \\cite{prestage}: \\begin{equation} \\frac{\\delta F_{rel}}{F_{rel}}=K \\frac{\\delta \\alpha}{\\alpha} \\, , \\label{dF} \\end{equation} \\begin{equation}\\label{K} K=\\frac{(Z \\alpha)^2 (12 \\gamma^2 -1)}{\\gamma^2 (4 \\gamma^2 -1)} \\, . \\end{equation} More accurate numerical many-body calculations \\cite{dzuba1999} of the dependence of the hyperfine structure on $\\alpha$ have shown that the coefficient $K$ is slightly larger than that given by this {}formula. For Cs ($Z$=55) $K$= 0.83 (instead of 0.74), {}for Rb $K$=0.34 (instead of 0.29) and finally for Hg$^+$ $K$=2.28 (instead of 2.18).  The last bracket in Eq.~(\\ref{A})  contains the dimensionless nuclear magnetic moment $\\mu$ (i.e., the nuclear magnetic moment $M=\\mu\\frac{e\\hbar}{2 m_p c}$), electron mass $m_e$ and proton mass $m_p$. We may also include a small correction arising from the finite nuclear size. However, its contribution is insignificant.  Recent experiments measured the time dependence of the ratios of the hyperfine structure intervals of $^{199}$Hg$^+$ and H \\cite{prestage}, $^{133}$Cs and $^{87}$Rb \\cite{marion} and the ratio of the optical frequency in Hg$^+$ to the hyperfine frequency of $^{133}$Cs~\\cite{bize}. In the ratio of two hyperfine structure constants for different atoms time dependence may appear from the ratio of the factors $F_{rel}$ (depending on $\\alpha$) as well as from the ratio of nuclear magnetic moments (depending on $m_q/\\Lambda_{QCD}$). Magnetic moments in a single-particle approximation (one unpaired nucleon) are: \\begin{equation}\\label{mu+} \\mu=(g_s + (2 j-1) g_l)/2 \\, , \\end{equation} for $j=l+1/2$. \\begin{equation}\\label{mu-} \\mu=\\frac{j}{2(j+1)}(-g_s + (2 j+3) g_l) \\end{equation} {}for $j=l-1/2$. Here the orbital g-factors are $g_l=1$ for a valence proton and $g_l=0$ for a valence neutron. The present values of the spin g-factors, $g_s$, are $g_p=5.586$ for proton and $g_n=-3.826$ for neutron. They depend on $m_q/\\Lambda_{QCD}$. The light quark masses are only about $1 \\%$ of the nucleon mass ($m_q=(m_u+m_d)/2 \\approx$ 5 MeV) and the nucleon magnetic moment remains {}finite in the chiral  limit, $m_u=m_d=0$. Therefore, one might think that the corrections to $g_s$ arising from the finite quark masses would be very small. However, through the mechanism of spontaneous chiral symmetry breaking, which leads to contributions to hadron properties from Goldstone boson loops, one may expect some enhancement of the effect of quark masses~\\cite{Leinweber:2001ui}. The natural framework for discussing such corrections is chiral perturbation theory and we discuss these chiral corrections next. ",
        "conclusions": ""
    },
    "0402/hep-th0402015_arXiv.txt": {
        "abstract": "We consider the effect of the ultraviolet (UV) or short wavelength modes on the background of Brane Gas Cosmology. We find that the string matter sources are negligible in the UV and that the evolution is given primarily by the dilaton perturbation.  We also find that the linear perturbations are well behaved and the predictions of Brane Gas Cosmology are robust against the introduction of linear perturbations. In particular, we find that the stabilization of the extra dimensions (moduli) remains valid in the presence of dilaton and string perturbations. ",
        "introduction": "Understanding the behavior of strings in a time dependent background has been a subject of much interest and has been pursued in a number of differing ways. One scenario, known as Brane Gas Cosmology (BGC), is devoted to understanding the effect that string and brane gases could have on a dilaton-gravity background in the early Universe \\cite{bv,vafa,bgc,isotropization,stable,extended}. In \\cite{bv}, it was suggested that the energy associated with the winding of strings around the compact dimensions would produce a confining potential for the scale factor and halt the cosmological expansion\\footnote{This was later shown quantitatively in \\cite{vafa}.}. The analysis of BGC was initially performed under the assumption of a homogeneous and isotropic cosmology. The results were recently extended to the case of anisotropic cosmology in \\cite{isotropization}. There, it was shown that string gases can give rise to three dimensions growing large and isotropic due to string annihilation while the other six dimensions remain confined. In \\cite{stable} it was shown that by considering both momentum and winding modes of strings, the six confined dimensions can be stabilized at the self-dual radius, where the energy of the string gas is minimal. This result demonstrated that, in BGC, the volume moduli of the extra dimensions can be stabilized in a natural and intuitive way.  In recent work \\cite{perturbations}, we considered the effect of string inhomogeneities and dilaton fluctuations on BGC.  The string sources of BGC are usually represented by a perfect fluid with homogeneous energy and pressure densities given by the mass spectrum of the strings (see e.g. \\cite{bgc,stable,subodh}). One may worry that inhomogeneities of string sources (in particular strings winding around the confined dimensions) as a function of the unconfined spatial directions could lead to serious instabilities which could ruin the main successes of BGC, namely the prediction that three directions become large leaving the other six confined uniformly as a function of the coordinates of the large spatial sections. In \\cite{perturbations}, we found that at the linear level BGC is robust with respect to long wavelength perturbations. In that paper it was found that at late times the inhomogeneities are subleading compared to the evolution of the background. In this paper we will extend our considerations to the ultraviolet or small wavelength perturbations. Our expectation was that on small wavelengths, the motion of the strings would smear out potential instabilities in a way analogous to how the motion of light particles (``free-streaming'') leads to a decay of short wavelength fluctuations in standard cosmology (see e.g. \\cite{Peebles} for a review). However, we will find that the string matter perturbations are actually sub-leading in the evolution and the dilaton perturbation is the primary driving force of instability.  For reference, in Section 2 and 3 we present the background solution and perturbed equations as found in \\cite{perturbations}.  The crucial new results appear in Section 4, where we derive the perturbation equations for the UV modes and then solve for their late time behavior.  The full equations are presented in the Appendix. We conclude with a discussion of our findings and future prospects in Section 5. ",
        "conclusions": "We have extended the analysis of perturbations in BGC to include the UV modes.  We have derived the evolution equations for the fluctuations at small wavelengths and at late times.  We then solved these equations using a perturbative approach, which we were able to check both analytically and numerically. We find a novel behavior for the perturbations, in that string matter sources are negligible compared with the dilaton perturbation and the resulting behavior is that of a decaying oscillator.  This has interesting consequences in regards to the worry of black hole formation and the usual worrisome behavior of Kaluza-Klein massive states on the background.  We have concluded that at the linear level and in the gas approximation these types of string matter sources will have a negligible effect. Moreover, we find that the predictions of BGC remain robust under the consideration of both long and short wavelength perturbations.  In particular, the prediction that $3+1$ dimensions will grow large while $6$ dimensions remain stabilized around the self dual radius remains intact.  Although these results are promising for BGC there is still much to be done. A more complete treatment of the perturbations would need to take into consideration the non-linear behavior.  It would also be interesting to test the string gas approach itself.  That is, how does one go from the consideration of the effects of individual strings to the known predictions of BGC? Finally, it is an important consideration to reexamine these perturbations in the presence of a frozen dilaton. We know that at very late times in the cosmological evolution the dilaton most likely acquired a mass.  Since the dilaton perturbation played such a vital role in this analysis it could be expected that the results would change dramatically in the massive dilaton case. However, if the perturbations do remain well behaved in this case, it would also be of interest to see if BGC could give rise to a method of structure formation or a unique signature to be observed in the Cosmic Microwave Background.  We leave these questions and concerns to future work."
    },
    "0402/astro-ph0402034_arXiv.txt": {
        "abstract": "{ Using a radiative transfer code (DUSTY) parameters of the circumstellar dust shells of 15 hot post-AGB stars have been derived. Combining the optical, near and far-infrared (ISO, IRAS) data of the stars, we have reconstructed their spectral energy distributions (SEDs) and estimated the dust temperatures, mass loss rates, angular radii of the inner boundary of the dust envelopes and the distances to these stars. The mass loss rates (10$^{-6}-10^{-5}$M$_{\\odot}$yr$^{-1}$) are intermediate between stars at the tip of the AGB and the PN phase. We have also studied the ISO spectra of 7 of these stars. Amorphous and crystalline silicate features were observed in IRAS14331-6435 (Hen3-1013), IRAS18062+2410 (SAO85766) and IRAS22023+5249 (LSIII +5224) indicating oxygen-rich circumstellar dust shells. The presence of unidentified infrared (UIR) band at 7.7$\\mu$, SiC emission at 11.5$\\mu$ and the \"26$\\mu$\" and \"main 30$\\mu$\" features in the ISO spectrum of IRAS17311-4924 (Hen3-1428) suggest that the central star may be carbon-rich. The ISO spectrum of IRAS17423-1755 (Hen3-1475) shows a broad absorption feature at 3.1$\\mu$ due to C$_{2}$H$_{2}$ and/or HCN which is usually detected in the circumstellar shells of carbon-rich stars. ",
        "introduction": "In the evolution of low and intermediate mass stars (0.8 $-$ 8M$_{\\odot}$), the post-asymptotic giant branch (post-AGB) or protoplanetary nebula (PPN) phase is a transition stage from the tip of the AGB to the planetary nebula (PN) stage (Kwok, 1993). The hot post-AGB stars form an evolutionary link between the cooler G,F,A supergiant post-AGB stars (Parthasarathy \\& Pottasch, 1986) and the hotter O-B central stars of PNe (Parthasarathy, 1993a).  Analysis of the UV(IUE) spectra of hot post-AGB stars (Gauba \\& Parthasarathy, 2003), revealed that in many cases, the hot (OB) central stars of PPNe are partially obscured by circumstellar dust shells. Stars on the AGB and beyond are characterised by severe mass loss (10$^{-8}$ $-$ 10$^{-3}$ M$_{\\odot}$ yr$^{-1}$) which results in the formation of circumstellar envelopes. The physical mechanisms responsible for the intensive mass loss from AGB stars are not well understood although the most promising mechanism to date involves radiation pressure on the dust grains (Tielens, 1983). While AGB stars appear to have spherically symmetric dust outflows (eg. Habing \\& Blommaert, 1993), PN tend to have axially symmetric inner regions and spherical outer halos (eg. Schwarz et al., 1992). Inorder to understand the mass loss mechanisms, wind velocities and time scales responsible for the evolution of PNe, we need to study the circumstellar environment of the stage intermediate between the AGB and the PN phase, i.e. the post-AGB/PPN phase. Circumstellar dust shells of some cooler post-AGB stars (eg. Hoogzaad et al., 2002; Hony et al., 2003) and PNe (eg. Siebenmorgen et al., 1994) have been modelled to derive the dust composition, mass loss rates and dynamical ages.  As a consequence of dredge-up of byproducts of helium burning to the surface of stars on the AGB, the oxygen-rich atmospheres of some of these stars may be transformed into carbon-rich atmospheres (see eg.  Iben \\& Renzini, 1983). This change of chemistry would also be reflected in the composition of the dust grains formed in the cirumstellar envelopes of AGB and post-AGB stars. With the resolution and wavelength coverage of the ISO mission (Kessler et al.,1996) the detection of prominent gas and solid state features specific to oxygen-rich and carbon-rich chemistries became possible. Amorphous and crystalline silicate features and crystalline water have been reported in the ISO spectra of some AGB and post-AGB stars and the nebulae surrounding [WC] central stars of PNe (see eg. Waters \\& Molster, 1999; Hoogzaad et al., 2002). Hrivnak et al. (2000) detected the \"21$\\mu$\" and \"30$\\mu$\" emission features besides the unidentified infrared (UIR) emission bands at 3.3, 6.2, 7.7 and 11.3$\\mu$ in the ISO spectra of a sample of carbon-rich PPNe.  Bogdanov (2000, 2002, 2003) modelled the complete spectral energy distribution (SED) of three hot post-AGB stars, IRAS18062+2410 (SAO85766), IRAS19590-1249 (LSIV-12 111) and IRAS20462+3416 (LSII+34 26) using radiative transfer codes and derived their mass loss rates, inner radii of the dust envelopes, optical depth of the envelopes and the distances to these stars. We need to study a bigger sample of such stars to understand the evolution of the infrared spectrum as the stars evolve from the cooler post-AGB phase to the hot central stars of PNe. In this paper, we have used the radiative transfer code, DUSTY (Ivezi\\'c et al., 1999) to model the circumstellar dust shells of 15 hot post-AGB stars. Additionally, 7 stars from our list were found to have ISO spectra. We also report the analysis of the ISO spectra of these stars. ",
        "conclusions": "We have modelled the circumstellar dust shells of 15 hot post-AGB stars using the radiative transfer code, DUSTY and derived their dust temperatures, distances to the stars, mass loss rates and angular radii of the inner boundary of the dust envelopes (Tables 5a and b). These stars have detached dust shells (as is evident from the SEDs, Fig. 3), OB-giant or supergiant spectra and cold dust between 100$-$315K, satisfying the observational properties of PPNe as defined by Kwok (1993, 2001). In addition to the cold dust, warm dust was detected in the case of IRAS12584-4837 (Hen3-847) and IRAS17423-1755 (Hen3-1475) indicating ongoing mass loss. From the grain types used for the model fits, we may infer the chemical composition of the circumstellar dust shells. The use of both silicate and amorphous carbon grains to model the SEDs of IRAS12584-4837 (Hen3-847) and IRAS17460-3114 (SAO 209306) suggests that the central stars in these two cases may have undergone a recent change from an oxygen-rich to a carbon-rich chemistry. Such hot post-AGB stars may evolve into the [WC] central stars of PNe. Recently, Waters et al. (1998) detected carbon-rich PAH features in the near-infrared and crystalline silicates in the far-infrared ISO spectra of two PNe with [WC] central stars, BD+30 3639 and He2-113.  Observational evidence (eg. Chu et al., 1991) suggests that three winds are involved in stripping the outer envelope of the AGB star on its way to becoming a PN (Marten et al., 1993; Frank, 1994) : the spherically symmetric AGB wind (eg. Habing \\& Blommaert, 1993) when the star loses mass at rates of $10^{-7}-10^{-6}$M$_{\\odot}$yr$^{-1}$ with a wind velocity of $\\sim$ 10 kms$^{-1}$; the superwind phase when the mass loss is thought to increase dramatically at the end of the AGB, upto $10^{-5}-10^{-3}$M$_{\\odot}$yr$^{-1}$, still with a wind velocity of $\\sim$ 10 kms$^{-1}$; once the superwind exhausts most of the AGB star's envelope, a fast wind with mass loss rate of $10^{-8}$M$_{\\odot}$yr$^{-1}$ and velocity of $\\sim$ 1000 kms$^{-1}$ develops at some point during the PPN phase. Velocities of 1000 kms$^{-1}$ and mass loss rates of $\\sim$ $10^{-8}$M$_{\\odot}$yr$^{-1}$ have been observed in the central stars of PNe (eg. Gauba et al., 2001). For our hot post-AGB stars, we derived mass loss rates of 10$^{-5}-10^{-6}$M$_{\\odot}$yr$^{-1}$. The mass loss rates ($\\dot M$) scale with the gas-to-dust mass ratio (r$_{gd}$). We have adopted r$_{gd}$ = 200. For carbon-rich AGB and post-AGB stars values between 200 and 250 are often used (eg. Jura, 1986; Meixner et al., 1997). For the cool (F3Ib) post-AGB star, HD161796 (Parthasarathy \\& Pottasch, 1986) with an oxygen-rich circumstellar environment, Hoogzaad et al. (2002) estimated r$_{gd}$ = 270.  Furthermore, our models assume that the dust density distribution falls off as y$^{-2}$ in the entire circumstellar dust shell. Such an assumption would break down in the case of episodic mass loss (Olofsson et al., 1990). Eg. episodic mass loss may have been responsible for the rapid evolution (30 $-$ 40 years) of IRAS17119-5926 (Hen3-1357) and IRAS18062+2410 (SAO85766) from B-type post-AGB supergiants to young PNe (Parthasarathy et al., 1993c, 1995; Bobrowsky et al., 1998, Parthasarathy et al., 2000b).  The proper motions ($\\mu$) of the stars from the Tycho-2 Catalogue (Hog et al., 2000) have been listed in Tables 5a and b. Using the derived distances (d) in conjunction with the proper motions we estimated the component of the stellar space velocities of the targets tangent to the line of sight (V$_{\\rm T}$). For IRAS17203-1534, IRAS18062+2410 (SAO85766) and IRAS18371-3159 (LSE63), the large V$_{\\rm T}$ values (Tables 5a and b) imply very high space velocities (V$_{\\rm s}$ = (V$_{\\rm T}^{2}$ $+$ V$_{\\rm r}^{2}$)$^{1/2}$; where V$_{\\rm r}$ is the radial velocity of a star), close to the escape velocity from the Galaxy of 290 kms$^{-1}$ near the Sun. Mooney et al. (2002) estimated a distance of 8.1 kpc to IRAS18062+2410 (SAO85766) which is much greater than our estimate of $\\sim$ 5 kpc. Such a large distance, if correct, would imply a still higher space velocity. We believe our distance estimates to be closer to the actual values for these stars. However, the assumption of spherical density distributions in our models, may be an over simplification for some of these objects. Eg. IRAS17423-1755 (Hen3-1475) has IRAS colors moderately close to those of HD233517. HD233517 is unresolved and may have a disk instead of a spherical outflow (see eg. Jura, 2003, Fisher et al., 2003). The predicted angular sizes of the inner radii of the dust shells (Tables 5a and b) suggests that these objects should be easily resolvable in the mid-IR images with large ground based telescopes. Imaging of these objects in the IR would serve to test the basic assumptions such as those of spherical symmetry for our models.  In Table 6 we have compared the predicted and observed (V$-$J) values for these stars ($\\Delta$(V-J)=(V-J)$_{\\rm predicted}$-(V-J)$_{\\rm obs}$) where, (V-J)$_{\\rm predicted}$=(V-J)$_{\\rm o}$ $+$ A$_{\\rm V}$ $-$ A$_{\\rm J}$. The intrinsic (V-J) colors ((V-J)$_{\\rm o}$), for the spectral types of the stars are from Ducati et al. (2001). For stars with emission lines in their optical spectra, IRAS12584-4837 (Hen3-847), IRAS17423-1755 (Hen3-1475), IRAS22023+5249 (LSIII+5224)and for the PN, IRAS22495+5134 (LSIII+5142), (V-J)$_{o}$ could not be assumed. In the case of IRAS14331$-$6435 (Hen3-1013), IRAS17203-1534, IRAS18062$+$2410 (SAO85766) and IRAS18379-1707 (LSS5112), we find significant differences between the values of (V-J)$_{\\rm predicted}$ and (V-J)$_{\\rm obs}$. On first sight, this would then raise a suspicion about the adopted E(B$-$V)$_{\\rm total}$ values. However, we would like to point out that the V and J magnitudes of these stars have not been recorded simultaneously. Many of these stars are variable as evidenced from the J,H,K data on IRAS12584-4837 (Hen3-847), IRAS18062+2410(SAO85766) and IRAS18379-1707 (LSS5112). The V and J magnitudes may have be recorded at different epochs of the variability cycle and hence it may not be suitable to compare (V-J)$_{\\rm predicted}$ and (V-J)$_{\\rm obs}$.  We also studied the ISO spectra of 7 hot post-AGB stars, IRAS14331-6435 (Hen3-1013), IRAS16206-5956 (SAO243756), IRAS17311-4924 (Hen3-1428), IRAS17423-1755 (Hen3-1475), IRAS18062+2410 (SAO85766), IRAS22023+5249 (LSIII +5224) and IRAS22495+5134 (LSIII +5142). A weak amorphous silicate feature (10.8$\\mu$) alongwith crystalline silicate features was found in the dust shells of IRAS14331-6435 (Hen3-1013) and IRAS22023+5249 (LSIII +5224). The 17.6$\\mu$ amorphous silicate feature was missing in these two stars. The post-AGB star IRAS18062+2410 (SAO85766) did not show evidence for the presence of crystalline silicates but strong amorphous silicate features at 10.8$\\mu$ and 17.6$\\mu$ were detected. Volk \\& Kwok (1989) predict that at dust temperatures of typically a few 100K for post-AGB stars, the spectrum should increase from 8 to 23$\\mu$ and the 10.8$\\mu$ and 17.6 $\\mu$ silicate features should cease to be observable. This appears to be consistent with the observed spectral features and the dust temperatures of 230K, 130K and 120K for IRAS18062+2410 (SAO85766), IRAS14331-6435 (Hen3-1013) and IRAS22023+5249 (LSIII +5224) respectively from our model fits. The presence of silicate features in these stars indicates the O-rich nature of the central stars. The formation of crystalline silicates in the circumstellar shells of post-AGB stars is still not well understood (see eg., Waters et al., 1996). In contrast, PAH emission at 7.7$\\mu$, the \"26$\\mu$\" and \"main 30$\\mu$\" features and 11.5$\\mu$ SiC emission in IRAS17311-4924 (Hen3-1428), are typical of circumstellar dust shells around carbon-rich post-AGB stars. However, the 21$\\mu$ emission feature detected in several carbon-rich PPNe (Hrivnak et al., 2000) was notably absent in the ISO spectrum of IRAS17311-4924 (Hen3-1428). Volk et al. (2002) pointed out that although all sources with the 21$\\mu$ emission feature also display the \"26$\\mu$\" and \"main 30$\\mu$\" features, the converse is not true. The hot post-AGB star, IRAS01005+7910 (Klochkova et al., 2002) which showed the \"26\" and \"main 30$\\mu$\" emission also did not show the 21$\\mu$ emission feature (Hrivnak et al., 2000). It may be that the dust grains responsible for the 21$\\mu$ emission are destroyed as the central star evolves towards hotter temperatures. The broad absorption feature at 3.1$\\mu$ in IRAS17423-1755 (Hen3-1475) attributed to C$_{2}$H$_{2}$ and/or HCN indicates that the central star may be carbon-rich."
    },
    "0402/astro-ph0402202_arXiv.txt": {
        "abstract": "We discuss a bias present in the calculation of the global luminosity function (LF) which occurs when analysing faint galaxy samples.  This effect exists because of the different spectral energy distributions of galaxies, which are in turn quantified by the $k$-corrections. We demonstrate that this bias occurs because not all galaxy types are visible in the same absolute magnitude range at a given redshift and it mainly arises at high redshift since it is related to large $k$-corrections. We use realistic simulations with observed LFs to investigate the amplitude of the bias.  We also compare our results to the global LFs derived from Hubble Deep Field-North and -South (HDF) surveys.  We conclude that, as expected, there is no bias in the global LF measured in the absolute magnitude range where all galaxy types are observable. Beyond this range the faint-end slope of the global LF can be over/under-estimated depending on the adopted LF estimator. The effect is larger when the reference filter in which the global LF is measured, is far from the rest-frame filter in which galaxies are selected.  The fact that LF estimators are differently affected by this bias implies that the bias is minimal when the different LF estimators give measurements consistent with one another at the faint-end.  For instance, we show that the estimators are discrepant in the same way both in the simulated and HDF LFs. This suggests that the HDF LFs are affected by the presently studied bias.  The best solution to avoid this bias is to derive the global LF in the reference filter closest to the rest-frame selection filter. ",
        "introduction": "The luminosity function (LF) is a fundamental and basic tool to understand and constrain the history of galaxy formation and evolution.  Moreover, the derived mean luminosity density at different redshifts allows to derive estimates of the cosmic star formation density.  In the distant Universe, LFs are measured in se\\-ve\\-ral redshift bins in order to quantify the evolution of galaxy populations.  In this paper, we focus on the relia\\-bi\\-lity of the statistical estimators usually used to measure the global LF.  We call {\\it global LF} the sum of the LFs per galaxy type \\citep{1988ARA&A..26..509B}.  Calculating LFs is not a trivial task since estimators must account for all biases or limits introduced by the observational selection effects. Most of the surveys are limited in apparent magnitude. This effect is accounted for in the 1/V$_{\\rm max}$ LF estimator \\citep{1968ApJ...151..393S}.  The drawback of the 1/V$_{\\rm max}$ method is the implicit assumption, in its formulation, of a uniform galaxy distribution (i.e. no significant over- or under-densities of galaxies). Nevertheless, because of its simplicity, this method is the most often used in high-redshift surveys.  \\citet{1971MNRAS.155...95L} developed the C$^{-}$ method to overcome the assumption of a uniform galaxy distribution.  The STY \\citep*{1979ApJ...232..352S} and the Step Wise Maximum-Likelihood LF estimators, hereafter SWML, \\citep*[][EEP]{1988MNRAS.232..431E} are both related to maximum-likelihood statistical methods. The C$^{-}$, STY and SWML methods make no assumptions on spatial distribution of galaxies, but the information about the normalization of the LF is lost. \\citet{1982ApJ...254..437D} reviewed various estimators to derive the normalization. In contrast to C$^{-}$ and SWML, the STY method does assume a parametric form to the luminosity distribution.  All LF estimators present both advantages and drawbacks. \\citet{1997AJ....114..898W} and \\citet*{2000ApJS..129....1T} compared several LF estimators using simulated ca\\-talogues.  Their mock catalogues did not tackle into detail the effects of $k$-corrections and of the mix of individual and different LF shapes for different morphological types in the measurement of the global LF. The Canada-France Redshift Survey \\citep[CFRS;][]{1995ApJ...455..108L} demons\\-trated that the evolution of the LF depends strongly on the studied galaxy population. In this paper we add the dependency of limiting absolute magnitudes on galaxy type in si\\-mu\\-lated catalogues, and at the same time we introduce an evolution of the LF per galaxy population to produce realistic simulations. These improvements enable us to identify an intrinsic bias in the estimators to measure the global LF.  The different visibility limits for the various galaxy types (mainly due to different $k$-corrections) affects all flux-limited surveys. Hence it can have an impact on statistical analyses, in particular the LF estimates.  The accepted idea is that certain galaxy types sometimes can not be visible in a given redshift bin, so that it would ob\\-vi\\-ous\\-ly underestimate the global LF.  However even though all galaxy types are visible in a given redshift bin, we show using realistic simulations that a bias still arises in the measurement of the global LF.  As noted by \\citet{1995ApJ...455..108L}, it occurs because different galaxy types are not visible in the same absolute magnitude range.  In the literature this bias has never been quantified. We use real and simulated data to investigate the amplitude and the behavior of this bias. In particular our analysis is focused on high-redshift data since the $k$-correction values are small at low redshift.  One possible solution to avoid this bias would be to sum the extrapolated LFs per galaxy type to measure the global LF. Unfortunately this solution is hazardous in the highest redshift bins of a deep survey for two reasons: the number of galaxies is often too small to derive LFs per galaxy type and not all the LF slopes per galaxy type are well constrained, which would imply a dangerous extrapolation.  The analysis of the global high-redshift LFs has been the framework of most of the previous analyses on deep surveys like, for instance, the Subaru Deep Field \\citep[SDF;][]{2003AJ....125...53K}, the Hubble Deep Fields (HDF; e.g. for example \\citealt*{1997AJ....113....1S}, \\citealt{2000ApJS..129....1T}, \\citealt*{2002A&A...395..443B}), the CFRS \\citep{1995ApJ...455..108L}.  With the on-going or earlier deep surveys, one needs to quantify in details this bias related to $k$-correction effects.  This paper is organized as follows.  Section~2 describes the origin of the bias linked to the spectral energy distribution dependency of absolute magnitudes. Section~3 reviews briefly the following estimators, 1/V$_{\\rm max}$, STY, SWML and C$^{+}$, and the bias linked to each of them. Section~4 quantifies the impact of the bias on the global LF from simulations and from the Hubble Deep Field surveys. Section~5 presents our conclusion.  Throughout this paper, we adopt an Einstein-de Sitter universe ($\\Omega_{\\rm 0} = 1$, $\\Omega_{\\rm \\Lambda} = 0$) and H$_{\\rm 0} = 100$~km~s$^{-1}$~Mpc$^{-1}$, but the results here discussed are not dependent on the adopted cosmological model. ",
        "conclusions": "Our study enabled us to describe when LF estimators are robust for the measurement of the global LF in the framework of the earlier and future deepest surveys. We demonstrated that the estimation of the global LF contains an intrinsic bias due to the fact that, in a magnitude limited sample, different galaxy types have different limits in absolute magnitude because of different k-corrections. The importance of the effect is larger when the range of k-correction between the different galaxy types is wide. For this reason this bias mainly arises in high redshift samples.  The STY and SWML estimators are not affected in the same way by this bias as the 1/V$_{\\rm max}$ and C$^{+}$ estimators.  If the STY, SWML and the 1/V$_{\\rm max}$, C$^{+}$ methods are not in good agreement with each other, this is an indication that the bias in the global LF estimators is present. A good indication of the presence of a significant bias is when the differences between different estimators (Vmax and STY for instance) is larger than the statistical uncertainties (Poisson errors for instance).  We quantified it using realistic simulations and observations for galaxies selected in the $I$ filter, and measuring the LF in various reference filters (UV, B, I). We obtain the following results. \\begin{itemize} \\item[(i)] Case $1+z_{low} < \\lambda^S/\\lambda^{Ref}$ (e.g., a reference-frame UV LF for galaxies selected in I): the studied estimators underestimate the faint-end slope of the global LF for $z_{low} \\la 2 $. This underestimate is particulary significant for the 1/V$_{\\rm max}$ and C$^{+}$ methods (i.e. for instance, the $UV$-LF of the SDF). \\item[(ii)] Case $1+z_{low} \\sim \\lambda^S/\\lambda^{Ref}$ (e.g., a reference-frame B LF for galaxies selected in I): the estimators of the global LF are robust up to $z_{low} \\la 1.3$. In this redshift range the bias is minimal (i.e. for instance, the CFRS case). \\item[(iii)] Case $1+z_{low} > \\lambda^S/\\lambda^{Ref}$ (e.g., a reference-frame I LFs with galaxies selected in I): the STY and SWML methods overestimate the faint-end slope of the global LF, while the 1/V$_{\\rm max}$ method roughly recovers well the global LF (e.g., for instance, the redshift bin [1.25, 2] of the HDF). \\end{itemize}  The ways to reduce the intrinsic bias of the global LF estimators are the following: \\begin{itemize} \\item[(a)] The selection of galaxy subsamples in the closest rest-frame filter to the reference filter in which the LF is measured \\citep[see e.g.][2003]{2001ApJ...551L..45P}.  This method is also the best to reduce the SED dependency in the measurement of absolute magnitudes since in this case the term [color+$k$-correction] is little dependent on the SED.  Only multi-color surveys allow to derive the same rest-frame band LF at different redshifts using this strategy. \\item[(b)] In principle, the estimate of the global LF using the sum of the extrapolated LF per galaxy type.  However it requires a good knowledge of the slope for all the LFs per type, and in practice the use of extrapolated LFs may be hazardous. \\item[(c)] The estimate of the global LF using a filter in which the differences between $k$-corrections are small, as for instance in the $K$-filter, e.g. \\cite{2002A&A...395..443B}, \\cite{poz}. \\item[(d)] The estimate of the global LF within an absolute magnitude range in which all galaxy types are detected \\citep[see e.g.][]{1997ApJ...487..512S}. This method is appropriate for very large surveys like the VVDS for instance, at the cost of the loss of the faintest bins of the global LF. \\end{itemize}"
    },
    "0402/astro-ph0402491_arXiv.txt": {
        "abstract": "{Nonlinear growth of one-dimensional density structures with a frozen-in magnetic field is investigated in Newtonian cosmology. A mechanism of magnetic field amplification is discussed. We discuss the relation between the initial conditions for the velocity field and the basic time-scales determining the growth of the magnetized structure.} ",
        "introduction": "During the last decade the evidence large-scale cosmological magnetic fields has systematically grown. Fields of  several microgauss have been measured beyond the galaxy clusters (Kim et al. 1991); the Rotation Measure (Kronberg 1994) confirms the existence of coherent magnetic field on Mpc scales or larger. The recent discovery of large-scale diffuse radio emission testifies to the presence of magnetic fields of $\\sim 0.1~\\mu $G, along the 6 Mpc filament (Bagchi et al. 2002), with evidence for their coherent nature.  Understanding of the magnetic field behaviour at different phases of the matter-dominated era is crucial to explain the microgauss fields observed in high redshift objects  $z \\geq 2$; starting from the damped Ly $\\alpha$ systems, through the distant radiogalaxies and the galaxy clusters up to the scales typical of galaxy superclusters. While the magnetic field of galaxies may result from dynamo effects, the fields at larger scales cannot be explained by the same mechanism.  Here there is either no rotation, necessary for dynamo action, or the structures  are dynamically too young to leave dynamo action enough time to operate. Magnetic fields at Mpc scales are likely to be primordial. Some pre-dynamo mechanisms of the primeval magnetic field amplification must be at work at least in the linear and {\\wnr}.  In this paper we discuss  the mutual relationships between density growth and the magnetic field evolution in early nonlinear stages, when the wall or filamentary structure is formed. We investigate whether the growing planar density structure may drag and amplify the magnetic field.  We employ the exact solutions for 2-D (\\pan) inhomogeneity evolution in the Newtonian description\\footnote{The problem is opposite to that formulated by Wasserman (1978) and Kim et al.(1996), where the magnetic fields are expected to actively support the structure formation processes. } and emphasize the role of initial conditions, in particular, the large-scale primordial flows. The magnitude of primordial velocity fields at recombination determines the time and the growth rate of density fluctuation. As a consequence, it defines the duration time of the pre-dynamo and dynamo amplification phase. To avoid problems with a mathematical definition of the {\\wnr} we work with fully nonlinear equations and their solutions. Although finally we refer to the regime where the density contrast $\\x$ is between $1$ and $100$ (which is relevant for the cosmological structures we discuss), dynamical equations are true for $\\x>100$.  Magnetohydrodynamic equations in the covariant notation are given in Section~\\ref{magneto}. Simplifying  physically relevant assumptions and the resulting nonlinear perturbation equations are discussed in Section~\\ref{planar}. Nonlinear solutions for the density contrast and the magnetic field enhancement are given in Section~\\ref{evolution}. Section~\\ref{numerical}  contains numerical estimations and graphical presentation of the magnetized \\pan formation. ",
        "conclusions": "We provided  the nonlinear exact solutions for the density,  velocity and magnetic fields for the \\pan-type structures in the Newtonian expanding universe. The approximations of the potential velocity field and vanishing matter pressure have been employed. The time when the compressing flat structure enters the regime of nonlinear growth is controlled by the initial value of velocity field at the recombination. The structures accompanied by large hydrodynamic flows collapse earlier, i.e. the moment when dynamo mechanism may switch on occurs at higher redshifts, which eventually results in stronger magnetic field enhancement. For presently observed velocity fields, $10^{-1}$, (see, e.g. Dekel 1997) in supergalactic structures of $100$~Mpc and $\\x \\sim 1$ the initial inflows $\\y_\\ini \\simeq 10^{-3}$ and initial magnetic fields $H_\\ini\\simeq 10^{-9} - 10^{-8}$~Gauss are expected. The result is compatible with the simulation estimations (Gramann et al., 2002).  Firm evidence of primordial magnetic fields in structures at pre-virial stages is of particular importance, as these fields \"remember\" the initial conditions and thus set constraints on the seed magnetic fields, density and velocity fields at the recombination. The observational techniques become more important (Faraday rotation measurements and the indications coming from the propagation of cosmic radiation UHE in the Local Supercluster), which potentially might distinguish between the large scale magnetic seed component from other magnetic fields of astrophysical origin (i.e. resulting from galactic dynamo, outflows from radiogalaxies etc.). The rotation measure which have the same sign along the Supercluster plane would suggest a coherent, relic field at this scale."
    },
    "0402/astro-ph0402172_arXiv.txt": {
        "abstract": "We present the peculiar near-infrared spectrum of the newly discovered brown dwarf 2MASS~J05185995$-$2828372, identified in the Two Micron All Sky Survey.  Features characteristic of both L and T dwarfs are present, namely strong carbon monoxide absorption in $K$-band, strong methane absorption in $J$- and $H$-bands, and red near-infrared colors.  We consider several scenarios that could produce these features and conclude that the object is most likely to be an unresolved L/T binary system.  We discuss how the estimated photometric properties of this object are consistent with the observed $J$-band brightening of brown dwarfs between late-L and early-T dwarfs, making detailed study of this system an important probe of the L/T transition. ",
        "introduction": "\\label{sec:intro}  The existence of brown dwarfs, low-mass (M~$\\la0.075$~M$_{\\sun}$) objects that form like stars but are incapable of maintaining core hydrogen fusion, was first postulated by \\citet{Kumar}. After several decades of unsuccessful searches, well over 100 brown dwarfs are currently known, primarily as a consequence of the availability of deep far-red and infrared sky surveys \\citep[DENIS, 2MASS,SDSS]{DENIS,2MASS,SDSS}.  Without a long-lived energy source, brown dwarfs cool rapidly, exhibiting spectra dominated by a sequence of complex molecular bands. Metal hydride absorption (e.g., FeH, CrH, and CaH) replaces titanium oxide at optical wavelengths as the effective temperature falls below 2100 K, and the spectral type evolves from type M to type L \\citep{K99,Chabrier}. As the temperature drops below 1300 K, methane forms in the outer atmosphere \\citep{Tsuji64} and the strong absorption at 1--3~$\\micron$ leads to significantly bluer near-infrared colors. These are T dwarfs \\citep{B02,Geballe}.  We are currently using near-infrared photometry from the Two Micron All-Sky Survey \\citep[2MASS]{2MASS} to search for all late-type M and L dwarfs  lying within 20 parsecs of the Sun \\citep{paper5}. In the course of follow-up observations, we have identified a cool dwarf which appears to break the current spectral classification paradigm. 2MASS~J05185995$-$2828372 (hereafter, 2M~0518) was selected for observation based on its red ($J-K_s$) color of 1.82 magnitudes and its relatively bright apparent magnitude, $J=15.98$.  The peculiar near-infrared spectrum of this object, however, exhibits both L and T dwarf spectral features.  The following section describes our observations and \\S~\\ref{sec:discussion} discusses possible explanations for the observed properties of this intriguing object.  \\begin{figure}[b] \\centering \\includegraphics[width=3.25in]{f1} \\caption{Color-color diagram for 2M~0518 (\\emph{five-pointed star}), late-type stars (\\emph{triangles}), L dwarfs (\\emph{crosses}), and T dwarfs (\\emph{circles}).}\\label{fig:color} \\end{figure} ",
        "conclusions": ""
    },
    "0402/astro-ph0402458_arXiv.txt": {
        "abstract": "{We present imaging circular polarimetry and near-infrared photometry of the suspected ultra-short period white-dwarf binary RX\\,J0806.3+1527 obtained with the ESO VLT and discuss the implications for a possible magnetic nature of the white dwarf accretor and the constraints derived for the nature of the donor star. Our $V$-filter data show marginally significant circular polarization with a modulation amplitude of $\\approx 0.5$\\% typical for cyclotron emission from an accretion column in a magnetic field of order 10\\,MG and not compatible with a direct-impact accretor model. The optical to near-infrared flux distribution is well described by a single blackbody with temperature $kT_{\\rm bb} = 35000$\\,K and excludes a main-sequence stellar donor unless the binary is located several scale heights above the galactic disk population. }    \\addkeyword{Polarization} \\addkeyword{Stars: binaries: close} \\addkeyword{Stars: individual: RXJ0806.3+1527} \\addkeyword{Stars: late-type} \\addkeyword{Stars: magnetic fields}   \\begin{document} ",
        "introduction": "\\label{sec:intro}  The soft X-ray discovered system RX\\,J0806.3 +1527 (Beuermann et al. 1999) has recently been suggested to be a semidetached white dwarf binary with a helium-degenerate secondary and the shortest known orbital period Israel et al. (2002). If the 321 s pulse period (Israel et al. 1999, Burwitz \\& Reinsch 2001) is indeed the orbital period this system would lie close to the theoretical minimum period of white dwarf binaries and would be a suitable candidate for gravitational wave detection.  Three flavors of the double-degenerate model (polar, direct accretor, electric star) have been advanced to account for the observational characteristics (Cropper et al. 1998, Marsh \\& Steeghs 2002, Wu et al. 2002). Alternatively, an interpretation of \\rxj{} as a face-on, stream-fed intermediate polar has been advocated (Norton et al. 2003).  We present the results of our recent imaging circular polarimetry and near-infrared photometry and discuss the implications for a possible magnetic nature of the white dwarf accretor and the constraints derived for the nature of the donor star. ",
        "conclusions": "\\label{sec:discussion}  The circular polarization detected in the $V$-filter data is comparable to that seen in some intermediate polars (e.g. Buckley et al. 1995). This is characteristic of cyclotron emission originating from an accretion column in a magnetic field of order 10\\,MG and would be compatible with all currently discussed models except for the direct accretor.  The optical-to-infrared flux distribution excludes a main-sequence stellar donor unless the binary is located several scale heights above the galactic disk population. If the system is at a distance of 1\\,kpc (i.e. at approximately twice the galactic scale height) the secondary must be of spectral type L8 or later that its flux contribution can be hidden in the observed spectrum. This is difficult to conceal in an intermediate polar interpretation of \\rxj{} which implies a Roche-lobe filling stellar companion.  On the other hand, the question whether other periods are present in the system (as expected for an intermediate polar) is not well settled and needs further observations. Furthermore, the evidence for Hydrogen in the optical spectrum of \\rxj{} implies severe problems for all ultra-short period binary models which require a Helium rich or a Helium degenerate secondary.  Concluding, none of the models proposed so far fits well with all available observations and the true nature of RX\\,J0806.3+1527 must still be considered open."
    },
    "0402/astro-ph0402635_arXiv.txt": {
        "abstract": "{ We report on multi-epoch HST/WFPC2 images of the XZ~Tauri binary, and its outflow, covering the period from 1995 to 2001. Data from 1995 to 1998 have already been published in the literature. Additional images, from 1999, 2000 and 2001 are presented here. These reveal not only further dynamical and morphological evolution of the XZ~Tauri outflow but also that the suspected outflow source, XZ~Tauri North has flared in EXor-type fashion. In particular our proper motion studies suggests that the recently discovered bubble-like shock, driven by the the XZ~Tauri outflow, is slowing down (its tangential velocity decreasing from 146\\,km\\,s$^{-1}$ to 117\\,km\\,s$^{-1}$). We also present simulations of the outflow itself, with plausible ambient and outflow parameters, that appear to reproduce not only the dynamical evolution of the flow, but also its shape and emission line luminosity. ",
        "introduction": "XZ~Tau is a classical T-Tauri binary system with a separation of 0$\\farcs$3 (\\cite{Haas90}) and is located in the well known Lynds 1551 star-forming region some 140\\,pc away (\\cite{Elias78}). The system was first found to have an associated Herbig-Haro (HH) outflow through ground-based CCD imaging and spectroscopy (\\cite{Mundt88}; \\cite{Mundt90}). These early observations revealed a bipolar optical flow that could be traced to at least 10$\\arcsec$ on either side of the binary at a position angle of 15\\,$\\degr$. The first Hubble Space Telescope (HST) Wide Field Planetary Camera 2 (WFPC2) images of XZ~Tauri taken in 1995 show a bubble of emission nebulosity extending 4$\\arcsec$ to the north of the system (\\cite{Krist97}, hereafter K97).  Further images, taken 3 years later, show dramatic structural changes as the bubble expanded and altered from being centre-brightened to limb-brightened, suggesting the formation of a HH bowshock (\\cite{Krist99}, hereafter K99).  Ground-based photometry of XZ~Tau from 1962 to 1981 (\\cite{Herbst94}) has shown variations, of almost two magnitudes in the V band, for the binary as a whole. Such variations are common amongst young stellar objects (YSOs). Of the two components, the southern one has been, at least until recently, optically brighter and thought to be the more evolved star. Its companion, however, dominates at infra-red wavelengths and is probably of higher overall luminosity (\\cite{Haas90}). Recent Faint Object Spectrograph (FOS) observations, however, unexpectedly found the northern component to be optically brighter (\\cite{White01}, hereafter WG01), a result that we will discuss further in the light of our findings. For this system therefore it seems more appropriate not to use the terms primary and secondary but instead we will adopt the nomenclature, used in K97, of XZ Tau North and South.  We report here analysis of further HST Archive WFPC2 images of the XZ~Tau system and outflow from 1999, 2000 and 2001. These data show not only ongoing changes in the outflow but a dramatic brightening of XZ~Tau North in the optical suggesting that it may be an EXor. We also simulate the outflow in an attempt to reproduce its dynamical and morphological evolution. ",
        "conclusions": "Multi-epoch HST/WFPC observations of the XZ~Tau binary and its associated outflow have shown considerable changes in the system within only 6 years, from 1995 to 2001. The presence of {\\em two} limb-brightened shock fronts is now clearly evident, with a deceleration of the outer shock from 146\\,km\\,s$^{-1}$ to 117\\,km\\,s$^{-1}$.  Stellar photometry revealed that the suspected source of the outflow, XZ~Tau North, has flared in EXor-type fashion increasing in brightness by 3 magnitudes in R between 1998 and 2001. Finally, numerical simulations of the outflow produced reasonable agreement with observation in terms of morphology, dynamical evolution and emission line luminosity, using plausible ambient and outflow parameters. Deceleration by the amount observed, caused by the ambient medium, should have produced a much brighter bowshock apex than that seen. The cause of this discrepancy is not obvious.   \\vspace {0.3in} {\\bf Acknowledgements} \\newline  We wish to thank the referee, Dr S. Cabrit, for useful comments and suggestions. D.C. and T.P.R. would like to acknowledge support for their research from Enterprise Ireland. This work was carried out as part of the CosmoGrid project, funded under the Programme for Research in Third Level Institutions (PRTLI) administered by the Irish Higher Education Authority under the National Development Plan and with partial support from the European Regional Development Fund."
    },
    "0402/astro-ph0402403_arXiv.txt": {
        "abstract": "The main goal of this research is to get better insights into the properties of the plasma filled magnetospheres of black holes by means of direct numerical simulations and, ultimately, to resolve the controversy surrounding the Blandford-Znajek mechanism. Driven by the need to write the equations of black hole electrodynamics in the form convenient for numerical applications, we constructed a new system of 3+1 equations which not only has more traditional form than now classical 3+1 system of Thorne and Macdonald but which is also more general. To deal with the magnetospheric current sheets, we also developed a simple model of radiative resistivity based on the inverse Compton scattering of background photons. The results of numerical simulations combined with simple analytical arguments allow us to make a number of important conclusions on the nature of the Blandford-Znajek mechanism. We show that, just like in the Penrose mechanism and in the MHD models of Punsly and Coroniti, the key role in this mechanism is played by the black hole ergosphere. The poloidal currents are driven by the gravitationally induced electric field which cannot be screened within the ergosphere by any static distribution of the electric charge of locally created pair plasma. Contrary to what is expected in the Membrane paradigm, the energy and angular momentum are extracted not only along the magnetic field lines penetrating the event horizon but along all field lines penetrating the ergosphere. In dipolar magnetic configurations symmetric relative to the equatorial plane the force-free approximation breaks down within the ergosphere where a strong current sheet develops along the equatorial plane. This current sheet supplies energy and angular momentum at infinity to the surrounding force-free magnetosphere.  The Blandford-Znajek monopole solution is found to be asymptotically stable and causal. The so-called horizon boundary condition of Znajek is shown to be a regularity condition at fast critical surface. ",
        "introduction": "As the result of great advances in observational astrophysics during the previous several decades black holes are no longer regarded as peculiar solutions allowed by the Einstein equations which may or may not have anything to do with real astronomical phenomena. It is now widely believed that black holes are common in the Universe and play a key role in the most violent space events such as activity of galactic nuclei and gamma-ray bursts. Enormous amounts of energy released during such events can have two different origins. First of all this can be the gravitational energy of matter released during accretion onto an already existing black hole or during the gravitational collapse leading to formation of a new black hole. On the other hand, this can be the rotational energy of a black hole itself.  The notion of rotational energy of a black hole emerged as the result of theoretical discovery made by Penrose(1969) who found that a fraction of the Kerr hole mass can be converted, at least in principle, into the energy of surrounding matter or radiation (see also Cristodolou 1970). He has shown that two particles may interact in the neighbourhood of the black hole in such a way that one of the particles acquires negative energy and eventually disappears into the hole whereas the other particle carries away excessive positive energy. Unfortunately, as this was shown later, such interactions are rather special and must be very rare under typical astrophysical conditions, rendering the Penrose process inefficient (Bardeen et al. 1973). However, it was also found that the electromagnetic field can be used to extract the rotational energy of black holes too.  First, Goldreich and Julian(1969) analysed the vacuum solution for a rotating neutron star with a dipolar magnetic field aligned with the rotational axis. They argued that the rotationally induced electric field was strong enough to pull charged particles from the stellar surface and, thus, fill the surrounding space with plasma. Using the force-free approximation to describe the produced magnetosphere they argued that an electromagnetically driven wind would carry away rotational energy and angular momentum of the star.  Then, Wald(1974) found a rather interesting particular solution of the vacuum Maxwell equations in the Kerr spacetime.  Far away from the hole this solution described a uniform magnetic field aligned with the rotational axis of a black hole. However, near the black hole it described a strong electric field as well. Moreover, just like in the problem of an aligned rotator this ``gravitationally induced'' electric field had a significant component along the magnetic field. Bisnovatyi-Kogan and Ruzmaikin (1976) argued  that, in the case of astrophysical black holes, a rather strong magnetic field of such kind can be generated in their accretion discs.  Finally, Blandford \\& Znajek (1977) realised that the similarity between the vacuum solution for a Kerr black hole and the vacuum solution for a rotating neutron star meant the possibility of electromagnetically driven wind from a rotating black hole, provided the space around the black hole could be filled with plasma. Moreover, they argued that, under the typical astrophysical conditions, the vacuum solutions were, in fact, unstable to cascade pair production, ensuring a plentiful supply of charged particles. Then they developed a general theory of force-free steady-state axisymmetric magnetospheres of black holes and found a perturbative solution for a slowly rotating black hole with monopole magnetic field. The key element of this solution was Znajek's ``boundary condition'' \\cite{Z77} imposed on the event horizon. As expected, this solution exhibited outgoing electromagnetic fluxes of energy and angular momentum. Moreover, the electromagnetic mechanism seemed to be very robust and the estimated power of the wind was high enough to explain the energetics of radio galaxies and quasars.  In their analysis, Blandford \\& Znajek (1977) used covariant equations of electrodynamics and operated with components of the electromagnetic field tensor and four-potential.  Later, Thorne and Macdonald \\cite{TM,MT} developed a 3+1 approach, where the equations of black hole electrodynamics were written in more or less traditional form in terms of the spatial vectors of electric and magnetic field as measured by the so-called local fiducial observer (FIDO). In this work they adopted the system of coordinates due to Boyer and Lindquist(1967) which has a number of useful properties but also one important drawback. Just like the system of Schwarzschild coordinates, it is singular at the event horizon. As the result, the event horizon appears as a peculiar inner boundary of of physical space which required rather special treatment.  A number of authors studied the properties of electromagnetic field near the event horizon (Hanni \\& Ruffini 1973; Hajicek 1974; Znajek 1977,1978; Damour 1978) and gradually a picture emerged, according to which the event horizon could be treated as a rotating conducting surface with surface charges, surface currents, and a finite surface resistivity. This perfectly suited the quest of Thorne and Mackdonald for a new formulation of the black hole electrodynamics which would make it look similar to the classical electrodynamics.  Surprisingly enough, the drawback of the Boyer-Lindquist coordinates seemed to turn into an advantage. This theory of the event horizon have made a great impact on the current perception of the Blandford-Znajek mechanism which is now widely associated with a mental picture of magnetic field lines originating from the horizon and being torqued by its rotation. It stimulated the development of The Membrane Paradigm \\cite{TPM} where the event horizon is attributed with a whole range of physical properties. However, one has to admit that, in spite of all its attractive simplicity and mathematical correctness, this construction is purely artificial. Moreover, it is rather worrying that the emphasis put on the role of the event horizon makes the Blandford-Znajek mechanism completely alien to the Penrose mechanism, where the key role is played by the black hole ergosphere.  The electrodynamic mechanism together with the horizon theory is now widely accepted by the astrophysical community.  In great contrast to this mainstream trend, Punsly and Coroniti \\cite{PC,PC1} and later Punsly (see the review in Punsly,2001) completely rejected both these theories. They argued that the event horizon cannot be regarded as a unipolar inductor because it is causally disconnected from the outgoing wind.  Indeed, both the fast and the Alfv\\'en waves generated at the event horizon can propagate only inwards and cannot effect the events in the outer space. The apparent lack of a proper unipolar inductor in the Blandford-Znajek solution and its reliance on Znajek's boundary condition made Punsly and Coroniti to conclude that this solution is nonphysical and structurally unstable.  They developed completely different MHD models which seemed to be based on clearer physical ideas. In brief, they argue that gravity forces magnetospheric plasma to rotate inside the black hole ergosphere in the same sense as the black hole and that the magnetic field exhibits a similar rotation because it is \"frozen\" into this plasma.  Koide (2003) carried out MHD simulations of dipolar black hole magnetospheres with plasma energy-density only $1\\div2$ order of magnitude smaller than the energy density of electromagnetic field. Although the numerical solution did not reached a steady state in these simulations, the obtained results seemed to indicate that in this regime the inertial effects accounted for approximately half of the extracted energy. However, in the black hole magnetospheres filled with plasma via pair production the characteristic number density of particles is given by the so-called Goldreich-Julian density \\cite{GJ,Bes,Hir-Oka}.  Under the typical conditions of supermassive black holes in AGNs, this corresponds to the mass density of pair plasma hardly exceeding $10^{-13}$ of the energy density of the electromagnetic field. Thus, even if the particle density is several orders of magnitude larger than the Goldreich-Julian one, the inertial effects are still expected to be negligibly small. Given the fact that, in such a degenerate regime, the full system of relativistic MHD presents a real challenge for numerical methods (see Komissarov, 2001a, Gammie et al., 2003), it is difficult not to conclude that the electrodynamic approach is more suitable. This is one  of the reasons for the renewed interest towards the theory of force-free electrodynamics \\cite{Kom02a,Bland02}. (This does not mean that the pair production is the only possible way of plasma supply in the neighbourhood of astrophysical black holes and that the inertial effects are always that small. For example, the coronas of black hole accretion discs are likely to be filled with dense plasma pulled from the disc surface.)  The recent numerical studies of the force-free magnetospheres of black holes added to the controversy surrounding the Blandford-Znajek mechanism. On one hand, they revealed the asymptotic stability of the Blandford-Znajek monopole solution \\cite{Kom01} and, thus, raised doubts about the validity of the causality arguments by Punsly and Coroniti.  On the other hand, the results for the dipolar magnetospheres questioned the virtues of the horizon theory as well. Indeen, they clearly indicated that the key role in the electrodynamic mechanism is played not by the black hole event horizon but by its ergosphere \\cite{Kom02b}. Moreover, these simulations also revealed a certain deficiency of the force-free approximation as a dissipative current sheet was formed in the equatorial plane within the ergosphere.   This paper is an attempt to understand the nature of the electrodynamic mechanism and to resolve the controversy surrounding it.  In Sec.II we show that the 3+1 equations of the black hole electrodynamics can be written in a more general and somewhat simpler form than in \\cite{MT,TM}. In Sec.III the basic results for the force-free magnetospheres are re-derived using the new 3+1 formulation. In order to handle the ergospheric current sheets, one has to go beyond the force-free approximation and consider the resistive electrodynamics. A model of resistivity based on the inverse Compton scattering of the background photons is described in Sec.IV.  In Sec.V we present the results of numerical studies of the black hole magnetospheres obtained within the framework of resistive electrodynamics.  The implications of these results for the perception of the Blandford-Znajek mechanism are discussed in Sec.VI. The more technical information about the properties of light surfaces, the properties of the Kerr-Schild coordinate system, and the details of our numerical method is presented in the Appendix.  Throughout this paper we adopt $(-+++)$ signature for the spacetime and assume that the Greek indices range from 0 to 3 and refer to spacetime tensors whereas the Latin ones range from 1 to 3 and refer to purely spatial tensors. In addition, we adopt such units that the speed of light and $4\\pi$ do not appear in the equations of electrodynamics. ",
        "conclusions": "\\begin{itemize}  \\item The 3+1 equations of black hole electrodynamics can be written in more general and simpler form than the classical formulation of Thorne and Macdonald \\cite{TM,MT,TPM}. Our equations are very similar to the classical equations of electrodynamics in matter and hold in any system of coordinates which delivers time-independent metric form, e.g. Kerr-Schild coordinates.  \\item The inertia of magnetospheric plasma is not an essential element in the mechanisms of extraction of rotational energy of black holes involving electromagnetic field. We have seen no indications of any major deficiency of the electrodynamic model which would suggest a need for the more complex MHD model, at least at the black hole end. The force-free approximation, however, may break down locally via formation of dissipative current sheets. We have constructed a simple model for radiative resistivity based on the inverse Compton scattering of background photons and used this model to show the development of the equatorial current sheet within the ergosphere of a black hole with a dipolar magnetic field. This current sheet supplies energy and angular momentum for the surrounding force-free magnetosphere.  Having said that, we need to stress that the role of particle inertial has not been considered in this paper and requires further investigation. For example, one would expect the inertial effects to become important in the remote part of the outgoing wind, like they seem to do in the winds from neutron stars, e.g. \\cite{Mestel,KenCor}.  \\item  The gravitationally induced electric field is the ultimate cause of the poloidal currents in the black hole magnetospheres. Within the ergosphere this field cannot be screened by any static distribution of electric charge. This makes a ``magnetized'' black hole very different from the classical unipolar inductor (or any other battery), where the potential difference between terminals is created and sustained by the electromotive force which first separates electric charges and then drives electric currents against the electric field. In fact, within the ergospheric current sheet the cross-field current flows along the electric field but not against it. The special role played by the ergosphere allows us to call it  the ``driving force'' of the Blandford-Znajek mechanism. Since the same applies to the otherwise rather different Penrose mechanism \\cite{Penrose}, this suggests that it is the existence of the ergosphere that makes possible energy extraction from a black hole in any form.  \\item The Blandford-Znajek solution in particular, and the electromagnetic mechanism of extraction of rotational energy of black holes in general, do not clash with causality. First of all, the Znajek horizon condition is not a boundary condition but a regularity condition at the fast critical point. This perfectly justifies its utilization in the Blandford-Znajek monopole solution. Secondly, the ergosphere is causally connected to the outgoing wind by means of both fast and Alfv\\'en waves. This conclusion is strongly supported by numerical simulations that show that the Blandford-Znajek monopole solution is asymptotically stable.  \\item Our results fully agree with the conclusion of Punsly and Coroniti \\cite{PC,PC1} and Beskin \\& Kuznetsova \\shortcite{Bes-Kus} on the superficial nature of the interpretation of the event horizon as a unipolar inductor of black holes. The failure of the Membrane paradigm to predict the outflow of energy and angular momentum along all magnetic field lines penetrating the ergosphere has clearly exposed its limitations. Although the analogy with conducting sphere is based on mathematically sound grounds, it does not account for all important properties of black hole electrodynamics and does not reveal the true physical nature of the Blandford-Znajek mechanism.  \\end{itemize}"
    },
    "0402/astro-ph0402259_arXiv.txt": {
        "abstract": "{ We have built two samples of galaxies selected at 0.2$\\mu$m (hereafter UV) and 60 $\\mu$m (hereafter FIR) covering a sky area of 35.36 deg$^2$. The UV selected sample contains 25 galaxies brighter than $AB_{0.2}=17$. All of them, but one elliptical, are detected at 60 $\\mu$m with a flux density larger or equal to  0.2~Jy. The UV counts are significantly lower than the euclidean extrapolation towards brighter fluxes of previous determinations. The FIR selected sample contains 42 galaxies brighter than $f_{60}$=0.6 Jy. Excepting four galaxies, all of them have a UV counterpart at the limiting magnitude $AB_{0.2}=20.3$~mag. The mean extinction derived from the analysis of the FIR to UV flux ratio is $\\sim 1$~mag for the UV selected sample and $\\sim 2$~mag for the FIR selected one. For each sample we compare several indicators of the recent star formation rate (SFR) based on the FIR and/or the UV emissions and we find linear relationships with slopes close to unity, meaning that no trend with the SFR exists when converting between each other. Various absolute calibrations for both samples are discussed in this paper. A positive correlation between extinction and SFR is found when both samples are considered together although with a considerable scatter. A similar result is obtained when using the SFR normalized to the optical surface of the galaxies. ",
        "introduction": "Tracing the star formation activity in galaxies at all redshifts is a fundamental step towards the understanding of the formation and evolution of the Universe. The star formation activity is commonly quantified by the Star Formation Rate (SFR) defined as the stellar mass formed per unit of time. In order to efficiently constrain the models of galaxy  formation and evolution it is important to measure a  SFR as current as possible in order not to integrate on too large lookback times. Quite naturally, the light emitted by young stars can be used to measure this SFR. The most commonly used tracers of the SFR in galaxies, at least in the nearby Universe, are the UV, FIR and H$\\alpha$ emissions (e.g. Kennicutt 1998). In the present study we focus our attention on the UV and FIR emissions.  Concerning the FIR window, the IRAS mission provided us with several now well studied samples such as the Point Source Catalog (PSC, Joint IRAS Science 1994), the Faint Source Catalog (FSC, Moshir et al. 1990) and the redshift survey of the Point Source Catalog (PSCz, Saunders et al. 2000). The UV window, less explored mostly because limited observations have been available until now (e.g. Donas et al. 1987; Deharveng et al. 1994; Milliard et al. 1992; Kinney et al. 1993; Bell \\& Kennicutt 2001), has been extensively studied (e.g. Milliard et al. 1992; Donas et al. 1995; Treyer et al. 1998). This situation should change dramatically in the near future once the GALEX data will be available.  The first step towards the determination of statistical properties of UV and FIR selected samples of galaxies is the determination of the galaxy counts and the luminosity function (LF). Accurate determinations of these observables are important for constraining models of galaxy formation and evolution. Nevertheless very few models predict the luminosity distribution and spectral energy distribution of galaxies over a large range of wavelength from UV to FIR (Totani \\& Takeuchi 2002 and references therein; Xu et al. 1998). This is due, at least in part, to our poor knowledge of the dust extinction in the universe (amount of dust, dust emission, mechanism of stellar absorption on large scales). The reconstruction of the whole SED of galaxies from UV to FIR is difficult even in the well sampled nearby Universe because of the lack of multiwavelength data on large and homogeneous samples of galaxies (Boselli et al. 2003; Flores et al. 1999; Rigopoulou et al. 2000; Cardiel et al. 2003). This situation should evolve  dramatically in the future with the planned observations of large fields with telescopes working at very different wavelengths (SIRTF, GALEX, VIMOS, ASTRO-F).  The FIR to UV flux ratio has been proved to be the best indicator of the dust extinction in normal galaxies (e.g. Buat \\& Xu 1996; Meurer et al. 1999), becoming a fundamental parameter in the determination of their present SFR (e.g. Buat et al. 1999,2002; Hirashita et al. 2003, hereafter HBI).  The aim of the present paper is to study the statistical properties of two samples extracted from the same area of the sky, selected according to UV and FIR criteria. The two selections (UV and FIR) appear very complementary since the sample of UV selected galaxies will be biased towards active star forming galaxies with low extinction, and on  the contrary a FIR selection is likely to favor galaxies with high extinction. An accurate knowledge of the statistical properties of UV and FIR selected samples is crucial for the analysis of similarly selected samples at higher $z$, where the lack of complementary data prevents a precise determinations of the dust extinction and SFR.  With these two samples in hand, we study the extinction and SFR related properties and their dependence on the selection method. More precisely we address the following issues: (1) determine the mean extinction of purely UV or FIR selected samples of galaxies, (2) observationally constrain models of galaxy formation and evolution, and (3) provide the best recipes for the determination of the SFR, valid for UV and/or FIR selected samples of galaxies.  The paper is organized as follows: Section~2 presents the FIR and UV data and the region of the sky where the data were taken. Sections~3 and 4 describe the selection procedure of the UV and FIR selected samples. Section~5 discusses  the extinction properties of the samples and Section~6 is devoted to the comparison between the different SFR tracers. A final summary of the main results of the paper is presented in Section~7.  \\section {The FIR and UV data}  An accurate comparison between the properties of UV and FIR galaxies can be done if both samples are extracted from the same region of the sky. With this purpose we selected some fields covered by the FOCA experiment (Milliard et al. 1991) at UV wavelengths (0.2$\\mu$m). A total of 9 FOCA fields were chosen. The area of each field is circular with a radius of 1.13~$\\deg$. Accounting for some overlapping of the fields, a total area of 35.36~$\\deg^{2}$ was covered. Table~\\ref{tabla} shows the basic properties of the selected fields. Four out of the nine observed fields are centered on nearby clusters of galaxies: m010 (Cancer), m028 (Coma), m067 (Abell~1367) and m050 (Virgo), and two of them point to the external parts of the Coma cluster (m030, m031). The Cancer cluster is however in Hubble flow (Gavazzi et al. 1991), thus its members are similar to field galaxies. The fractions of cluster galaxies in our UV selected and FIR selected (hereafter UVsel and FIRsel respectively) samples are 64\\% and 50\\% respectively. The possible contribution of the cluster environment to the properties of the UVsel and FIRsel galaxies are discussed in the text.  The detection and flux extraction of the UV objects in the FOCA plates was carried out in an automatic way. Only UV sources with surface brightness (averaged over $15 \\times 15$~arcsec$^{2}$) brighter than 2.8 times the sky $\\sigma$ (the detection limit established by the automatic detection algorithm for each frame, given in Col.~5 of Table~1) were considered as detections. The automatic determination of the UV flux is accurate for point like sources but uncertain for extended objects, so the UV fluxes of the FIRsel galaxies were determined in two ways: (1) for the galaxies which were not resolved by the automatic detection algorithm (i.e. only one UV source was detected) we used the UV flux provided by the detection algorithm\\footnote{The automatic algorithm performs aperture photometry around the selected UV sources using increasing radii. The adopted aperture is selected when the magnitude stabilizes. More details about this procedure are given in Moulinec (1989).}. (2) for the galaxies which were splitted in several sources (i.e. several H{\\sc ii} regions were detected as individual sources) we performed aperture photometry by hand. The optical images were used to set the size of the apertures in order to be sure that all the UV flux corresponding to the galaxies, even the most external star forming regions, was included in the total UV flux. In both cases the apertures used to measure the UV fluxes are close to the optical diameters of the galaxies. The average zero point uncertainty of the FOCA data is about 0.2~mag. The spatial resolution of the FOCA frames is $\\approx 20$~arcsecs FWHM and the $1\\sigma$ astrometric uncertainty is $\\approx 2.5$~arcsecs.  Concerning the FIR data, the IRAS all sky survey ensures a total coverage of our selected FOCA fields. The total photometric uncertainty of the FIR data of the IRAS mission is $\\approx 10$\\% at 60 and 100$\\mu$m. The positional uncertainty is variable and of elliptical shape with major axis $\\approx 40$~arcsecs. The spatial resolution shows irregular shape and depends on the bandwidth and coordinates of the object. In the direction where it is maximal being $\\approx 4.5$~arcmin at 60$\\mu$m (see Moshir et al. 1990 for details). ",
        "conclusions": "We have presented the properties of two samples of galaxies selected from the same region of the sky by their UV and FIR fluxes. The detection rate for the UVsel sample at FIR wavelengths was found to be 96\\% whereas that for the FIRsel sample at UV wavelengths equals 90\\%. The UV counts are lower than the expected from the extrapolation of previous determinations at fainter magnitudes, even when including the contribution of the cluster galaxies. We showed that their dust extinction properties are different, the UV selected galaxies exhibiting a lower extinction than the FIR selected ones ($\\approx$1~mag on average for the UVsel vs. $\\approx$2~mag for the FIRsel). Four galaxies of the FIRsel sample do not have a UV counterpart, implying lower limits to the UV extinction of $3.8$~mag. These galaxies could be part of the population of very extincted objects already reported in the literature.  We compared  different indicators of the SFR calculated with the FIR and/or UV  luminosities and we showed that they correlate well with each other for both samples. The relations between the different estimators of the SFR present a slope close to unity for both samples, meaning that no trend with the SFR exists when converting between each other. For both samples we found the best agreement between the following quantities: (a) the SFR calculated from the UV luminosities corrected for dust extinction using the FIR/UV ratio and (b) the sum of the SFR calculated from the dust luminosities corrected for the average contribution of the dust heating due to old stars ($\\approx 40$\\%) and of the SFR calculated from the observed UV luminosities.  Putting both samples together we find the correlation  between SFR and extinction already reported for other samples of galaxies but with a very large scatter. Most of the trend is due to the galaxies selected in FIR.  The results of this work seem not affected by the cluster environment since we have shown that the global properties of cluster and field galaxies present in our samples are similar. This is expected from the similarities between field and cluster LFs at UV and FIR wavelengths (Bicay \\& Giovanelli 1987; Cortese et al. 2003)."
    },
    "0402/astro-ph0402545_arXiv.txt": {
        "abstract": "Distributions of Faraday rotation measure (FRM) and the projected magnetic field derived by a 3-dimensional simulation of MHD jets are investigated based on our \"sweeping magnetic twist model\". FRM and Stokes parameters were calculated to be compared with radio observations of large scale wiggled AGN jets on kpc scales. We propose that the FRM distribution can be used to discuss the 3-dimensional structure of magnetic field around jets and the validity of existing theoretical models, together with the projected magnetic field derived from Stokes parameters. In the previous paper, we investigated the basic straight part of AGN jets by using the result of a 2-dimensional axisymmetric simulation. The derived FRM distribution has a general tendency to have a gradient across the jet axis, which is due to the toroidal component of the magnetic field generated by the rotation of the accretion disk. In this paper, we consider the wiggled structure of the AGN jets by using the result of a 3-dimensional simulation. Our numerical results show that the distributions of FRM and the projected magnetic field have a clear correlation with the large scale structure of the jet itself, namely, 3-dimensional helix. Distributions, seeing the jet from a certain direction, show a good matching with those in a part of 3C449 jet. This suggests that the jet has a helical structure and that the magnetic field (especially the toroidal component) plays an important role in the dynamics of the wiggle formation because it is due to a current-driven helical kink instability in our model. ",
        "introduction": "To explain the formation of active galactic nucleus (AGN) jets and other astrophysical jets, various models have been proposed. Among them, the magnetohydorodynamic (MHD) model is one of the most promising models, since it can explain both the acceleration and the collimation of the jets (see, e.g., Meier, Koide, \\& Uchida 2001, and references therein). Lovelace (1976) and Blandford (1976) first proposed the theoretical model of the magnetically driven jet from accretion disks, and Blandford \\& Payne (1982) discussed magneto-centrifugally driven outflow from a Keplerian disk in steady, axisymmetric and self-similar situation. Time-dependent, 2-dimensional axisymmetric simulations were performed by Uchida \\& Shibata (1985), Shibata \\& Uchida (1986), Uchida \\& Shibata (1986). They pointed out that large amplitude torsional Alfv\\'en waves (TAWs) generated by the interaction between the accretion disk and a large scale magnetic field play an important role. The toroidal magnetic field propagates along the large scale magnetic field while squeezing it into a collimated jet-shape by the pinching effect of the Lorentz force. In this paper, we refer this model as a \"sweeping magnetic twist model\". After these papers, many authors have performed time-dependent, 2-dimensional axisymmetric simulations (e.g. Stone \\& Norman 1994, Ustyugova et al. 1995, Matsumoto et al. 1996, Ouyed \\& Pudritz 1997, Kudoh, Matsumoto, \\& Shibata 1998).  Using the numerical data of MHD models, observational quantities such as Faraday rotation measure (FRM) or Stokes parameters have been derived to be compared with observations of AGN jets: Laing (1981) computed the total intensity, the linear polarization, and the projected magnetic field distributions, assuming some simple magnetic field configurations and high energy particle distributions in the cylindrical jet. Clarke, Norman, \\& Burns (1989) performed 2-dimensional MHD simulations in which a supersonic jet with a dynamically passive helical magnetic field was computed, and derived distributions of the total intensity, the projected electric field, and the linear polarization. Hardee \\& Rosen (1999) calculated the total intensity and the projected magnetic field distributions, using 3-dimensional MHD simulations of strongly magnetized conical jets. Hardee \\& Rosen (2002) calculated the FRM distribution and discussed that the observation of the radio source 3C465 in Abell cluster A2634 (Eilek \\& Owen 2002) suggests helical twisting of the flow.  FRM is given by the integral of $n_e B_\\parallel$ along the line-of-sight between the emitter and the observer (where $B_\\parallel$ is the line-of-sight component of the magnetic field, and $n_e$ is the electron density there). It is, in principle, not possible to specify which part on the line-of-sight the contribution comes from. However, in recent high-resolution radio observations (e.g. Eilek \\& Owen 2002, Asada et al. 2002), the FRM distribution seems to have good correlation with the configuration of the jet; this suggests that the FRM variation is due to the magnetized thermal plasma surrounding the emitting part of the jet. In fact, sharp FRM gradients seen in 3C273 can not be produced by a foreground Faraday screen (Taylor 1998, Asada et al. 2002). If this is the case, we can get a new information, that is, the line-of-sight component of the magnetic field, and thus can predict the 3-dimensional configuration of the magnetic field around the jet, together with the projected magnetic field.  Uchida et al. (2004) (hereafter paper I) carried out a 2-dimensional axisymmetric simulation, and investigated the model counterpart distributions of FRM and the projected magnetic field in the basic straight part of AGN jets. It was described how a systematically helical field configuration is produced in the \"sweeping magnetic twist model\". It was also shown as a result that the model can reproduce fairly well the characteristic distribution of FRM having the gradient across the jet axis (Perley, Bridle, \\& Willis 1984, Asada et al. 2002, Gabuzda \\& Murray 2003). The systematic FRM gradient was caused by the gradient of the line-of-sight component of the magnetic field. This means that the existence of the helical magnetic field (and the propagation of TAWs) is plausible. On the basis of this success, in this paper the treatment is advanced to the non-axisymmetric situation, the wiggled structure, in AGN jets.  The morphological structures in AGN jets, such as \"wiggles (kinks)\" or \"bends\" are frequently seen not only on kpc scales, but also on pc scales and smaller (Hummel et al. 1992). Such a helical distortion might be caused either by plasma instabilities or precession of jet ejection axis due to the gravitational interaction between binary Black Holes (BBHs) (Begelman, Blandford, \\& Rees 1980), BH/disk, or galaxies. Magnetically driven jets possess a toroidal field component, which is equivalent to an axial electric current, and such \"current-carrying\" jets are susceptible to MHD instabilities, and moreover, Kelvin-Helmholtz instabilities must be also taken into account for jet dynamics. MHD instabilities are usually divided into pressure-driven instabilities and current-driven instabilities (Bateman 1980). Kelvin-Helmholtz instabilities have been considered for past decades by many theoretical and numerical work (for reviews, see Birkinshaw 1991; Ferrari 1998, and references therein). One of the promising possibilities is a magnetic mechanism based on the \"sweeping magnetic twist model\". It was shown by 3-dimensional MHD simulations that the formation of the wiggled structure can be explained by the current-driven helical kink instability (Todo et al. 1993). In these simulations, the magnetic field was a force-free helical field from the beginning and the propagation of TAWs was not dealt with. Nakamura, Uchida, \\& Hirose (2001) extended the treatment of the \"sweeping magnetic twist model\" to the part far from the gravitator and the accretion disk. They investigated the behavior of TAWs propagating far from the AGN core. They showed that the current-driven helical kink instability can explain the production of the observed wiggles. If the mechanism of the wiggle formation is clarified, it can become a clue to the physics of the jet formation.  In this paper, we calculate FRM, the projected magnetic field, and the total intensity from the numerical data of a MHD simulation based on our ``sweeping magnetic twist model'', and discuss these model counterparts comparing with an observation. Here we consider the wiggled structure of the jet, and thus use the same kind of 3-dimensional simulation as in Nakamura et al. (2001). In section 2, we describe the application of the model to the distant part of the jet from the AGN core, and show the formation of a helical structure of the jet by the current-driven helical kink instability. We introduce the method to calculate model counterparts of observational quantities in section 3, and show the results in section 4. We discuss comparison between the calculated distributions and observed ones in section 5. The conclusion is summarized in section 6. ",
        "conclusions": "We performed a 3-dimensional MHD simulation based on our ``sweeping magnetic twist model'', which was applied to the situation far from the AGN core by Nakamura et al. (2001). In the ``sweeping magnetic twist model'', the disk rotation generates the toroidal magnetic field and it propagates into two directions along a large scale magnetic field as torsional Alfv\\'en waves (TAWs). In this paper, the situation TAWs are propagating far from the gravitator was dealt with. It was assumed that there is a lower Alfv\\'en velocity region ahead of the propagating TAWs. The toroidal magnetic field becomes accumulated after TAWs enter the low Alfv\\'en velocity region. This causes the growth of the current-driven helical kink instability and the wiggled structure is formed.  We also calculated the observational quantities (FRM, Stokes, $I$, $Q$, and $U$ parameters) by integrating the numerical data along the line-of-sight. The Faraday rotation screen and the emitting region were defined separately. The integration for FRM was done only in the Faraday rotation screen and that for Stokes parameters only in the emitting region. The projected magnetic field was determined from $Q$ and $U$.  Before the formation of the wiggled structure, the FRM distribution has a gradient across the jet axis. An asymmetry in the total intensity distribution exists except the case where $\\theta$ is equal to $90^\\circ$. Same features are seen in the result of paper I and these features can explain some observations (Perley et al. 1984, Asada et al. 2002, Gabuzda \\& Murray 2003). The projected magnetic field is parallel to the jet axis because the pitch angle of the helical magnetic field is not so large.  After the formation of the wiggled structure, the FRM distribution has a clear correlation with the large scale structure of the jet itself. This is caused by the following; the magnetic field becomes less twisted around the wiggled structure owing to the current-driven helical kink instability. This is equivalent to that the magnetic field is almost along the structural axis of the jet in the wiggled structure. The total and polarization intensities are low in the wiggled structure when the jet inclination to the line-of-sight is large. However, they become higher as the jet inclination becomes small because the structure on the backside and that on the frontside of the 3-dimensional helix overlap on the same ray as seen by observer. We found that when we see the jet at a certain angle, we can reproduce the characteristics of the observation of the 3C449 jet (Feretti et al. 1999). This suggests that the FRM distribution could be strongly affected by the magnetic field in or around the jet and that the jet has a helical structure. If this is the case, it is also suggested that the magnetic field (especially the toroidal component) plays an important role in the formation of astrophysical jets."
    },
    "0402/astro-ph0402223_arXiv.txt": {
        "abstract": "{Evolution of a coronal loop in response to an impulsive energy release is numerically modelled. It is shown that the loop density evolution curves exhibit quasi-periodic perturbations with the periods given approximately by the ratio of the loop length to the average sound speed, associated with the second standing harmonics of an acoustic wave. The density perturbations have a maximum near the loop apex. The corresponding field-aligned flows have a node near the apex. We suggest that the quasi-periodic pulsations with periods in the range 10--300~s, frequently observed in flaring coronal loops in the radio, visible light and X-ray bands, may be produced by the second standing harmonic of the acoustic mode. ",
        "introduction": "Wave activity of the solar corona attracts attention in relation with coronal heating and solar wind acceleration problems, and as an efficient tool for MHD coronal seismology (e.g. Nakariakov 2003). The observational evidence of coronal waves and oscillations is abundant. Low period coronal oscillations, in the range between a few seconds to several minutes,  are believed to be associated with magnetohydrodynamic waves. In particular, propagating slow magnetoacoustic waves have been identified in polar plumes (Ofman, Nakariakov \\& DeForest 1999) and over loop footpoints (Nakariakov et al. 2000, Tsiklauri \\& Nakariakov 2001), and standing global slow modes in loops (Ofman \\& Wang 2002). The majority of confidently interpreted examples of the coronal wave activity has been found in the EUV coronal emission.  Observations in other spectral windows, in particular in the radio band, also demonstrate various kinds of oscillations (e.g., the quasi-periodic pulsations, or QPP, see Aschwanden 1987 for a review), usually with periods from a few seconds to several tens of seconds. It is commonly accepted that the waves with periods about several seconds may be produced by either propagating or standing sausage modes (Roberts, Edwin \\& Benz 1984; Nakariakov, Melnikov \\& Reznikova 2003, Aschwanden, Nakariakov \\& Melnikov 2004). However, it is known that the existence of the non-leaky sausage mode cut-off imposes an upper limit on the possible wave periods: waves with the periods longer than the cut-off period cannot be trapped in the loops. For a magnetic cylinder of the radius $a$, the cut-off period $P_\\mathrm{c}$ is estimated (see, e.g. Roberts et al. 1984) as \\begin{equation}\\label{p_c} P_\\mathrm{c} \\approx \\frac{2.62 a}{C_\\mathrm{Ae}} \\sqrt{\\frac{C_\\mathrm{Ae}^2-C_\\mathrm{A0}^2}{C_\\mathrm{s0}^2+ C_\\mathrm{A0}^2}}, \\end{equation} where $C_\\mathrm{s0}$ is the sound speed in the loop, and $C_\\mathrm{A0}$ and $C_\\mathrm{Ae}$ are the Alfv\\'en speeds inside and outside the slab, respectively. For example, for typical flaring loop parameters $a=6$~Mm, $C_\\mathrm{A0}=1$~Mm/s, $C_\\mathrm{Ae} = 3C_\\mathrm{A0}$ and $C_\\mathrm{s0}=0.5$~Mm/s, the cut-off period is about 15~s, and it is difficult to make it greater than 20~s. However, X-ray band observations often give much longer periodicities, frequently in association with a flare. For example, periodicities from 20~s to 25~min have been presented by Harrison (1987), McKenzie \\& Mullan (1997) and Terekhov et al. (2002). Similar periodicities have been observed in the decimeter and microwave bands. In particular, Wang \\& Xie (2000) observed QPP with the periods of about 50~s at 1.42 and 2~GHz (in association with an M4.4 X-ray flare). The coincidence of QPP periods observed in the X-ray and in radio bands is not a surprise, as the higher frequency radio bursts are found to correlate very well with X-ray bursts (e.g., Benz \\& Kane 1986). Moreover, pulsations with the periods significantly greater than the estimated sausage mode cut-off period have been found in both hard X-ray and microwave bands simultaneously (e.g. Fu, Liu \\& Li 1996, Tian, Gao \\& Fu 1999).  Similar oscillations, with the period of 220~s, have recently been observed in the white-light emission associated with stellar flaring loops (Mathioudakis et al. 2003). Possibly, the 160~s periodic oscillations observed by Houdebine et al. (1993) during a flare on Ad-Leonis and 26 and 13~s coherent oscillations observed by Zhilyaev et al. (2000) in EV Lac flares have the same nature.  A possible interpretation of these medium period QPPs may be connected with kink or torsional modes (Zaitsev \\& Stepanov 1989). However, these modes are practically incompressible and, in the case of a small amplitude, the produced perturbation of the magnetic field is also very weak. (E.g., the direct observations of kink modes in EUV, Aschwanden et al. 1999, Nakariakov et al. 1999). Thus, it is not simple to link these weak perturbations of the magnetic field with observed QPPs.  In this study, we suggest an alternative mechanism for the generation of long-period QPPs. We demonstrate that in a coronal loop an impulsive energy release  generates the second spatial harmonics of the acoustic mode. These oscillations are of high quality and do not experience dissipation. This mode, producing noticeable perturbations of the loop density and generating field-aligned flows, is shown to be responsible for QPP with medium and long periods. ",
        "conclusions": "We suggest that the second standing acoustic mode may be responsible for QPP with  periods of about 10--300~s (estimated by Eq.~(\\ref{ther})) observed in solar and stellar flare light curves. This mechanism is similar to the interpretation of coronal loop oscillations observed with SUMER, proposed by Ofman \\& Wang (2002). The main new element in our study is connected with the mode excitation. We demonstrate that the second standing acoustic harmonic appears as \\textit{a natural response} of the loop to an impulsive energy deposition. The SUMER oscillations are likely to be associated with some other excitation mechanism, as only a small fraction of SUMER oscillations are observed in association with solar flares (Wang et al. 2003).  Traditionally, the acoustic wave interpretation was excluded as these waves were supposed to be highly dissipative. However, numerical simulations discussed above, as well as recently gained abundant observational evidence of the presence of acoustic waves in the solar corona, suggest that the observed periodicities can be associated with this mode. The physical mechanism responsible for the induction of the quasi-periodic pulsations can be understood in terms of auto-oscillations generated by an electric-circuit generator. Indeed, the physical system modelled here contains all the necessary ingredients of a generator: the DC power supply (thermal instability), the nonlinear element (the plasma) and the resonator (the loop). This may explain why the oscillations may be observed to be dissipationless. However, proper analytical theory of the excitation of this mode is still to be developed.  The oscillation period is determined by the ratio of the loop length and the average sound speed in the loop. The typical amplitude of the density and velocity perturbations is 2-10\\% of the background. This value is consistent with observed amplitudes of X-ray QPPs (McKenzie \\& Mullan 1997). However, the observed amplitudes of radioband QPPs are sometimes higher. This discrepancy would be resolved by taking into account the specific mechanism responsible for the modulation of radio emission, and the line-of-sight angle (e.g. Cooper, Nakariakov \\& Williams 2003). In certain condition, the modulation mechanism can amplify the pulsations up to the observable level.  The model developed here is quite approximate, as it does not take into account two-dimensional MHD effects such as centrifugal force and the perturbation of the loop cross-section. However, we believe that the neglected effects do not change the qualitative picture described here and that Eq.~(\\ref{2sah}) gives a correct estimation for the oscillation period.  Finally, we would like to emphasize that the generation of the second spatial harmonic of an acoustic wave, by means of which we try to explain the observed quasi-periodic oscillations in flaring loops, is a consistent feature, seen for a wide range of physical parameters, including the case of non-symmetric heating functions. More detailed analysis will be presented elsewhere."
    },
    "0402/astro-ph0402015_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:1} Whether the formation of massive stars significantly differs from the formation of low mass stars is not yet fully understood. While disk accretion seems possible under special circumstances (e.g. \\cite{jia96}), coalescence of smaller objects into one massive star is an alternative model (e.g. \\cite{bon98}). Observations of molecular outflows in regions of massive star formation indicate, however, that the accretion model is indeed applicable. Even though little is known about the initial conditions of \\index{massive star formation} \\cite{eva02}, young massive protostars are expected not to be detectable at NIR/MIR wavelengths, but mainly in the submm continuum. \\\\ \\\\ Using SCUBA and IRAM bolometers, we investigated a sample of 47 luminous IRAS sources, searching for massive protostellar objects in their close vicinities \\cite{bpda, kle0X}. Near IRAS 07029-1215, we discovered a deeply embedded object in a distance of about 1 kpc \\cite{wbr89}, powering a high-velocity bipolar CO outflow. The IRAS source is located in the bright H\\,{\\sc ii} region S 297, it has an IRAS luminosity of 1700 L$_{\\odot}$ \\cite{hen92}. ",
        "conclusions": ""
    },
    "0402/astro-ph0402479_arXiv.txt": {
        "abstract": "{The excursion set model provides a convenient theoretical framework to derive dark matter halo abundances. This paper generalizes the model by introducing a more realistic merging and collapse process. A new parameter regulates the influence of the environment and thus the coherence (non-Markovianity) of the merging and the collapse of individual mass shells.  The model mass function also includes the effects of an ellipsoidal collapse. Analytic approximations of the halo mass function are derived for scale-invariant power spectra with the slopes $n=0,-1,-2$. The $n=-2$ mass function can be compared with the results obtained from the `Hubble volume' simulations. A significant detection of non-Markovian effects is found for an assumed accuracy of the simulated mass function of 10\\%. ",
        "introduction": "The hierarchical growth of virialized cosmic structures provides a useful physical paradigm for the understanding of the formation of galaxies and clusters of galaxies in the Universe (White \\& Rees 1978). The Zel'dovich (1970) theory also includes a description of partially virialized structures like filaments and walls. In both models, cosmic structures grow from initial Gaussian density fluctuations via gravitational instability, leading after merging and their (partial) virialization to mass functions which are used as powerful statistical diagnostics.  Theoretical mass functions of dark matter halos are estimated from N-body simulations with accuracies of 10-30\\,\\% (e.g. Jenkins et al. 2001). More physical insights can be obtained from analytic treatments based on Press-Schechter (PS) -like arguments (Press \\& Schechter 1974; Bond et al. 1991, hereafter BCEK). Alternative derivations are based on, e.g., non-Gaussian statistics (Lucchin \\& Matarrrese 1988), or treat both fluctuations and interactions of density perturbations within one process (Cavaliere \\& Menci 1994), or directly consider the non-linear regime (Valageas \\& Schaeffer 1997). The inclusion of non-spherical dynamical approximations (Monaco 1995, Lee \\& Shandarin 1998) and ellipsoidal collapse models (Sheth, Mo \\& Tormen 2001, hereafter SMT, 2002) modify the assumption of a spherical collapse while preserving the simplicity of the original excursion set idea.  The excursion set model of BCEK assumes a Gaussian density field which is smoothed at a given spatial location with a top-hat filter in wavenumber space (sharp $k$-space filter) using different comoving filter radii $R$. The resulting filtered density contrasts $\\delta(R)$ perform a highly jagged diffusion trajectory (Fig.\\,\\ref{FIG_T}) where the mass function is derived -- as a function of the standard deviation $\\sigma(R)$ of the mass density fluctuations -- from the loss rate of trajectories at the barrier $\\delta_{\\rm c}$ at their lowest $\\sigma$ level (highest $R$ or mass scale). For the spherical collapse model and for the Einstein-de Sitter Universe we have $\\delta_{\\rm c}=1.686$, weakly dependent on cosmology. The halo mass function is thus directly related to the first passage time distribution of the trajectories.  There is a problem with this approach related to the assumption of a sharp $k$-space filter. The filter gives a quite unrealistic mass assignment scheme, with a growth of cosmic structure which depends only on the mass of a halo it has within an infinitesimally small time interval at a given cosmic epoch, and without any dependency on past or future properties of the halo (Markov assumption). Therefore, the merging events occur as completely uncorrelated, sudden, jumps in the formation history. We will replace the sharp $k$-space filter by a non-sharp filter. This leads to a more realistic mass assignment scheme. The corresponding growth of cosmic structure depends on the properties of the halo over a finite and future-directed time range (non-Markov assumption).  The idea to invoke non-Markovian processes is not new. A discussion of the physical consequences of the Markov assumption and why related processes could fail can be found in White (1996, 1997). BCEK pointed out the relation between the shape of a mass filter and the Markov assumption. Their discussion of non-Markovian processes is, however, mainly restricted to results obtained with Monte-Carlo experiments.  Using the traditional excursion set model as a guideline, Sect.\\,\\ref{FIRST} introduces a simple analytic model which describes a more uniform mass assignment scheme, and thus a more uniform spherical collapse of dark matter halos. This model is further improved in Sect.\\,\\ref{ELLIP}, by including the effects of an ellipsodial collapse similar to SMT. The combined model can be regarded as a simple though typical example of a non-Markovian process. It generalizes the traditional excursion set model in a manner such that non-Markovianity can now be gradually increased by a new filter parameter. The same filter parameter also increases the smoothness of the profile of the mass filter (see Eq.\\,\\ref{FILTER} in Sect.\\,\\ref{FF}).  The standard mirror image method is used in Sect.\\,\\ref{MASSF} to derive an analytic form for the halo mass function. We show why this method, which in general does not work for non-Markovian processes, can be used in our specific case. The resulting mass function has the same functional form as the standard excursion set result in terms of the variance of the mass distribution and the critical density threshold. However, non-Markovian effects change the relations between filter radius, mass and variance. These relations become a function of the power spectrum of the underlying mass distribution. After the corresponding transformations of the mass functions they also become apparent in the halo mass function itself. This is the reason why in our non-Markovian context, the profile of the mass filter and thus how much mass is sweeped in from the surrounding mass of a collapsing region becomes a function of the power spectrum.  How much mass the filter sweeps in for a given filter radius can be a quite complex function, especially when general power spectra are considered. We could derive approximate analytic results for scale-invariant power spectra with the slopes $n=0,-1,-2$. The latter case is close to the observed value and allows a comparison with the Jenkins et al. (2001) mass function obtained from the `Hubble Volume' simulations (Sect.\\,\\ref{DISCUSS}). The basic aim is to test under the given assumptions the presence of non-Markovian effects in the simulations and to determine how accurately mass functions should be measured to detect the effects. A discussion of general power spectra goes beyond the scope of our analytic treatment and is postponed to a further paper.  \\begin{figure} \\vspace{-1.7cm} \\centerline{\\hspace{0.0cm} \\psfig{figure=ito3.ps,height=8.5cm,width=8.5cm}} \\vspace{-2.50cm} \\caption{\\small Trajectories for the extended PS process (Markovian process, jagged curves, $T=0$) and for the non-Markovian process (smooth curves, $T=0.23$).} \\label{FIG_T} \\end{figure} ",
        "conclusions": "\\label{DISCUSS}  The process (Eq.\\,\\ref{F2}) corresponds to a smooth filter profile and allows a simple discussion of non-Markovian effects within the excursion set model. The model includes the standard excursion set result as a limiting case, and allows non-Markovian effects to be increased gradually by the new filter parameter $T$.   \\subsection{Effects of non-Markovianity}  Figure \\ref{FIG_MF} shows halo mass functions $f(M)$ for different power spectra and values of the $T$ parameter. In all cases the non-Markovian effects described by Eq.\\,(\\ref{F2}) increase (decrease) the number density of high (low) mass halos. At high masses the ellipsoidality slightly decreases the number density of high mass halos and thus partially compensates non-Markovianity. At small masses both non-Markovianity and ellipsoidality decrease the number density (Fig.\\,\\ref{FIG_MF3}).  The result of equating the mass and the process variances in Eq.\\,(\\ref{EQ}) is a mass filter with a profile which depends on the power spectrum of the mass distribution and thus on the mean profile of density peaks. To show that this is a generic property of non-Markovian processes with increments determined by integrals of the form (Eq.\\,\\ref{MEM}), we generalize the $[1-e^{-(t-s)/T}]$ term in Eq.\\,(\\ref{F2}) by a function $K_T(s,t)$ so that Eq.\\,(\\ref{EQ}) reads \\begin{equation}\\label{G1} t\\,=\\,\\int_0^tds\\,K_T^2(s,t)\\,=\\,\\int_0^\\infty \\frac{4\\pi k^2\\,dk\\,P(k)}{(2\\pi)^2}\\,|W_T(kR)|^2\\,. \\end{equation} Moreover, from the derivation which leads to Eq.\\,(\\ref{RESOL}) we found that the resolution variable $t$ is the same for the Brownian process and our non-Markovian process. The same resolution variable was also obtained for the Ornstein-Uhlenbeck process (Schuecker et al. 2001a). We thus generalize Eq.\\,(\\ref{RESOL}) to a universal resolution variable \\begin{equation}\\label{G2} t\\,=\\,\\int_0^t\\,ds\\,=\\,\\int_0^{1/R} \\frac{4\\pi\\,k^2\\,dk\\,P(k)}{(2\\pi)^3}\\,=\\,s[R,P(k)]\\,. \\end{equation} The combination of Eqs.\\,(\\ref{G1}) and (\\ref{G2}) then suggests the equivalence $K_T(R,P(k))=W_T(kR)$ for $kR\\le 1$ and zero elsewhere. This shows that the kernel of Eq.\\,(\\ref{F1}) which determines the increments of the process is closely related to the profile of the mass filter.  We further conclude that for merging and accretion processes described by the comparatively simple integration scheme (Eq.\\,\\ref{F1}), the mass filter and thus how much mass is sweeped in from the surrounding mass by a collapsing region must also depend on the global mass distribution, and thus on the power spectrum.  Equations (\\ref{VOL1}-\\ref{RM1}) follow a ``natural'' choice of the relation between the filter radius, used in the process of halo detection, and the halo mass as determined by the material which eventually collapse to form a virialized structure. However, other choices are also possible (BCEK) and still a matter of debate.  We further note that the application of filters to all spatial points of a density field, as proposed in all PS-like models, leads to a large scatter between filter mass and group mass. The introduction of non-Markovianity follows the same local filtering scheme and is thus not expected to significantly reduce the scatter on the halo-by-halo level.   \\subsection{Comparison with simulated mass functions}  On the statistical level, PS-like models predict mass functions, merger rates, formation times, biasing schemes etc.  remarkably accurately. A comparison of the theoretical mass functions (Eq.\\,\\ref{F9}) with the results from large and high-resolution N-body simulations (e.g., Jenkins et al. 2001) should thus be more fruitful and should give more important information about the significance of non-Markovian effects.  The Jenkins et al. (2001) mass function can be used for fluctuation fields with effective power spectrum slopes between $n_{\\rm eff}=-2.5$ and $-1$ at $\\sigma=0.5$ in the mass range $10^{11}-10^{16}\\,M_\\odot$. Observations suggest a slope of about $n=-2.0$ (see, e.g., Schuecker et al. 2001b who found $n=-1.8$ on scales $<100\\,h^{-1}\\,{\\rm Mpc}$ for X-ray clusters of galaxies -- quite consistent with the power spectrum of galaxies, see their Fig.\\,16). We thus compare with our $n=-2$ model.  To be consistent with our ``natural'' choice of the radius-mass relation (Eq.\\,\\ref{RM1}), the Jenkins et al. mass function (e.g. their Eq.\\,9) has to be transformed accordingly by using the $\\sigma^2(R(M))$ relation of the non-Markovian process for $n=-2$. To say it in another way, the non-Markovian effects of the process (Eq.\\,\\ref{F1}) can only become apparent after the transformation of $f(\\sigma^2(R))$ to $f(M)$ because non-Markovianity changes the radius-mass assignment scheme. Therefore, when we want to search for non-Markovian effects we have to transform $f(\\sigma^2)$ of both the theoretical and the simulated mass function and test different values of $T$. For this test we multiply the prefactor $16\\pi^2$ in the normalization (Eq.\\,\\ref{VOL3}) by 1.234 to be consistent with the standard $\\sigma_8$ normalization (obtained with the top-hat filter) for $T=0$.  The comparison with the Jenkins et al. (2001) mass function within the mass range $10^{12}-10^{16}\\,M_\\odot$ (for the lower limit see Sect.\\,\\ref{FAIL}) gives the best fit parameter values $T=0.23$ and $\\beta=0.12$ with $\\chi^2=1.4$ for three degrees of freedom (Fig.\\,\\ref{FIG_M2}). The statistical significance of the non-Markovian effects strongly depends on the assumed error of the Jenkins et al. mass function. For errors of 20-30\\%, non-Markovian effects must be regarded as insignificant, whereas for 10\\% errors, non-Markovianity is clearly detected (Fig.\\,\\ref{FIG_X2}). We thus conclude that for the most optimistic error estimates of the Jenkins et al. mass function, its shape suggests the presence of non-Markovian effects, i.e., effects of the environment on the coherence of the collapse. Moreover, non-Markovian effects seem to be not very large and cosmic mass functions with errors better than 10\\% are needed for their clear detection.  \\begin{figure} \\vspace{-0.5cm} \\centerline{\\hspace{0.0cm} \\psfig{figure=chi2_10_a.ps,height=7.0cm,width=7.0cm}} \\vspace{-0.5cm} \\caption{\\small $\\chi^2$ distribution ($1-3\\sigma$ contours for three parameters) in the $\\beta=0.12$ plane, assuming 10\\% errors of the Jenkins et al. (2001) mass function.} \\label{FIG_X2D} \\end{figure}  It should be mentioned that we simultaneously tested in Fig.\\,\\ref{FIG_X2} for the significance of the scaling parameter $a$ introduced by SMT to account for a failure of the simple counting argument of the excursion set formalism (R. Sheth, private communication). The best fit has $a=0.99$ (Fig.\\,\\ref{FIG_X2D}) which means that the inclusion of non-Markovian effects does not require this correction. Therefore, the likelihood contours in Fig.\\,\\ref{FIG_X2} are only plotted in the $a=1$ plane.   \\subsection{Effects of ellipsoidality}\\label{FAIL}  The effects of ellipsoidality on the halo mass function are in general much larger and are detected for all random errors $\\le 30\\%$ with clear significance. The best fit $\\beta=0.12$ is lower than $\\beta=0.3$ obtained from simulations (see Sect.\\,\\ref{FIRST}). In addition, the model starts to deviate from the Jenkins et al. function at masses smaller than $10^{12}\\,M_\\odot$. In order to understand these failures of the model one should not forget that our goal to give a full analytic treatment of the problem forces us to approximate the ellipsoidal collapse by a linear drift. Sheth \\& Tormen (2002) give a simple prescription for approximating the solution to the general barrier in the Markov case which better fits the low-mass range. Therefore, a useful direction for future work is to see if their method works in our particular non-Markovian context also.  Future observed mass functions from e.g. X-ray clusters of galaxies and lensing studies have high enough precision for detailed studies of non-Markovian effects expected in the high-mass regime. On the theoretical side, further studies are in preparation to predict mass functions also for Cold Dark Matter power spectra to improve the comparison with observed and simulated mass functions."
    },
    "0402/astro-ph0402564.txt": {
        "abstract": "With a handful of measurements of limb-darkening coefficients, galactic microlensing has already proven to be a powerful technique for studying atmospheres of distant stars. Survey campaigns such as OGLE-III are capable of providing $\\sim\\,10$ suitable target stars per year that undergo microlensing events involving passages over the caustic created by a binary lens, which last from a few hours to a few days and allow to resolve the stellar atmosphere by frequent broadband photometry. For a caustic exit lasting 12~h and a photometric precision of 1.5~\\%, a moderate sampling interval of 30~min (corresponding to $\\sim\\,$25--30 data points) is sufficient for providing a reliable measurement of the linear limb-darkening coefficient $\\Gamma$ with an uncertainty of $\\sim 8\\,$\\%, which reduces to  $\\sim 3\\,$\\% for a reduced sampling interval of 6~min for the surroundings of the end of the caustic exit. While some additional points over the remaining parts of the lightcurve are highly valuable, a denser sampling in these regions provides little improvement. Unless an accuracy of less than $5\\,$\\% is desired, limb-darkening coefficients for several filters can be obtained or observing time can be spent on other targets during the same night. The adoption of an inappropriate stellar brightness profile as well as the effect of acceleration between source and caustic yield distinguishable characteristic systematics in the model residuals. Acceleration effects are unlikely to affect the lightcurve significantly for most events, although a free acceleration parameter blurs the limb-darkening measurement if the passage duration cannot be accurately determined. ",
        "introduction": "The light we receive from a star originates from layers at different depths from its surface, where the contribution of inner layers decreases from the center towards the limb yielding the phenomenon of limb darkening of the observed surface brightness across the stellar face, which reflects the temperature of the stellar atmosphere as function of distance from the center. Decreasing temperatures from the center to the surface cause a stronger effect for shorter wavelengths, making stars bluer near their center and redder near their limb.  Using caustics of gravitational lenses to measure the brightness profile of closely-aligned background sources was proposed by \\citet{SW1987} with the view toward active galactic nuclei, and several authors have later discussed its application to source stars undergoing galactic microlensing events that involve fold-caustic passages \\citep[e.g.][]{Rhie:LD,GG:detLD}. Making use of both the large total magnification and a strong differential magnification across the stellar face during the caustic passage, galactic microlensing has successfully been demonstrated to be a powerful technique for revealing the brightness profile of distant stars by providing several measurements of limb-darkening coefficients characterizing linear or square-root laws \\citep{PLANET:M28, PLANET:M41, PLANET:O23, joint, PLANET:EB5, Yooetal} with uncertaintes of a few percent.   As discussed in detail by \\citet{Do:Fold}, the lightcurve of a source star in the vicinity of a fold caustic is completely characterized by local properties, so that its brightness profile can be obtained already from data taken around the caustic passage without the need to find a model for the full lightcurve, which describes all properties of the lens system and involves a larger number of parameters along with possible parameter ambiguities. Nevertheless, data outside the caustic passage region provide stronger constraints on the differential blending between observing sites and filters as well as a better assessment of acceleration effects between source and caustic, which can be caused by the parallactic motion of the Earth around the Sun or by orbital motion of the binary lens or a possible binary source. This additional information can be used to reduce uncertainties in the measurement of the stellar brightness profile and its characterizing limb-darkening coefficients.  Some aspects of the measurement of limb-darkening coefficients from microlensing events involving fold-caustic passages have previously been studied by \\citet{Rhie:LD}. This article provides a deeper and furthergoing discussion on this issue and on some general properties of lightcurves of microlensed sources near fold caustics with respect to revealing their brightness profiles.  \\begin{figure*} \\includegraphics[width=168mm]{sensitivities2.eps} \\caption{The sensitity ${\\bmath {\\mathcal T}}(\\eta, \\rho)$ of the fold caustic to the shape of the stellar brightness profile $\\xi$ as a function of the caustic passage phase $\\eta = \\pm\\,(t-t_\\rmn{f}^\\star)/t_\\star^\\perp$ (left) and the fractional stellar radius $\\rho$ (right), as given by Eq.~(\\ref{eq:sensfunc}).} \\label{fig:sensitivity} \\end{figure*}   Sect.~\\ref{sec:lightcurve} addresses the relation between the lightcurve and the stellar brightness profile and discusses the sensitivity of microlensing as function of phase during the caustic passage and fractional radius of the source star. With linear limb darkening as example, Sect.~\\ref{sec:linear} deals with parametrized profiles. A simulation of typical sampled data for different sampling rates is described in Sect.~\\ref{sec:sampled}, which are used in Sect.~\\ref{sec:reconstruction} to investigate the amount of information about the linear limb-darkening coefficient contained in different regions of the lightcurve. In Sect.~\\ref{sec:wrongprofile}, consequences arising from the adoption of an inadequate limb-darkening profile profile are discussed, whereas Sect.~\\ref{sec:acceleration} addresses the influence of an effective acceleration between source and caustic. The paper finishes with a summary of the results and recommendations on the observing strategy in Sect.~\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  As laid out by \\citet{SW1987}, the differential magnification across the face of a source during its passage over a caustic created by an intervening gravitational lens can reveal its radial brightness profile through frequent accurate photometric observations. Galactic microlensing surveys such as OGLE-III are capable of providing $\\sim\\,10$ such passages of stars in the Galactic bulge per year from which a measurement of the stellar brightness profile can be obtained, which reflects the variation of temperature with distance from the center providing a powerful technique for probing stellar atmosphere models.  Contrary to caustic entries, the corresponding exits are predictable from data obtained during the characteristic rise in magnification to a peak \\citep{PLANET:sol,Mao:exit}, which allows to schedule frequent observations before the caustic passage begins.  While the caustic passages are likely to last from a few hours to a few days, photometric measurements with 1m-class telescopes with an uncertainty of less than 1.5~\\% with sampling intervals of a few minutes can be obtained for stars brighter than $\\sim\\,17$th magnitude as currently being carried out by the PLANET collaboration \\citep{PLANET:EGS}, offering the possibility to collect a few hundred data points during the course of the passage.   In the vicinity of the caustic passage, the lightcurve can be described by means of a characteristic profile function $G_\\rmn{f}^\\star(\\eta;\\xi)$, which depends solely on the dimensionless normalized stellar brightness profile $\\xi(\\rho)$ as function of the fractional radius $\\rho$ and the caustic passage phase $\\eta$. $G_\\rmn{f}^\\star(\\eta;\\xi)$ can be seen as the response delivered by the caustic to the specific form of the brightness profile. The weight of the contribution of the brightness at a specific fractional radius $\\rho$ to the caustic profile function $G_\\rmn{f}^\\star(\\eta; \\xi)$ is given by a function ${\\bmath {\\mathcal T}}(\\eta, \\rho)$, which measures the sensitivity of the lightcurve at the point of time that corresponds to the caustic passage phase $\\eta$ to a local variation of the stellar brightness profile at fractional radius $\\rho$. An inspection of ${\\bmath {\\mathcal T}}(\\eta, \\rho)$ shows that the caustic passage provides a one-dimensional scan of the brightness profile where each fractional radius is most efficiently probed as it touches the caustic, which happens twice during the course of the caustic passage, so that the majority of information about outer radii is provided at the beginning and the end of the caustic passage, while inner regions of the source reveal their identity close to times when the source center passes. However, the integrated sensitivity over the full duration of the caustic passage increases with fractional radius.  In general, the extraction of the stellar brightness profile  $\\xi(\\rho)$ from the observed lightcurve $m_\\rmn{fold}(t)$ involves the determination of 5 model parameters ($t_\\rmn{f}^\\star$, $t_\\star^\\perp$, $m_\\rmn{f}^\\star$, $g_\\rmn{f}^\\star$, $\\hat \\omega_\\rmn{f}^\\star$), which relate $m_\\rmn{fold}(t)$ and the caustic profile function $G_\\rmn{f}^\\star(\\eta;\\xi)$, as well as the solution of the integral equation, which relates $G_\\rmn{f}^\\star(\\eta;\\xi)$ and $\\xi(\\rho)$ by means of the caustic sensitivity function ${\\bmath {\\mathcal T}}(\\eta, \\rho)$.  The variation of coefficients characterizing a parametrized stellar brightness profile causes variations in the observed lightcurve through variations of the brightness profile at all fractional radii. The sensitivity of the lightcurve to variations in such a coefficient $\\Gamma$ therefore depends both on the properties of ${\\bmath {\\mathcal T}}(\\eta, \\rho)$ and $\\partial \\xi(\\rho; \\Gamma)/\\partial \\Gamma$. For linear limb darkening, this sensitivity is largest near the beginning of a caustic exit, followed by the surroundings of its end, whereas the smaller integrated sensitivity over the caustic passage for smaller fractional radii and the fact that $\\partial \\xi(\\rho; \\Gamma)/\\partial \\Gamma$ shows little variation with $\\rho$ for the inner parts of the source imply a small sensitivity while these pass the caustic. However, the identification of the end of the caustic exit with the characteristic feature of a jump discontinuity in its slope allows a direct measurement of the corresponding point of time $t_\\rmn{f}^\\star$ and magnitude $m_\\rmn{f}^\\star$, while data points just after the end of the caustic exit can be used to measure the parameter $\\hat \\omega_\\rmn{f}^\\star$ directly, which characterizes the variation of magnification due to non-critical images of the source. For this reason, sampling of the surroundings of the end of the caustic exit is more valuable than of its beginning. Monitoring of both of these regions provides direct measurements of the caustic passage half-duration $t_\\star^\\perp$ and the caustic rise parameter $g_\\rmn{f}^\\star$.  Unless a precision measurement on the linear limb-darkening coefficient $\\Gamma$ is desired, a sampling interval at the limit of the capabilities of the monitoring campaign is not required. For a caustic passage lasting 12~h, a sampling interval of $\\sim\\,$30~min (corresponding to  $\\sim\\,$25--30 data points) is sufficient to provide $\\Gamma$ with an uncertainty of less than $\\sim\\,$8~\\%, while a dense sampling corresponding to an interval of $\\sim\\,$6~min (with $\\sim\\,$150 data points) would reduce the uncertainty to less than $\\sim\\,$3~\\%. This offers the possibility for measuring limb-darkening coefficients in several broadband filters with the same telescope on the same microlensing event or to monitor other microlensing events during the same night (e.g.\\ to look for anomalies caused by planets around the lens stars).  Since coverage of the surroundings of the end of a caustic exit is most effective in providing an accurate measurement of the linear limb-darkening coefficient, it should be the goal of any observation strategy to try to obtain data in this region, where the sampling rate should be chosen with regard to the desired accuracy of the measurement and the number of broadband filters. Moderate sampling over other regions of the light curve is more valuable than an increased sampling over the end of the caustic exit, where a very dense sampling over these other regions however does not provide much additional information but can be used to compensate for a missing coverage of some parts of the caustic passage.   The adoption of an inappropriate stellar brightness profile or the neglect of acceleration effects, which include contributions due to the revolution of the Earth around the Sun or the orbital motion within the binary lens or a possible binary source, leads to well-distinguishable characteristic systematics in the model residuals. Although acceleration is unlikely to cause significant effects on the lightcurve for most events, if used as a free model parameter, however, it blurs the measurement of the stellar brightness profile if the duration of the caustic passage cannot be determined with sufficient accuracy, which is the case if only the surroundings of the end of the caustic exit are sampled. Nevertheless, the related degeneracy can be limited by applying a reasonable upper limit to the absolute amount of acceleration resulting from an assessment of the dynamical properties of the binary lens and the Earth's motion at the time of observation."
    },
    "0402/astro-ph0402147_arXiv.txt": {
        "abstract": "We assume that all the energy loss of the putative pulsar in SN 1987A would contribute to the luminosity of the remnant, which acts like a bolometer. The bolometric luminosity of SN 1987A provides an upper bound on the pulsar's rate of energy loss. An isolated pulsar spinning down by magnetic dipole radiation alone, with initial rotation periods of 10-30 ms, as extrapolated for galactic young pulsars, can have a luminosity below the bolometric bound if either the magnetic field is weak, B $\\sim$ 10$^{9}$-10$^{10}$ G or if it is so strong that the pulsar luminosity decays rapidly, B $\\sim$10$^{16}$ G. ",
        "introduction": "The observed Balmer absorption lines class SN 1987A as a Type II supernova. The identification of the pre-supernova star, the neutrino flux detected at the time of the explosion, the evolution of the optical light curve and the observations of the emission of soft $\\gamma$ rays all confirm this classification. Therefore, it is expected that a neutron star lies at the center of the remnant. Simulations of the neutrino output in core collapse leading to a black hole differ strongly from the Kamiokande neutrino detections, which agree with several neutron star models (Burrows 1988). Any detection of a neutron star formed in SN1987A would arouse great interest since the properties of very young neutron stars are largely unknown. Searches for pulsed emission from the remnant have yielded no results (\\\"{O}gelman et al. 1990). A claimed pulsar period of 2.14 ms has been discussed by Middleditch et al (2000). A synchrotron nebula would indicate a young pulsar even if its beam does not sweep our direction so that a pulsed signal is not received. No synchrotron nebula has been detected in the remnant of SN 1987A, with a Chandra upper limit of 5.5 $\\times$ 10$^{33}$ ergs s$^{-1}$ to the luminosity of a central X-ray source (Park et al. 2003).  The newborn neutron star is characterized by its initial rotation rate and magnetic dipole moment. It will spin down, generating radiation and accelerating charged particles at the expense of its rotational energy. The power output of the neutron star is the luminosity of a rotating dipole, regardless of the detailed processes by which the dipole emission is converted to the energy of charged particle acceleration and high-enegy electromagnetic radiation. The remnant is optically thick to the radiation emitted by the young pulsar. Thus, the remnant will act as a bolometer, absorbing and reprocessing the pulsar luminosity. Most of the remnant's central luminosity is in the optical, IR and UV bands, exceeding its X-ray luminosity by several orders of magnitude.  The observed bolometric luminosity of the remnant is dominated and well fitted with the radioactive decay luminosity, so the contribution from the pulsar must be a small admixture. We shall assume that the pulsar loses energy as an isolated rotating dipole, of constant magnetic dipole moment, without any other torques (as would arise, for example, from the presence of a fallback disk). We use the bolometric luminosity to place bounds on the dipole emission of the putative pulsar and on its birth properties. ",
        "conclusions": "By using the bolometric luminosity of the central remnant in SN 1987A, one can obtain interesting constraints on the properties of the neutron star at birth. The constraints on initial rotation period and on magnetic dipole moment are not independent of each other. To place the options in context, let us recall the possibilities presented by the conventional interpretation of the radio pulsar population.  Regarding the initial periods, the population of ``fresh\" normal radio pulsars (as distinct from the ``recycled\" binary and millisecond radio pulsars (Alpar et al. 1982, Radhakrishnan \\& Srinivasan 1982) implies P$_0 \\sim$ 10-30 ms, by extrapolation from the youngest pulsars. An alternative suggestion, particularly for the possible pulsar in SN 1987A, is that it was formed by a core merger (Chen \\& Colgate 1995, Middleditch et al. 2000). The pulsar is expected to have an initial period in the millisecond range, and millisecond pulsars in general can be young objects according to this scenario.  The suggestion that some pulsars are born with slow rotation rates to explain the distribution in the P-$\\dot{P}$ diagram has proved difficult to check conclusively because of a complex of selection effects and has led to conflicting results. Slow initial rotation rates are not supported by the later work (Lorimer et al. 1993). The distribution in the P-$\\dot{P}$ diagram can be obtained if pulsars are born with short initial periods of 10-30 ms and evolve under varying torques that include contributions other than the dipole radiation torque at some epochs in the pulsars' evolution. This is supported by the observation of braking indices less than the dipole braking index n=3 (Kaspi et al. 2001) and ages that differ significantly from the characteristic age P / 2$\\dot{P}$ of dipole spindown (Gaensler \\& Frail 2000).  The dipole magnetic moments of the majority of radio pulsars are on the order of 10$^{30}$ G cm$^3$ (B $\\sim$ 10$^{12}$ G). There is no evidence for decay of the magnetic moment during the normal radio pulsars' active lifetimes (Bailes 1989, Bhattacharya et al. 1992). There are a few pulsars in the tail of the distribution, with dipole magnetic moments in the 10$^{31}$ G cm$^3$ range (Camilo et al. 2000, Morris et al. 2002, McLaughlin et al. 2003) and surface dipole magnetic fields in the magnetar range, higher than the quantum critical field B$_c \\equiv ({m_{e}}^2 c^3)/(e \\hbar)$ = 4.4 $\\times$ 10$^{13}$ G. PSR J1847-0130 (McLaughlin et al. 2003) has the highest dipole moment, $\\mu$ = 9.4 $\\times$ 10$^{31}$ G cm$^3$, corresponding to a dipole magnetic field of 9.4 $\\times$ 10$^{13}$ G on the magnetic equator and 1.9 $\\times$ 10$^{14}$ G on the poles. These pulsars lie in the upper right hand corner of the P-$\\dot{P}$ diagram, among the anomalous X-ray pulsars (AXPs) and the soft gamma-ray (burst) repeaters (SGRs) for which the magnetar model has been developed (see Thompson 2000 for a review). While the magnetar model successfully addresses many properties of the SGRs and AXPs, the period clustering and the presence of radio pulsars with periods and inferred magnetic dipole moments similar to those of AXPs and SGRs has motivated another class of models involving fall-back disks and a combination of dipole radiation and disk torques (Alpar 2001, Chatterjee, Hernquist \\& Narayan 2000, Ek\\c{s}i \\& Alpar 2003). In these models the magnetic moments of pulsars like PSR J1847-0130, and of AXPs and SGRs can be in the conventional 10$^{30}$ G cm$^3$ range. The absence or presence and mass of such disks introduces a third neutron star parameter at birth, in addition to P$_0$ and $\\mu$. We shall assume here that the possible pulsar in SN 1987A has no disk or a light enough disk mass that its early evolution is determined by magnetic dipole radiation, as must be the case for the typical radio pulsars. An analysis of the present bolometric luminosity constraints allowing for fallback disks in the initial conditions will be the subject of separate work. Inferred dipole magnetic moments of 10$^{27}$ - 10$^{28}$ G cm$^3$ are typical of the observed millisecond pulsars. The recycling hypothesis (Alpar et al 1982, Radhakrishnan \\& Srinivasan 1982) is supported by the locations of binary and millisecond pulsars in the P-$\\dot{P}$ diagram, and particularly by the discovery of millisecond rotation periods from low mass X-ray binaries (Wijnands \\& van der Klis 1998), of which now five are known. These links, as well as the abundance of millisecond and binary pulsars in globular clusters, suggest that the radio pulsars with such weak inferred magnetic dipole moments are from an old population. The possible pulsar in SN 1987A is not expected to have a magnetic moment on the order of 10$^{28}$ G cm$^3$ if only old pulsars can have such weak fields. However, if the pulsar was born as a result of a core merger, allowing it a millisecond period at birth, then the magnetic field could be weak. Another possibility is that pulsars are born with weak initial magnetic fields, and the dipole magnetic fields of conventional $\\mu \\sim$ 10$^{30}$ G cm$^3$ pulsars and of magnetars are generated subsequently on timescales longer than the 16 yr time span of SN 1987A observations (see Reisenegger (2003) for a review and Michel (1994) for applications to the possible pulsar in SN 1987A ).  Allowing all pulsar birth scenarios with appropriate choices of initial periods, we found that for 10, 15, and 30 ms initial periods, the upper limits on weak initial magnetic moments are 2.8 $\\times$ 10$^{27}$, 6.4 $\\times$ 10$^{27}$, or 2.5 $\\times$ 10$^{28}$ G cm$^3$. With the assumption of a slow rotator, say with an initial period of 0.3 s, it was found that the constraint at the weak field end was not very stringent, $\\mu < 2.5 \\times$ 10$^{30}$ G cm$^3$. If the pulsar was born from a core merger, with an initial period of 2 ms, then the upper limit on a weak magnetic moment is as low as $\\mu < 1.1 \\times$ 10$^{26}$ G cm$^3$.  For all choices of initial period considered above, if the initial magnetic moment is not weak, and if the putative pulsar is spinning down under a constant magnetic dipole radiation torque, without field generation, we conclude that the putative pulsar has a magnetic dipole moment $\\mu$ greater than 2.4 $\\times$ 10$^{34}$ G cm$^3$. If weak initial magnetic moments can be ruled out for all young pulsars, then we find that the possible pulsar in SN 1987A has to be a strong magnetar. If that is the case, then we may need to change our strategy for detecting the putative pulsar and look for its transient features, such as bursts observed from anomalous X-ray pulsars and soft gamma-ray repeaters, if the uniting property of all these sources is the possession of magnetar fields."
    },
    "0402/astro-ph0402237_arXiv.txt": {
        "abstract": "A generic prediction of hierarchical gravitational clustering models is that the distribution of halo formation times should depend relatively strongly on halo mass, massive haloes forming more recently, and depend only weakly, if at all, on the large scale environment of the haloes.  We present a novel test of this assumption which uses the statistics of weighted or `marked' correlations, which prove to be particularly well-suited to detecting and quantifying weak correlations with environment.  We find that close pairs of haloes form at slightly higher redshifts than do more widely separated halo pairs, suggesting that haloes in dense regions form at slightly earlier times than do haloes of the same mass in less dense regions. The environmental trends we find are useful for models which relate the properties of galaxies to the formation histories of the haloes which surround them. ",
        "introduction": "The excursion set model of hierarchical clustering (Epstein 1983; Bond et al. 1991) has been remarkably successful.  It provides useful analytic approximations for the abundance of haloes of mass $m$ at time $t$ (Bond et al. 1991; Sheth, Mo \\& Tormen 2001), for the conditional mass function of $m$ haloes at $t$ which are later (at $T>t$) in more massive haloes $M>m$ (Bond et al. 1991; Lacey \\& Cole 1993; Sheth \\& Tormen 2002), for the abundance of haloes as a function of the larger scale environment (Mo \\& White 1996; Lemson \\& Kauffmann 1999; Sheth \\& Tormen 2002) for the distribution of halo formation times (Lacey \\& Cole 1993) and masses (Nusser \\& Sheth 1999; Sheth \\& Tormen 2004).  Here, formation is typically defined as that time when the most massive progenitor contains at least half the final mass.  In the simplest and most used approximation, this approach ignores most correlations between different spatial scales.  In this approximation, the approach predicts that there should be no correlation between halo formation and the large scale environment in which the halo sits (White 1996).  This is because, in the model, formation refers to a smaller mass than the final virial mass, and hence to a smaller spatial scale than that associated with the Lagrangian radius of an object, whereas the larger scale environment, by definition, refers to scales which are larger than that of the halo.  Lemson \\& Kauffmann (1999) presented evidence from measurements in numerical simulations of clustering that halo formation times were indeed independent of environment.  They interpreted this as evidence that the excursion set neglect of correlations was justified. (Lemson \\& Kauffmann also presented evidence that a number of other physical properties of haloes were also independent of environment, and this evidence has been used to justify an assumption which enormously simplifies semi-analytic models of galaxy formation: that the properties of galaxies are determined by the haloes in which they form, and not by the surrounding larger-scale environment.) Their conclusion is somewhat surprising for the following reason. It is quite well established that the ratio of massive to low mass haloes is larger in dense regions, and that the excursion set model is able to quantify this dependence quite well (see the references given earlier). It is also well established that, on average, low mass haloes form at higher redshifts (see references given earlier). Together, these suggest that if one averages over the entire range of halo masses in any given region, then the mean formation time in dense regions should be shifted to lower redshifts, simply because these regions contain more massive haloes which, on average, form later.  In Figure~4 of their paper, Lemson \\& Kauffmann averaged over the entire range of halo masses accessible to them in their simulations, and found no dependence of formation time on environement; at face value, this is {\\it in}consistent with the simplest excursion set prediction!  The main goal of this paper is to repeat the test for environmental effects on halo formation.  Section~\\ref{density} shows that a simple plot of formation time versus local density does not show strong trends, suggesting that the excursion set approximation is rather accurate. But then, Section~\\ref{marks} presents evidence, from what we feel is a more senstive test, which indicates that low mass haloes in dense environments form slightly earlier than haloes of the same mass in less dense environments.  Thus, it may be that, when one averages over a range of halo masses, the shift to later formation times associated with the fact that the dense regions contain the most massive haloes is approximately compensated-for by a density-dependent shift to slightly earlier formation times, with haloes in denser regions having larger formation redshifts than their counterparts (of the same mass) in the field.  This second test uses a technique known as marked correlation functions:  our results indicate that marked correlation functions are a powerful means of detecting and quantifying environmental dependences.  Section~\\ref{alt} illustrates that our results do not depend sensitively on what definition of halo formation we use.  A final section summarizes our findings, and argues that our results may have important implications for studies of halo structure.  It also presents evidence which suggests that the density profiles of close halo pairs are neither more nor less centrally concentrated than are the profiles of their counterparts in less dense regions. ",
        "conclusions": "\\label{discuss} We presented evidence that halo formation is weakly correlated with the surrounding density field.  The weakness of the correlation suggests that the usual excursion set model neglect of correlations is a good approximation, but the fact that a correlation does exist means that, as the available data on galaxy properties and the surrounding large scale structure becomes more precise, it will become necessary to build a more sophisticated model.  Our demonstration that environmental effects do leave a mark on the halo population, and that marked correlation functions are a useful method of detecting and quantifying this, suggests that a marked correlation analysis of halo concentrations, spins, shapes and alignments should yield interesting results. For instance, the particular definition of halo formation we chose to study in Section~\\ref{alt} is motivated by the work of Navarro, Frenk \\& White (1997) who argue that $z_{0.01}$ correlates strongly with the central concentration of the parent halo. Since our analysis indicates that halo formation depends on environment, it is plausible that the structural properties of haloes will also depend on environment.   As a first step, we have attempted a measurement which uses the halo concentration as the mark.  This is not entirely straightforward, because the GIF simulations we have used in this paper do not have particularly good mass-resolution, so they are not well suited to estimating the density run around halo centres; as a result, the halo concentrations we estimate are rather noisy.  Nevertheless, if we select a sample for which the Navarro et al. formula is a better fit, then we find reasonable agreement with previous higher resolution studies (e.g. Jing 2000) which were restricted to considerably smaller halo samples.  Specifically, we find that on average, massive halos have slightly smaller concentration parameters:  $\\langle\\ln c\\rangle = \\ln[c_*\\,(M/M_*)^{-0.1}]$, with $c_*=8$ and 9 at $z=0$ in the SCDM and $\\Lambda$CDM runs. More importantly, the distribution around the mean concentration is approximately independent of halo mass, and is well approximated by a lognormal with rms $\\sigma_{{\\rm ln}\\,c}=0.3$. This is shown in Figure~\\ref{plnc}.  \\begin{figure} \\centering \\epsfxsize=\\hsize\\epsffile{st2004f10.eps} \\vspace{-4cm} \\caption{Similar to Figures~\\ref{xiwf} and~\\ref{xiwnfw}, but now with scaled concentration $C$ as the mark.  There is no evidence that, once mass dependent trends have been removed, halo concentrations depend on environment.  It will be interesting to see if this conclusion remains when this analysis is repeated with higher resolution simulations.  } \\label{xiconc} \\end{figure}  Because the same lognormal shape provides a good description of the distribution of concentrations at all masses, we can use it to remove mass-dependent trends, as we did in our study of halo formation in main text.  Figure~\\ref{xiconc} shows a marked correlation function in which this scaled concentration was used as the mark.  There is no evidence that the concentrations of close pairs are any different than those of more widely separated pairs, suggesting that, once mass-dependent trends have been removed, halo concentrations do not depend on environment.  Since this strongly suggests that the connection between halo formation and concentration is not as straightforward as is generally assumed, it will be interesting to see if this conclusion remains when the analysis is repeated on higher resolution simulations.  Although we believe a marked correlation analysis of halo concentrations, spins, shapes and alignments will yield interesting results, we think that there is one measurement which is more interesting still. Recent studies of pure dark-matter simulations suggest that galaxies are likely associated with the substructure components of dark matter haloes (Kravtsov et al. 2003).  A marked correlation correlation analysis which uses the number of subclumps as the mark would be extremely interesting, because one of the crucial assumptions in current interpretations of the galaxy correlation function is that the number of galaxies which form in a halo depends on halo mass, but does not depend on halo environment.  For similar reasons, a marked correlation analysis of the sites in which gas cools in dark-matter plus hydro- simulations will be very interesting, particularly in view of the fact that conventional tests find little evidence for environmental trends (Berlind et al. 2003).  Finally, it is worth mentioning that models based on the work of Efstathiou \\& Rees (1988) which relate the formation of massive black holes to halo formation assume that, at fixed mass, the clustering of halos is independent of whether or not they formed recently.  Measurements in simulations of the large-scale clustering of merger sites indicate that, on scales larger than a few Mpc, this assumption is accurate (Percival et al. 2003). Similarly, simulations show that the large-scale cross-correlation between halos that formed recently and the entire halo population is similar to the auto-correlation function of the entire halo population (Kauffmann \\& Haehnelt 2002). Both measurements indicate that, at least on large scales, the only trends with environment are those which arise from the correlation between mass and environment, consistent with the simplest excursion set prediction.  Although they emphasized what they saw on large scales, on smaller scales both Kauffmann \\& Haehnelt (2002) and Percival et al. (2003) do see weak evidence for small additional trends with environment.  Their measurements can be reconciled with the excursion set approach if we recall that the excursion set prediction is based on the assumption that different scales are uncorrelated---in the jargon, this comes from using a smoothing filter which is sharp in $k$-space. The predictions of the excursion set approach do depend on the choice of filter.  However, the precise choice of filter cannot matter on scales larger than the correlation length.  Thus, the predicted importance of halo mass rather than environment {\\em on large scales}, although it is derived from consideration of the sharp $k$-space filter, is a generic prediction of the approach.  On the other hand, one does expect filter-dependent environmental trends on smaller scales.  Since the sharp $k$-space filter is not expected to be a reasonable choice on small scales, one generically expects to find correlations with environment on small scales (over and above those associated with halo mass).  Presumably, it is these correlations which our analysis is well-suited to detecting.   \\bigskip  We would like to thank the Aspen Center for Physics for support, and for providing the stimulating environment in which this work was begun. We would also like to thank the Virgo consortium for making the simulation data used here publically available at {\\tt http://www.mpa-garching.mpg.de/Virgo}, and Andy Connolly, Bob Nichol and the other members of the Pittsburgh Computational Astrostatistics (PiCA) Group for providing a fast code with which to evaluate marked correlation functions. This work was supported by NASA grant NAG5-13270, and by an FR Type II Grant from the University of Pittsburgh."
    },
    "0402/astro-ph0402551_arXiv.txt": {
        "abstract": "{ The long-term variability of the multiwavelength blazar emission can be interpreted in terms of orientation variations of a helical, inhomogeneous, non-thermally emitting jet, possibly caused by the orbital motion of the parent black hole in a binary system (Villata \\& Raiteri \\cite{vilrai99}). The helical-jet model is here applied to explain the quasi-periodic radio-optical light curves and the broad-band spectral energy distributions (SEDs) of the BL Lac object \\object{AO 0235+16}. Through a suitable choice of the model parameters, the helix rotation can well account for the periodicity of the main radio and optical outbursts and for the corresponding SED variability, while the interspersed minor radio events could be interpreted as due either to some local distortions of the helical structure or to other phenomena contributing to the source emission. In particular, the probable existence of flow instabilities provides a viable interpretation for the non-periodic features. ",
        "introduction": "Blazars, namely BL Lacertae objects and flat-spectrum radio quasars, belong to the class of active galactic nuclei (AGNs). According to the unified model of AGNs (see e.g. Urry \\& Padovani \\cite{urrpad95}), their emission is dominated by radiation produced in relativistic plasma jets, which are oriented at small angles with respect to the line of sight; the radiation is therefore strongly beamed towards the observer.  Details of the jet origin are still unknown, but their formation is thought to be triggered by the presence of a supermassive black hole surrounded by an accretion disc, possibly belonging to a binary black hole system (BBHS) hidden in the centre of the AGN. On one hand, this scenario is suggested by the elliptical morphology of BL Lac host galaxies, which are believed to originate from merging phenomena between spirals, naturally leading to the formation of massive binary black holes (Begelman et al. \\cite{beg80}; Wilson \\& Colbert \\cite{wilcol95}; Huges \\& Blandford \\cite{hug03}). On the other hand, several observational evidences such as bending, misalignment, wiggling and precession of jets (Begelman et al. \\cite{beg80}; Camenzind \\& Krockenberger \\cite{camkro92}; Kaastra \\& Roos \\cite{kaaroo92}; Conway \\& Wrobel \\cite{conwro95}; Villata et al. \\cite{vil98}; Villata \\& Raiteri \\cite{vilrai99} and references therein; Abraham \\cite{abr00}; Britzen et al. \\cite{bri01}), often associated with knots superluminally moving along different-scale curved trajectories, as well as the periodicity discovered in the multiwavelength light curves of some of these sources, have in some cases been interpreted in terms of helical structures tightly related to BBHSs.  The helical-jet model proposed by Villata \\& Raiteri (\\cite{vilrai99}) describes how orientation variations of the different-frequency emitting parts of a curved jet with respect to the line of sight, possibly caused by the orbital motion of the parent black hole in a BBHS, can be the cause of the observed changes in the spectral energy distribution (SED) of a blazar. The emission from the jet is non-thermal: the relativistic electron population is responsible for producing both the low-energy (from radio to UV--X-rays) synchrotron radiation and the high-energy (up to $\\gamma$-rays) one through inverse-Compton (IC) scattering of the same synchrotron photons (synchrotron-self-Compton process, hereafter SSC).  This model has already provided an interpretation for the huge X-ray spectral brightening of \\object{Mkn 501} (Villata \\& Raiteri \\cite{vilrai99}), the low-energy SED variations of \\object{S4 0954+65} (Raiteri et al. \\cite{rai99}), and the changes in the overall SEDs of \\object{S5 0716+71} (Ostorero et al. \\cite{ost01}) and \\object{ON 231} (Sobrito et al. \\cite{sob01}). The modelling of a constant helix rotation allows us now to simulate the long-term behaviour of the multifrequency light curves together with the SED time evolution. If the data sampling is good enough and the source emission exhibits a periodic (or quasi-periodic) behaviour, one can apply the model and find the set of parameters which best reproduces the observed multifrequency data, thus providing both a test for the model and a description of the corresponding jet features.  In this paper the model, fully described in Sect.\\ 2, is applied to the case of the BL Lac object \\object{AO 0235+16}, whose radio (and optical) light curves have recently revealed a $\\sim 5.7$ year quasi-periodicity (Raiteri et al. \\cite{rai01}). Model light curves and SEDs are compared with observational data in Sect.\\ 3. Possible interpretations of minor, non-periodic events are presented in Sect.\\ 4. A final discussion is performed in Sect.\\ 5. ",
        "conclusions": "We found that the helical-jet model is able to describe the long-term behaviour of the multiwavelength emission of the BL Lacertae object AO 0235+16. Both the periodic occurrence and the mean shape of the main radio and optical outbursts, as well as the corresponding remarkable variations of the broad-band SED, can be explained in terms of the orientation change of an inhomogeneous, steadily-emitting, rotating helical jet. The radiation from the different-frequency emitting regions of the jet is affected by relativistic beaming, whose amount depends on the angle between the velocity vector of the emitting plasma and the line of sight, which changes along the helical path and also varies with time. The assumption of flow instabilities (or other phenomena implying local emission enhancement) in the jet provides a viable interpretation for the non-periodic outbursts observed in the radio light curves.  The twisting of the jet in a helical structure and its rotation can originate from the orbital motion of the parent black hole in a BBHS, the main signature of which would indeed be the periodicity of the source light curves.  The periodicity of AO 0235+16 is about 2.9 years in the host galaxy rest reference frame, taking into account the $P=P_{\\mathrm{obs}}/(1+z)$ relation with $z=0.94$. According to Begelman et al. (\\cite{beg80}), this period (assumed to be the orbital period) enables to estimate the mass of the primary black hole for any given value of the mass ratio $M/m$ between the primary and secondary components: $M \\sim P_\\mathrm{yr}^{8/5}(M/m)^{3/5}\\,10^6\\,M_{\\sun}$. For $M/m \\sim$ 1--100 one can infer $M \\sim 5 \\times 10^6\\,$--$\\,9 \\times 10^7\\,M_{\\sun}$; the binary separation would be in the range $2\\times 10^{-3}\\,$--$\\,5\\times 10^{-3}\\,\\mathrm{pc}$. If the jet is emitted by the primary, its orbital radius could vary from $5\\times 10^{-5}\\,$ up to $\\,10^{-3}\\,\\mathrm{pc}$.  Several VLBI images of the source were produced over a wide range of observing radio wavelengths in the past years. A large fraction of these maps shows no evidence of extended structures apart from the compact core, regardless of the different resolution (Gabuzda et al. \\cite{gab92}; Gabuzda \\& Cawthorne \\cite{gabcaw96}); on the other hand, some of them reveal a faint jet north of the core (Jones et al. \\cite{jon84}; Chu et al. \\cite {chu96}; Shen et al. \\cite{she97}). The presence of a weak extension is confirmed by the sub-milliarcsecond maps obtained with the VSOP at 5 GHz (Frey et al. \\cite{fre00}) and, more recently, with the VLBI at 43 GHz (Jorstad et al. \\cite{jor01}), the latter also displaying a couple of components superluminally moving along bent trajectories.  Our model interprets the jet radio knots as the jet regions where the helical pattern presents the minimum viewing angle, with a maximization of the beaming effect. According to the above orbital radius estimate and taking the pitch angle $\\zeta$ and the viewing angle of the helix axis $\\psi$ into account, the observed separation between radio knots should be less than 10 $\\mu\\mathrm{as}$, well below the resolution of the most detailed available maps. All this under the assumption that the helix pitch does not vary when moving from the ``one-turn'' emitting region we considered, which is likely not true: the helix pitch could be smaller close to the black hole and growing outside, up to the observed sub-mas scales.  A huge observing effort is currently ongoing on this source, in order to closely follow its variability behaviour around the time of the next predicted outburst (first half of 2004; Raiteri et al. \\cite{rai01}): optical and radio telescopes of the WEBT collaboration ({\\tt http://www.to.astro.it/blazars/webt/}; Villata et al. \\cite{vil00}, \\cite{vil02}) are intensively monitoring it since summer 2003, together with the Effelsberg 100 m radio telescope and VLBA. Moreover, the ground-based observing effort will be intensified in 2004, during the optical/UV/X-ray pointings of the source by the instruments onboard XMM-Newton, which will provide more details on the shape of the SED in the UV--X-ray band, possibly sheding light on the origin of X-rays in AO 0235+16."
    },
    "0402/astro-ph0402598_arXiv.txt": {
        "abstract": "{ We present the general analytic solution for the evolution of radiative supernova remnants in a uniform interstellar medium, under thin-shell approximation. This approximation is shown to be very accurate approach to this task. For a given set of parameters, our solution closely matches the results of numerical models, showing a transient in which the deceleration parameter reaches a maximum value of 0.33, followed by a slow convergence to the asymptotic value 2/7. Oort (1951) and McKee and Ostriker (1977) analytic solutions are discussed, as special cases of the general solution we have found. ",
        "introduction": "In recent years numerical modelling of the structure and evolution of supernova remnants (SNRs) has reached an unprecedented level of detail. Nonetheless analytic models still play a very important role, when general properties have to be investigated, as well as when direct relations have to be drawn between pure observational quantities (like size, flux, etc.) and intrinsic physical parameters.  The adiabatic phase of SNR evolution (in a uniform and homogeneous medium) is well described by the Sedov (\\cite{Sedov-59}) analytic solution, which reproduces both the SNR radial evolution and its inner structure. This exact solution has been made possible by the fact that during this phase the SNR evolution is self-similar. This is no longer the case when radiative losses become important, and therefore no exact analytic solution is known for the late SNR evolution.  Approximated solutions in the adiabatic regime and beyond may be also obtained using a ``thin-shell'' model (see e.g.\\ Zel'dovich \\& Raizer \\cite{Zel-Raizer-66}, Ostriker \\& McKee \\cite{Ostr-McKee-88}). This approach assumes that the whole mass (and therefore kinetic energy) of the SNR is located in a rather thin shell just behind the outer shock; while the inner region is filled with a very hot and rarefied gas, of negligible total mass, but containing most of the SNR internal energy.  For the adiabatic phase this approximation is only moderately accurate (see e.g.\\ Zel'dovich \\& Raizer \\cite{Zel-Raizer-66}). In fact, according to the Sedov solution the gas density vanishes in the inner regions while its pressure keeps finite; however, the outer layer containing most of the mass is geometrically rather thick. On the other hand, numerical works (e.g.\\ Falle \\cite{Falle-75}, Blondin et al.\\ \\cite{Blondin-et-al-98}, hereafter BWBR) trace the formation of a much thinner shell in the radiative phase, therefore indicating that a thin-shell approximation should be far more accurate in describing the late evolution.  Oort (\\cite{Oort-51}) presented a first thin-shell approach to a radiative SNR expansion. By assuming momentum conservation in the shell, he found the SNR radius to evolve as $R\\propto t^{1/4}$. This solution, also known as ``momentum-conserving snowplow'', assumes that cooling is extremely efficient everywhere (and therefore that the interior pressure vanishes). However, numerical models (e.g.\\ Chevalier \\cite{Chev-74}) show that, even in the radiative phase, the gas in the central regions becomes so rarefied that its cooling time still keeps considerably longer than the SNR age. This led McKee \\& Ostriker (\\cite{McKee-Ostr-77}) to introduce a ``pressure-driven snowplow'' model, in which a fossil pressure in the hot interior has a substantial dynamical effect on the outer shell: in this case the radial evolution is $R\\propto t^{2/7}$ (for adiabatic index $\\gm=5/3$).  Even though the ``pressure-driven snowplow'' formula gets closer than the ``momentum-conserving snowplow'' one to the numerical results, some discrepancy still remains. For instance, by defining the ``deceleration parameter'' as $m=d\\log R/d\\log t$, numerical models obtain an asymptotic value ranging from 0.31 (Chevalier \\cite{Chev-74}) to 0.33 (BWBR). These values are significantly different from the analytic value, $2/7$ (namely 0.286), and various authors have discussed the origin of such discrepancy. Cioffi et al.\\ (\\cite{Cioffi-et-al-88}) ascribe it to a ``memory'' of the previous Sedov phase, leading to an actual internal pressure larger than that derived from the analytic model. BWBR, instead, attribute this discrepancy to the influence of the reverse shock, which moves towards the center raising the thermal energy, thus leading to a milder deceleration.  Other authors have estimated analytically the radial evolution under more general conditions than those given above.  Ostriker \\& McKee (\\cite{Ostr-McKee-88}) have shown that, for a general $\\gm$ as well as a power-law ambient density profile ($\\rhoa(r)\\propto r^{-\\om}$), $m=1/(4-\\om)$ for a ``momentum-conserving snowplow'', while $m=2/(2+3\\gm-\\om)$ for a ``pressure-driven snowplow''.  Liang \\& Keilty (\\cite{Liang-Keilty-2000}) have considered the case in which only a (constant) fraction $\\eps$ of the kinetic energy of the incoming flow is radiated in the outer shock. For $\\gm=5/3$, $m$ is found to decrease quasi-linearly with $\\eps$, from $2/5$ for the adiabatic case ($\\eps=0$) to $2/7$ for the fully radiative case ($\\eps=1$); and a value of $\\eps$ of about 0.8 (0.6) is required in order to obtain $m=0.31$ (0.33), as indicated by the numerical models. However, while $\\eps<1$ may be appropriate to describe gamma-ray burst afterglows (Cohen et al.\\ \\cite{Cohen-98}), SNR radiative shocks should be described as fully radiative ones (namely with $\\eps$ very close to unity).  The effect of cooling in the hot interior on the deceleration parameter has been studied by Gaffet (\\cite{Gaffet-83}), with the following results. Adiabaticity holds throughout most of the volume occupied by the hot gas, while cooling occurs only near the boundary with the radiative shell, giving as effect a net mass transfer from the hot interior to the shell. Assuming that the gas in the hot interior follows a cooling law $\\Lambda\\propto T^{-c}$, this paper discusses different regimes for different choices of $\\gm$ and $c$, showing that the asymptotic value of $m$ must be in the range between the values $1/4$ (Oort limit) and $2/(2+3\\gm)$ (McKee and Ostriker limit).  When $c>(2/3)(5\\gm-8)/(3-2\\gm)$ ($c>-2/3$ for $\\gm=5/3$) the asymptotic value of $m$ is 1/4. This is the case for cooling functions typical of the SNR regime, where $c>0$ and is usually taken in the range from 0.5 (e.g.\\ Cioffi et al.\\ \\cite{Cioffi-et-al-88}) to 1.0 (BWBR). Therefore, according to this result, after the commonly known radiative phase there could be a very late evolutive phase, in which radiative losses of hot interior begin to be prominent and the SNR evolves as ``momentum-conserving snowplow''. However, as it may be derived from numerical results, in a typical SNR, radiative cooling of the hot interior is negligible until very late times. Therefore the onset of the ``momentum-conserving snowplow'' regime should occur only near the end of a SNR lifetime, or not to occur at all.  A common limitation of all above-mentioned analytical models is that the radial evolution of the radiative shock has been approximated by a power-law behaviour $R\\propto t^{m}$ (with constant $m$). This allows a simplified treatment of radiative SNR evolution; however, it is natural to expect that a power-law expansion occurs only at late times (i.e.\\ at large $R$ values), after the transition from adiabatic to radiative expansion has been completed.  In this paper we shall show: 1) that the quoted difference between the numerical and asymptotic analytic value is just a consequence of the fact that the time needed to reach the asymptotic power-law regime is long compared with the age of the SNR; 2) that the SNR radial evolution during the transient phase is adequately described by a thin-shell model; 3) that a general analytic solution of this problem exists.  The plan of the paper is as follows. In Sect.~2 we derive a differential equation that describes the radiative SNR evolution, find its complete solution for an arbitrary adiabatic coefficient, discuss the general properties for the two branches of solutions that we find, and show that the Oort solution is just a special case of our general solution. Sect.~3 focuses on the conventional case $\\gm=5/3$, which allows simpler analytic formulae, and shows that our solution tends to the asymtotic regime given by the solution of McKee \\& Ostriker (\\cite{McKee-Ostr-77}); with the use of numerical results, we also derive the most appropriate initial conditions for our solution. Sect.~4 concludes. ",
        "conclusions": "Using a thin-shell approach, we have developed the definitive analytic treatment for the evolution of a SNR in the radiative phase, and we have also obtained a series of interesting relations. The main findings of the present work are the following.  The discrepancy between the analytic prediction of the asymptotic value of the deceleration parameter ($m=2/7$, McKee and Ostriker \\cite{McKee-Ostr-77}) and that derived numerically ($m=0.33$, BWBR) is only apparent. This discrepancy has been attributed to the presence of a reverse shock moving towards the center. We show, instead, that a thin-shell model, that by definition does not contain any information on inner structure details, closely fits the SNR evolution as derived numerically.  We then confirm that $2/7$ is the correct asymptotic value, even though the convergence towards this value is expected to be slow. We believe that, if BWBR numerical simulation had been runned until later stages of the SNR evolution, it would have shown that $m$ does not keep constant to $0.33$, but eventually decreases. This has been already pointed out by Chevalier (\\cite{Chev-74}) and can be seen in Fig.~5 of Cioffi et al.\\ (\\cite{Cioffi-et-al-88}), in Fig.~3 of Falle (\\cite{Falle-81}), and in Fig.~4b of Mansfield \\& Salpeter (\\cite{Mansf-Salp-74}). However, the convergence to the asymptotic value may need times longer than the SNR lifetime.  It might be expected that the evolution will eventually change from a ``pressure-driven snowplow'' ($m=2/7$, McKee and Ostriker \\cite{McKee-Ostr-77}) to a ``momentum-conserving snowplow'' ($m=1/4$, Oort \\cite{Oort-51}), as a consequence that the right side of Eq.~(\\ref{singleeq}) vanishes when $R\\tendsto\\infty$. However, Cioffi et al.\\ (\\cite{Cioffi-et-al-88}) have noted that, even at very late times ($\\sim10^2\\ttran$), the deceleration parameter $m$ is still closer to $2/7$ than to $1/4$. We have shown that such evolutive transition may in fact not occur, because the two kinds of evolution are associated with two different branches of solutions, corresponding to different initial conditions. In other words, for a deceleration parameter smaller than $2/7$, the right side of Eq.~(\\ref{singleeq}) (with $\\gm=5/3$) vanishes more slowly than the left side. Therefore, unless the pressure term is negligible from the beginning, or part of the internal energy of the hot interior is lost by other channels (e.g.\\ by electron conduction, or radiative processes), pressure effects must play an important role in the evolution at any time, until the SNR merges with the ambient medium.  Some of the conclusions we have presented here could have been reached long before. In fact, Blinnikov et al.\\ (\\cite{Blinn-et-al-82}) have reduced the system of equations for the thin-shell evolution to the single equation, equivalent to our Eq.~(\\ref{singleeq}), and have obtained a solution equivalent to our Eq.~(\\ref{slowsol}). However, they have not discussed the properties of this solution, taking that it would have quickly relaxed to the asymptotic behaviour."
    },
    "0402/astro-ph0402284_arXiv.txt": {
        "abstract": "We  analyse the UV spectral energy distribution of a sample of normal galaxies  listed in  the {\\it IUE} INES Guide No. 2-Normal Galaxies (Formiggini \\& Brosch, 2000) using a Principal Component Analysis.  The sample consists of the {\\it IUE}-SW spectra of the central regions of  118 galaxies, where the {\\it IUE} aperture included more than 1 per cent of the galaxy size. The principal components are associated with the main components observed in the UV spectra of galaxies. The first component, accounting for the largest source of diversity, can be associated with the UV  continuum emission. The second component represents the UV contribution  of an underlying evolved  stellar population. The third component is sensitive to the amount of activity in the central regions of galaxies and measures the strength of star formation events.    In all the samples analysed here the principal component representative of star-forming activity accounts for a  significant percentage of the variance. The  fractional contribution to the SED by the evolved stars and by the young population are similar.   Projecting the SEDs onto their eigenspectra, we find that none of the coefficients of the principal components  can outline an internal correlation  or can correlate with the optical morphological types. In a sub-sample of 43 galaxies, consisting of almost only  compact and BCD galaxies, the third principal component defines a sequence related to the degree of starburst activity of the galaxy. ",
        "introduction": "sample}   A suitable set of UV SEDs of galaxies covering a wide range of Hubble types can be extracted from the {\\it IUE} final archive. The data sets provided by INES  ({\\it IUE} Newly Extracted Spectra) consist of low-resolution spectra extracted with an improved method from the line-by-line images of the {\\it IUE} Final Archive, and of high-resolution spectra resampled to the low-resolution wavelength step.  A collection of UV spectra of 274 normal galaxies has been  compiled as the INES Guide No. 2 (Formiggini \\& Brosch 2000). In this guide,  a representative UV spectrum for each galaxy has been selected, combining the longest short-wavelength (SW) and long-wavelength (LW) exposures of the galaxy, both taken trough the large aperture. The aperture position of the {\\it IUE} apertures for the SW and LW spectra, that were obtained independently and required repositioning of the satellite, were checked with  aperture overlays  on the galaxy images. This procedure assured that  the entrance apertures of both SW and LW spectra were centered on  the galaxy optical position or,  alternatively, refer to same physical region of the galaxy, such as the HII regions of NGC4449 and NGC5236.  We found some cases where the {\\it IUE} aperture coordinates did not correspond to the coordinates of the galaxy, and  the {\\it IUE} did not,  in fact, observe the object. Such is, for instance, the {\\it IUE} spectrum of NGC 3077, where the  misplaced  aperture contains  a foreground star instead of the galaxy.   The large  apertures of the {\\it IUE} spectrograph are 10\"$\\times$20\" ovals, each corresponding in area to a  circular diaphragm with a diameter of 15.1 arcsec (Longo \\& Capaccioli 1992), and many of the galaxies observed are extended with respect to these large apertures. In order to estimate the fraction of the galaxy area observed by {\\it IUE}, we calculated for each galaxy the ``coverage parameter'' C, defined  as the logarithmic  ratio between the surface area of the galaxy and the area of the large {\\it IUE} aperture.  \\begin{equation} C=log[\\pi \\times (D_{25}^{2} \\times R_{25}) / (15.1)^{2} \\pi/4] \\end{equation}   Here D$_{25}$ and R$_{25}$ are the major axis and the axial ratio of the optical image of the galaxy as listed in NED, in units of arcsec. The numerator approximates the surface area of the galaxy, represented as an ellipse with the major and minor axes of the galaxy.  A C value of zero implies that the entire galaxy was measured by  the {\\it IUE} spectrum. For faint galaxies, where the axes are not measurable, the C parameter is $\\leq 0$. The histogram of the number of galaxies as a function of the C parameter (see Figure 5 of Formiggini \\& Brosch, 2000) shows that, for 90 per cent of the sample of  274 normal galaxies in the INES Guide No.2, the {\\it IUE} aperture covered less than  10 per cent of the galaxy.     For the present investigation, only the galaxies with a C value up to 2, corresponding to a coverage of more than 1 per cent of the galaxy  were assembled from the sample of Formiggini \\& Brosch (2000). This means that the spectra are dominated by the content of the very center of the light distribution of the galaxy. A few cases  where the aperture position was miscentered were rejected. The total number of galaxies included in the sample analyzed here is 118. Table 1 lists the relevant data for the set. Column two indicates to which sub-sample the specific galaxy belongs, where 1 represents membership in the sub-sample of relatively high S/N and 2 represents membership in the sub-sample of high coverage parameter (see  section 4). This sample, resulting from many different research projects, contains galaxies of all  morphological types, although it is obviously biased toward the brightest UV galaxies. It is a representative sample  of UV galaxy spectra although it is not uniform, and the percentage of galaxies in each morphological class is not representative of the galaxy population at large. Consequently, this sample is not suitable for luminosity function and density investigation. A sub-sample of galaxies with C values up to 1, corresponding to a coverage of more than 10 \\% of the galaxy area, was selected for comparison. This reduced sample  of 43 galaxies, consisting  mostly of BCD and compact objects, is analyzed in section 4.2.  \\begin{figure} \\begin{minipage}{80mm} \\centerline{\\epsfxsize=2.5in\\epsfbox{MD869rv2fig1.eps}} \\caption{The distribution of spectral types} \\end{minipage} \\end{figure}  Figure 1  shows the morphological distribution of the data set. The galaxies have been binned in four morphological bins: ellipticals, spirals, irregulars, and a group containing all the compact objects, namely galaxies classified as BCD, EmLS, HII  and compact. Note that the morphological types  and redshifts  adopted here are from NED, while in the INES Guide  No. 2 all information is from LEDA. The sample includes all classical Hubble types, from elliptical to irregulars, and  a large proportion of compact/BCD/EmLS galaxies whose emission in the UV band is dominated by young, UV bright stars.    \\begin{figure} \\begin{minipage}{80mm} \\centerline{\\epsfxsize=2.5in\\epsfbox{MD869rv2fig2.eps}} \\caption{Histogram of the redshifts distribution} \\end{minipage} \\end{figure}   The redshift distribution  of the sample is shown in Figure 2. {\\it IUE} observed galaxies only in the very nearby Universe, hence the sample analyzed in this paper can be considered as representing the local UV galaxy population. This sample can be used to directly analyze the rest-frame UV SED of galaxies without requiring extrapolation from the optical data.  It is known that the galaxy spectra have a low signal-to-noise ratio in the {\\it IUE} LW region. Actually, a visual inspection of the spectra in the LW range, shows that the noise is the dominant feature. Furthermore, only part of the galaxy sample has been observed in the LW range.  Hence, we have restricted the analysis to the far-UV region covered by the  SW {\\it IUE} images (between 1150--1900 \\AA). This is the wavelength range where the most significant features probing the galaxy evolution occur. ",
        "conclusions": ""
    },
    "0402/astro-ph0402417_arXiv.txt": {
        "abstract": "Microlensing event \\event\\ is characterized by a smooth, slightly asymmetric single-lens curve with an isolated, secure data point that is $\\sim 0.6$ magnitudes brighter than neighboring points separated by a few days.  It was previously suggested that the single deviant data point and global asymmetry were best explained by a planetary companion to the primary lens with mass ratio $q=10^{-3}-10^{-2}$, and parallax effects induced by the motion of the Earth.  We revisit the interpretation of \\event, and show that the data can be explained by wide variety of models.  We find that the deviant data point can be fit by a large number of qualitatively different binary-lens models whose mass ratios range, at the $\\sim 3\\sigma$ level, from $q\\simeq 10^{-4}$ to $\\simeq 10^{-1}$.  This range is consistent with a planet, brown dwarf, or M-dwarf companion for reasonable primary masses of $M\\ga 0.8M_\\odot$.  A subset of these binary-lens fits consist of a family of continuously degenerate models whose mass ratios differ by an order-of-magnitude, but whose light curves differ by $\\la 2\\%$ for the majority of the perturbation.  The deviant data point can also be explained by a binary companion to the {\\it source} with secondary/primary flux ratio of $\\sim 1\\%$.  This model has the added appeal that the global asymmetry is naturally explained by the acceleration of the primary induced by the secondary. The binary-source model yields a measurement of the Einstein ring radius projected onto the source plane of ${\\hat r_{\\rm E}}=1.87\\pm 0.40~{\\rm AU}$.  \\event\\ is an extreme example that illustrates the difficulties and degeneracies inherent in the interpretation of weakly perturbed and/or poorly sampled microlensing light curves. ",
        "introduction": "}  Planetary companions to Galactic disk and bulge microlens stars can be discovered via the short duration perturbation they create to the smooth light curve induced by the parent star \\citep{mp91}.  The majority of these perturbations are expected to be relatively simple and grossly characterized by three observables: the duration, peak time, and magnitude of the perturbation.  In the ideal scenario, these three observables are simply related to the three parameters describing the planetary system: the planet/star mass ratio, the instantaneous projected separation in units of the angular Einstein ring radius, and the angle of the source trajectory relative to the planet/star axis \\citep{gl92}.  Unfortunately, reality is a bit more complicated, and a number of degeneracies have been identified which can hamper the ability to infer these parameters in practice. \\citet{gandg97} demonstrated that there exists an ambiguity in the physical mechanism that sets the width of low-amplitude perturbations which can result in an order-of-magnitude uncertainty in the inferred mass ratio.  \\citet{gaudi98} pointed out that a subset of binary sources with extreme flux ratios can reproduce the duration and magnitude of a subset of planetary microlensing perturbations, although \\citet{han02} demonstrated that this degeneracy could be resolved with astrometric observations during the perturbation. \\citet{gs98} discuss a two-fold discrete degeneracy in the projected separation prevalent in high-magnification planetary events.  Along with these anticipated degeneracies, a few have been uncovered in the process of detailed modeling of observed events.  \\citet{bennett99} invoked a planetary companion to explain a short-duration deviation seen on a close binary-lens light curve.  However, it was later shown by \\citet{albrow00} that this perturbation could also be fit by one of the secondary caustics of the close binary lens, when rotation of the binary is considered. \\citet{gaudi02} found a weakly asymmetric event that could be equally well-explained by a planetary companion, or parallax deviations arising from the motion of the Earth.  Many of these degeneracies are `accidental,' in the sense that they arise from chance similarities between deviations caused by different physical situations, rather than by intrinsic degeneracies in the lens equation itself. They are therefore generally only approximate degeneracies, and can be resolved with accurate, well-sampled light curves.  Given the short duration and unpredictability of planetary deviations, dense, continuous and accurate light curve coverage is necessary to both detect and accurately characterize planetary microlensing perturbations.  A lensing star with a planetary companion is just an extreme limit of a binary-lens.  As discussed by numerous authors, binary lenses are themselves subject to numerous degeneracies \\citep{md95,albrow99,jm01}, some of which are rooted in symmetries in the lens equation itself \\citep{dominik99}, and are therefore nearly perfect \\citep{afonso00,albrow02}.  Binary lenses in which the source does not cross any caustics can also be confused with binary sources. This is especially problematic when only single-band photometry is available.  This may partially account for the fact that, although they were predicted to be plentiful \\citep{gh92}, only one candidate binary-source lensing event has been identified\\footnote{See \\citet{dominik98} and \\citet{hj98} for additional discussions of the apparent lack of binary-source events.}, event OGLE-2003-BLG-095 \\citep{collinge04}.  Indeed, \\citet{collinge04} found that the binary-source model for OGLE-2003-BLG-095 is only preferred over a binary-lens model at the $\\sim 3\\sigma$ level.  Source and lens binarity are not the only regimes where degeneracies are plentiful; global deviations from the fiducial point-source, point-lens, uniform motion (i.e.\\ \\citealt{pac86}) light curve have also been found to be subject to degeneracies.  Such deviations come in a variety of forms.  The motion of the Earth produces departures from uniform relative motion which can induce observable deviations from the standard lightcurve form.  These deviations can be quite dramatic for events with timescales $\\te$ of order or larger than a year \\citep{smith02}. However, in the more usual case where $\\te\\ll {\\rm yr}$, the effect of the motion of the Earth can be approximated by a constant acceleration, which results in deviations that can be symmetric or asymmetric with respect to the peak of the event \\citep{gmh94,smp03}, although asymmetric deviations are generally easier to recognize.  Unfortunately, as demonstrated by \\citet{smp03}, any such weak parallax deviation can also be explained by acceleration of the source, thus making the parallax interpretation non-unique for such short timescale events. \\citet{smp03} furthermore demonstrated that constant-acceleration events are subject to a two-fold discrete degeneracy between the magnitude and direction of the acceleration and the event timescale.  \\citet{gould03} showed that, in fact, light curve degeneracies extend to even higher order, and are present when one takes into account not only the acceleration, but also the jerk.  This jerk-parallax degeneracy has been seen in one event toward the bulge \\citep{park04}, and has been invoked to resolve the discrepancy between the photometric and microlensing mass determinations of the microlens in the Large Magellanic Cloud event MACHO-LMC-5 \\citep{gould03}.  Galactic bulge microlensing event \\event\\ exhibits both a global asymmetry and single deviant data point.  \\citet{jp02} first considered the interpretation of this event; they fit binary-lens models and included non-uniform motion caused by either parallax or arbitrary uniform acceleration.  They argued that, although both uniform acceleration and parallax explain the global deviations equally well, the event was most naturally explained by parallax-induced deviations, and that the single deviant point was best fit by a binary-lens with mass ratio $10^{-3}-10^{-2}$.  This mass ratio implies a companion in the Jupiter mass range for likely primary masses, and they therefore concluded that \\event\\ was a possible planet candidate.  Here we revisit the interpretation of this event.  We demonstrate that, in fact, there are many possible interpretations of \\event.  We find that the short-timescale perturbation could arise from a stellar companion to the source, or a stellar, brown-dwarf, or planetary companion to the lens.  The global asymmetry could arise from parallax deviations, or from acceleration of the source induced by its companion.  All of these interpretations are indistinguishable at the $\\sim 3\\sigma$ level.  Therefore, the correct interpretation of this event is unclear, and it serves as a particularly extreme reminder of the degeneracies involved with poorly sampled and weakly-perturbed microlensing light curves.   \\begin{figure} \\epsscale{1.1} \\plotone{f1.eps} \\caption{\\label{fig:one} The points with errorbars show the light curve of \\event. The long-dashed magenta line shows the best single-lens, single-source, constant-velocity model fit to the data, with the single high point and two neighboring points removed.  The short-dashed blue line shows the best single-lens, single-source, constant-acceleration model fit to the same dataset.  The dotted green line shows the best single-lens, binary-source, constant-acceleration model to the entire data set. The solid red line shows the best single-source, binary-lens, constant-acceleration model to the entire dataset.  The inset shows a close-up near the deviation. }\\end{figure} ",
        "conclusions": "Microlensing event \\event\\ exhibits a slightly asymmetric, smooth light curve with a single data point that deviates by $\\sim 0.6$ magnitudes from neighboring points separated by several days. \\citet{jp02} concluded that the simplest interpretation of \\event\\ was an event with parallax deviations arising from the motion of the Earth, and a short-timescale deviation due to a binary lens with mass ratio $q=10^{-3}-10^{-2}$, thereby making \\event\\ a candidate planetary event.  Here we demonstrated that \\event\\ can be reasonably well-fit by several different classes of models, and a wide range of parameters within each model class.  We found that the data can be fit by many different binary-lens models whose mass ratios span three orders of magnitude, from $q=10^{-4}$ to $10^{-1}$, thereby making the secondary consistent with a planet, brown dwarf, or M-dwarf for reasonable primary masses.  A subset of these binary-lens fits form a family of continuously degenerate models, whose mass ratios differ by an order of magnitude.  Astonishingly, the light curves of these models differ by $\\la 2\\%$ for the majority of their duration.  A binary-source model is also consistent with the data, for a secondary/primary flux ratio of $\\sim 1\\%$.  This model also naturally explains the global asymmetry of the lightcurve as due to the acceleration of the primary induced by the secondary.  Under the assumption of a bound binary-source, this model yields an estimate of the Einstein ring radius projected on the source plane of ${\\hat r_{\\rm E}}=1.87\\pm 0.40~{ \\rm AU}$.  All of these fits differ by $\\la 3\\sigma$, and are essentially indistinguishable when the scatter due to likely source variability is considered.  Unfortunately, the lack of color information during the event precludes the discrimination of models based on the positions of the source and blend on a color-magnitude diagram.  Although the primary goal of the study by \\citet{jp02} was to affect a modification of the OGLE observation strategy to ensure good sampling of short-duration perturbations, rather than argue that \\event\\ was a bona fide planetary event, it is still somewhat disturbing that many binary-lens fits were missed, and the possibility of a binary-source interpretation was not discussed at all.  \\event\\ serves as an extreme reminder of the degeneracies inherent in microlensing events, and highlights the difficulties in interpreting poorly-sampled and weakly-perturbed events. These difficulties become especially important when attempting to detect planets with microlensing.  Here observers and modelers need to be especially vigilant: in order to produce convincing and reliable planetary detections, it is essential not only to achieve dense and accurate photometry of planetary perturbations, but also to acquire as much auxiliary information as possible, and to perform detailed, careful, and thorough modeling, in order to ensure that planetary detections are robust."
    },
    "0402/astro-ph0402621_arXiv.txt": {
        "abstract": "{ An XMM-Newton target of opportunity observation of the field around the transient 18.37 s pulsar \\xte\\ in the Small Magellanic Cloud (SMC) revealed two bright, long-period X-ray pulsars in the EPIC data. A new pulsar, \\xnorth, with a pulse period of 701.7 $\\pm$ 0.8 s was discovered and 500.0 $\\pm$ 0.2 s pulsations were detected from \\xsouth\\ (= CXOU\\,J005455.6$-$724510), confirming the period found in Chandra data. We derive X-ray positions of RA = 00$^{\\rm h}$54$^{\\rm m}$55\\fs88, Dec = --72\\degr45\\arcmin10\\farcs5 and RA = 00$^{\\rm h}$55$^{\\rm m}$18\\fs44, Dec = --72\\degr38\\arcmin51\\farcs8 (J2000.0) with an uncertainty of 0\\farcs2 utilizing optical identification with OGLE stars. For both objects, the optical brightness and colours and the X-ray spectra are consistent with Be/X-ray binary systems in the SMC. ",
        "introduction": "The Small Magellanic Cloud harbours a large number of high mass X-ray binary (HMXB) systems, more than are known in the Large Magellanic Cloud and the Milky Way \\citep[see compilations by ][]{2000A&A...359..573H,2003PASJ...55..161Y} despite the much smaller mass of the SMC. The catalogue of \\citet{2004A&A...414..667H} comprises 65 HMXBs and candidates in the SMC with at least 37 showing X-ray pulsations which indicate the spin period of the neutron star in orbit around a high mass early type star. The number of X-ray pulsars in the SMC has meanwhile further grown by five: For three of the known candidate HMXBs pulsations of 202 s, 500 s and 138 s were detected, an X-ray source known from ROSAT was discovered as 34.1 s pulsar and 18.37 s pulsations were found in RXTE observations of the SMC. In addition accurate Chandra positions enabled the location of two RXTE-discovered pulsars with 82.4 s and 7.78 s period, the latter identified with SMC\\,X-3 \\citep[see][ and references therein]{2004astroph0402053C}.  The non-imaging instruments on RXTE allowed the 18.37 s transient pulsar \\xte\\ to be located to an accuracy of only 0.1\\degr$\\times$0.06\\degr. Hence, we proposed an XMM-Newton target of opportunity observation which was performed on December 18, 2003. We could not detect 18.37 s pulsations from any of the weaker sources seen in the EPIC images due to insufficient counting statistics and were therefore not able to identify the target. However, the two brightest X-ray sources were found to exhibit pulsations with longer periods \\citep{2004ATel..219....1H}. \\xnorth\\ (hereafter \\xna) is a new pulsar and \\xsouth\\ (hereafter \\xsa) is identified with CXOU\\,J005455.6$-$724510 by position and X-ray period of $\\sim$500 s. Here we present the results of a temporal and spectral analysis of the X-ray data of these two SMC pulsars and propose optical counterparts with properties as expected for Be/X-ray binary systems. ",
        "conclusions": "The December 2000 XMM-Newton observation of the SMC region around \\xte\\ revealed two X-ray pulsars with long pulse periods. The 500 s pulsar \\xsa\\ was the brighter of the two and was detected by Chandra on July 4, 2002, with pulsations reported by \\citet{2004ATel..217....1E}. It was detected during three ROSAT observations in May 1993, April 1994 and April/May 1997. Using the parameters derived from the EPIC spectra and the ROSAT PSPC and HRI spectral response to estimate expected ROSAT count rates, shows that this pulsar exhibits variations in flux by about a factor of 3 with intensities lowest in May 1993 and highest in December 2000. During several other ROSAT observations the source was not detected. However, the detection threshold did not reach the low flux level at which the source was detected in 1993 (large off-axis angles, short exposures) and no additional information about variability can be inferred. The source is probably also identical to the ASCA source AX\\,J0054.8$-$7244 \\citep{2003PASJ...55..161Y} which was detected in Nov. 1998 at a flux level of 4.9\\ergcm{-13}, similar to the low intensity state in May 1993. \\xna\\ is a newly discovered X-ray pulsar in the SMC with a period of 701 s. While it was the fainter of the two pulsars during the XMM-observation, it was the slightly brighter ($\\sim$20\\%) one during the 1993 PSPC observation, the only time when it was detected by ROSAT. Assuming the EPIC spectral parameters shows that it was brighter by 40\\% during December 2000 compared to May 1993. Eight non-detections by ROSAT and also by ASCA suggest that \\xna\\ is overall fainter than \\xsa, falling below the detection thresholds of the ROSAT and ASCA observations (which were, however, close or even above the detection flux) most of the time. Both pulsars seem to belong to the class of Be/X-ray binaries with long pulse period and moderate intensity variations.  Spectral analysis of both pulsars shows that a power-law with photo-electric absorption does not reproduce the EPIC spectra properly. Including soft emission components or an exponential cut-off yields acceptable fits, but the statistical quality does not allow us to decide between these models. The major difference between the cut-off model and the two component models is the resulting absorption column density, which is consistent with the Galactic foreground value of 6\\hcm{20} for the cut-off model. However, the reddening inferred from the suggested optical counterparts is incompatible with such low absorption and favours the soft component models. On the other hand the absorption is very high for the thermal plasma model which leads to an implausibly high intrinsic source luminosity in particular for \\xna. A large fraction of the column density may be local to the source and the different emission components might suffer different amounts of absorption. This could indicate that the spectrum at lower X-ray energies is more complex as it is seen from other HMXBs in the Magellanic Clouds. E.g. the pulse phase averaged spectrum of EXO\\,053109-6609.2 \\citep{2003A&A...406..471H} shows power-law components attenuated by different column densities. A similar behaviour of the two pulsars presented here is expected if the absorption changes with pulse phase as it is indicated by the pulse profiles.  With the new discovery of the 701 s pulsar \\xna\\ the number of pulsars in the SMC has grown to 45 \\citep{2004astroph0402053C} and future X-ray observations of the SMC promise to find more."
    },
    "0402/astro-ph0402441_arXiv.txt": {
        "abstract": "The observed cooling rate of hot gas in clusters is much lower than that inferred from the gas density profiles. This suggests that the gas is being heated by some source. We use an adaptive-mesh refinement code ({\\sc Flash}) to simulate the effect of multiple, randomly positioned, injections of thermal energy within 50~kpc of the centre of an initially isothermal cluster with mass $M_{200}=3\\times10^{14}\\Msol$ and $kT=3.1$~keV.  We have performed eight simulations with spherical bubbles of energy generated every $10^8$ years, over a total of $1.5\\Gyr$.  Each bubble is created by injecting thermal energy steadily for $10^7$ years; the total energy of each bubble ranges from 0.1--3$\\times10^{60}\\erg$, depending on the simulation.  We find that 2$\\times10^{60}\\erg$ per bubble (corresponding to a average power of $6.3\\times10^{44}\\ergs$) effectively balances energy loss in the cluster and prevents the accumulation of gas below $kT=1$~keV from exceeding the observational limits.  This injection rate is comparable to the radiated luminosity of the cluster, and the required energy and periodic timescale of events are consistent with observations of bubbles produced by central AGN in clusters.  The effectiveness of this process depends primarily on the total amount of injected energy and the initial location of the bubbles, but is relatively insensitive to the exact duty cycle of events. \\endabstract  \\keywords galaxies:clusters:general--cooling flows--X-ray:galaxies:clusters \\endkeywords ",
        "introduction": "Gas cooling at the centre of a cluster halo is an inherently unstable process: cooling increases the gas density, which in turn enhances the cooling rate. X-ray observations show that the cooling time of gas in most cluster cores is less than the Hubble time \\citep[e.g.][]{1977MNRAS.180..479F}. Unless cooling is balanced by some form of heating, gas will flow into the cluster centre at rates up to $\\sim 1000\\Msolyr$ \\citep[e.g.][]{2001A&A...365L.104P}. What happens to the cooled gas is unclear; it could be consumed by star formation, or lead to a reservoir of low temperature ($<1\\keV$) material in the core. Although this may be the fate of a fraction of the gas, there are indications that most gas in fact does not follow either route.  First, the star formation rates in the central galaxies of clusters are much lower than the inferred mass inflow rates \\citep{1989AJ.....98..180O,1987MNRAS.224...75J}, rarely approaching $\\sim 100 \\Msolyr$. To estimate the star formation rate in a typical cluster, we can average over the large sample of clusters studied by \\citet{1999MNRAS.306..857C}. This suggests that the typical star formation rate is less than $10\\Msolyr$.  Secondly, we can compare the mass deposition rates with observations of the molecular gas content of clusters.  \\citet{2001MNRAS.328..762E} finds molecular gas masses ranging from $10^9$ to $2\\times 10^{11}\\Msol$, with an average of $2.6\\times10^{10}\\Msol$. Assuming a gas consumption timescale of $10^9$ years \\citep[see][]{2001MNRAS.328..762E}, this implies a deposition rate that may be as high as $200\\Msolyr$ in a few clusters, but is $\\sim 30\\Msolyr$ on average.  Similar limits are obtained by \\citet{2003A&A...412..657S}. Finally, recent spectroscopic X-ray observations show no evidence for significant gas cooling below $1\\keV$ \\citep{2001A&A...365L..99K,2001A&A...365L.104P,2001A&A...365L..87T}, and observations of molecular and neutral material reveal that the amount of cold gas in clusters of galaxies is also much less than expected from the integrated cooling flow rate \\citep[typically less than $30\\Msolyr$ -- ][]{2002MNRAS.337...49E,2003ApJ...594L..13E,2003A&A...412..657S}.  This cooling-flow paradox has led many authors to investigate mechanisms to quench gas cooling. Observations of merging clusters show little evidence for cooling flows, which suggests that the merger process might be implicated in disrupting cooling.  Recent simulations have shown that sub-halo merging can indeed heat up gas \\citep{2003astro.ph..9836B}; however, the amount of cold gas produced in these simulations is still too large compared to that observed, leading to the conclusion that additional heating processes must be involved.  Several alternatives have been proposed, including energy injection from radio sources or active galactic nuclei \\citep[AGN;][]{1995MNRAS.276..663B,2001MNRAS.328.1091Q,2002MNRAS.332..729C}, viscous dissipation of sound waves \\citep{2003MNRAS.344L..43F,2003astro.ph.10760R} and thermal conduction \\citep{2003astro.ph..8352V,2004astro.ph..1470D}.  Each of these has advantages and disadvantages.  In particular, it is difficult to balance cooling, with its $\\rho^2$ density dependence, with heating processes that typically scale as $\\rho$.  Since cooling and heating can then balance only at one particular density, some of these feedback mechanisms may lead to unstable solutions with some regions of the cluster being efficiently heated while others continue to cool catastrophically.  AGN are particularly promising candidates for balancing cooling, given their potentially large energy reservoir \\citep{1993MNRAS.263..323T,1995MNRAS.276..663B,2001MNRAS.325..497B}. In most numerical simulations of this process to date, energy injection produces bubbles at high temperature and low density that, after a short expansion phase, gain momentum by buoyancy \\citep{2001MNRAS.325..676B,2001MNRAS.328.1091Q,2003MNRAS.339..353B}. This mimics the effect of a jet which rapidly loses its collimation, as often observed in local clusters \\citep{2003astro.ph.10011E}.  In other simulations, gas is injected at high velocity to mimic jets which retain their large-scale coherence \\citep{2002MNRAS.332..271R,2003astro.ph..7471O}.  In this paper, we consider the first mechanism and show that such bursts of localised energy can induce convection of the intra-cluster medium (ICM), which leads to a quasi-stable cluster configuration and reduces the average mass deposition rate to within observational limits. We discuss whether the required energy and duty cycle of AGN activity are compatible with observational limits.  This paper is organised as follows. In Section~\\ref{sec:simul} we describe the setup of our simulations and discuss the parameters explored. Results are presented in Section~\\ref{sec:results} and, in Section~\\ref{sec:discuss}, we compare the required energy and duty cycle with observational constraints. Finally, our conclusions are summarised in Section~\\ref{sec:conc}.  \\begin{table} \\begin{center} \\begin{tabular}{ccccc} \\hline & & \\multicolumn{2}{c}{\\sc Injected/Radiated Power} \\\\ name  & $E$           & $\\dot E_i$      & $<\\dot E_r>_{100}$ \\\\ & $10^{60}\\erg$ & $10^{44}\\ergs$  & $10^{44}\\ergs$ \\\\ \\hline A     & 0             & 0               & no cooling \\\\ \\hline S0.0  & 0             & 0               & 28 \\\\ \\hline S0.1  & 0.1           & 0.32            & 18 \\\\ \\hline S0.3  & 0.3           & 0.95            & 18 \\\\ \\hline S0.6  & 0.6           & 1.9             & 13 \\\\ \\hline S1.0  & 1.0           & 3.2             & 13 \\\\ \\hline S1.5  & 1.5           & 4.7             & 8.7 \\\\ \\hline S2.0  & 2.0           & 6.3             & 7.9 \\\\ \\hline S3.0  & 3.0           & 9.5             & 6.8 \\\\ \\hline \\end{tabular} \\end{center} \\caption{Summary of the nine simulations performed. The different simulations are referred to as S$n$, where the energy $E$ of a single bubble is $n\\times 10^{60}$ erg. $\\dot E_i$ is the mean energy injection rate over a duty cycle, and $<\\dot E_r>_{100}$ is the mean emitted energy rate, averaged over the last $10^8\\yr$ of the simulation.} \\label{tbl:inject} \\end{table}  \\begin{figure*} \\includegraphics[width=41pc]{energy.ps} \\caption{The evolution of the total energy of each simulation is shown by plotting $\\Delta E(t)=E_{\\rm T}(t) - E_{\\rm T}(0)$, where $E_{\\rm T}(t)$ is the sum of internal, kinetic and potential energy at time $t$ and $E_T(0)$ is its initial value.  The saw-tooth shape of the curves results from the discrete AGN events and subsequent cooling. At the mean injection rate of simulation S2.0, the energy keeps an almost constant value within the simulation time. We tested that this behaviour is maintained up to $5\\Gyr$, though we show the evolution only to $3\\Gyr$ for clarity.} \\label{fig:energy} \\end{figure*} ",
        "conclusions": "  \\begin{itemize} \\item[1.] For a time averaged energy injection rate of $6\\times 10^{44}\\ergs$, the mass inflow rate is less than $30\\Msolyr$, compatible with available observational limits.  With this amount of heat input, the total energy of the cluster remains approximately constant over 5 Gyr. \\item[2.] The evolution of the total cluster energy depends primarily on the total amount of energy injected and the spatial distribution of bubbles, but is only weakly sensitive to the duty cycle of heating events or to whether the bubbles are produced singly or in pairs. \\item[3.] The bubble activity generates concentric sound waves that are clearly evident in unsharp--masked projections of cluster emissivity. \\item[4.] When the injected energy just balances cooling, the entropy and temperature profile of the cluster remain approximately unchanged from their initial configurations. \\end{itemize}  In summary, periodic energetic events of the kind we have simulated can reduce the mass flow rate and accumulation of cold gas in massive clusters to within observational limits.  However, this mechanism operating on a fully formed cluster does not result in a final luminosity consistent with observations.  It is likely that the structure of the progenitors from which the cluster formed was affected by heating events prior to the assembly of the final cluster."
    },
    "0402/astro-ph0402394_arXiv.txt": {
        "abstract": "{We discuss a model of X-ray variability of active galactic nuclei (AGN). We consider multiple spots which originate on the surface of an accretion disk following intense irradiation by coronal flares. The spots move with the disk around the central black hole and eventually decay while new spots continuously emerge. We construct time sequences of the spectra of the spotted disk and compute the corresponding energy-dependent fractional variability amplitude. We explore the dependence on the disk inclination and other model parameters. AGN seen at higher inclination with respect to the observer, such as Seyfert~2 galaxies, are expected to have fractional variability amplitude of the direct emission by a factor of a few higher than objects seen face on, such as the Seyfert~1s. ",
        "introduction": "Broad band spectra of active galactic nuclei show the presence of several components. The two principle contributions are: (i)~the Big Blue Bump extending from optical/UV band to (sometimes) soft X-ray band, and (ii)~the hard X-ray power law. The Big Blue Bump emission is conventionally interpreted as originating from an accretion disk (Czerny \\& Elvis 1987; Koratkar \\& Blaes 1999; however, see e.g.\\ Collin et al.\\ 2002 for discussion of problems encountered by this scheme). On the other hand, the nature and geometry of the region responsible for the hard X-ray emission is still under discussion. Several models were proposed, the most popular ideas being a hot extended corona overlaying a relatively cold disk (e.g. Czerny \\& Elvis 1987, Haardt \\& Maraschi 1991, \\Agata \\& Czerny 2000, Liu et al. 2002, Merloni 2003), an inner hot flow (e.g. Ichimaru 1977, Narayan \\& Yi 1994, Narayan et al. 2002) the lamp-post model (where a point-like X-ray source is located at a specified height above the disk, e.g. Henri \\& Pelletier 1991, Malzac et al. 1998), and the model of multiple hot flares produced via magnetic field reconnections (Galeev et al. 1979 and later papers). The field was reviewed, for example, by Leighly (1999), Collin (2001), and  Poutanen (1998) and Done et al. (2002) (in the context of X-ray binaries). Notice that these scenarios are not completely disparate and there may be a certain overlap between them. In all models the primary source of X-rays should be strongly variable.  Indeed, AGN are variable in the X-ray band (e.g.\\ Lawrence et al.\\ 1987; Taylor et al.\\ 1993) but the origin of variability cannot be examined directly, because the relevant central regions still remain unresolved to direct imaging. However, studies of the X-ray spectral variability offer a direct insight into the structure of accretion flows onto a central black hole, which is assumed to power AGN and determine their spectra.  In the flare model, sudden dissipation occurs in very localized regions above the disk surface (coronal loops, in analogy with the solar corona). In this case the local irradiation flux can be orders of magnitude higher then the steady level of the disk flux itself, but the irradiation does not last very long and only a small fraction of the disk surface is irradiated at every moment.  Such arrangement of X-ray emitting region has been suggested for the first time by Galeev et al.\\ (1979), and it was subsequently developed in many papers (e.g.\\ Abramowicz et al.\\ 1991; Haardt et al.\\ 1994; van Oss et al.\\ 1993; Poutanen \\& Fabian 1999; \\.Zycki 2002). The development of a flare leads to a burst of `primary emission' as well as to the formation of a hot spot underlying the flare where roughly half of the X-ray flux is reprocessed by the disk. Irradiation of the disk surface in hydrostatic equilibrium leads to the formation of strongly stratified medium -- a hot fully ionized skin covering a cooler, more neutral zone -- and so the X-ray spectrum resulting from reprocessing should contain signatures that are characteristic for multi-temperature gas (Nayakshin et al.\\ 2000; Ballantyne et al.\\ 2001; R\\'o\\.za\\'nska et al.\\ 2002). The irradiating flux is locally very large, exceeding considerably the stationary energy flux which is dissipated inside the disk (e.g.\\ Nayakshin 2000; Ballantyne et al.\\ 2001). The resulting ionized skin is relatively optically thick (Collin et al.\\ 2003).  The flare/spot model is a possible (although not unique) explanation of narrow emission features, which have been reported in $\\sim5$--$6$~keV X-ray spectra of several AGNs and interpreted in terms of in terms of localized iron-line emission (Turner et al. 2002, 2004; Guainazzi 2003; Yaqoob et al. 2003; Dovciak et al. 2004).  In the present paper we test the flare/spot model by analyzing its predictions for the fractional variability amplitude in the X-ray band. We model the local spot/flare spectrum as a sum of flare (primary) emission of a power law shape, and a spot (reflected) emission. The spot emission is determined by using the coupled {\\sc titan/noar} codes (Dumont et al.\\ 2000) to solve the radiative transfer and the code of R\\'o\\.za\\'nska et al.\\ (1999) to calculate the hydrostatic equilibrium. We assume random distribution of spots and flares across the disk and we account for their motion during the time-integrated observation. We consider both non-rotating (Schwarzschild) and rapidly rotating (Kerr) black holes, and we apply general relativity corrections using the {\\sc ky}  code of Dov\\v{c}iak et al.\\ (2003).  In our approach, a sequence of solutions is specified by fixing statistical properties which the flare/spot distribution is required to obey. Given a particular solution, the corresponding energy-dependent fractional variability amplitude $F_{\\rm{}var}$ is calculated. The model is described in Section~\\ref{sect:model}. Results are given in Section~\\ref{sect:results}. The relevance of the model for explaining spectrum and variability of a typical Seyfert~1 galaxy is discussed in Section~\\ref{sect:discussion}. ",
        "conclusions": "\\label{sect:discussion}  The flare/spot model is an attractive explanation of the X-ray spectra and variability of AGN. In this scenario X-ray emission is generated both in hot magnetic loops above an accretion disks and in the bright spots created under the loops by strong irradiation. In the present paper we tested this model by analyzing the predicted fractional variability amplitude, $F_{\\rm{}var}$.  We derive simple analytical formulae which allow to estimate the level of variability from the assumed mean number of flares, flare duration, integration time of a single observation, the ratio of inner to outer disk radius and the radial dependence of the flare luminosity. This energy-independent and inclination-independent formula roughly applies to the case of a non-rotating black hole and small size of the spots. The formula is based on assumption of the uniformly covered disk surface but it can be readily generalized to another kind of distribution.  Larger spots and/or fast rotating black hole, with inner disk radius close to the marginally stable orbit require numerical approach, and we show both the dependence of the resulting  $F_{\\rm{}var}$ on energy and inclination angle.  This leads to a firm prediction of our model which can be used to test the basic scenario. The model shows that if the disk extends close enough to the black hole, the general relativity effects lead to significant enhancement of the variability at large inclination angle of observation. Therefore, Seyfert~2 galaxies, if intrinsically identical to Seyfert~1 objects but viewed at a larger angle, should exhibit statistically higher $F_{\\rm{}var}$ when measured at the same energy. A factor of 3 difference between Seyfert~2 and Seyfert~1 galaxies is expected for a Kerr black hole and a factor of 1.4 for a Schwarzschild black hole.  The observational evidence of any trends in this direction is scarce at present, but it may support or reject our view more firmly in future. However, some values can be given already now. Mean variability level of Seyfert 1 objects is $\\sim 18$\\% on the short timescales of 1 day in the sample of Markowitz et al. (2003). Seyfert 2 galaxy NGC 4945 is seen through the torus and it displays variability at the level of $\\sim 40$\\% in the hard X-ray band. As another example -- NGC 7582 (Mihara et al.\\ 2000), has revealed the normalized variability amplitude in the hard X-ray band to be about $\\sim 30$\\%. Studies of more Seyfert 2 objects are clearly needed.  The overall intra-day variability level of Seyfert galaxies is well explained by the flare model if the mean number of flares is of the order of 30--100 for assumed non-rotating black hole, or 300--1000 for a fast rotating black hole. The number of flares requested is much larger than the usual expectation of $\\sim 10$ flares. Such a small number of flares is predicted when an assumption is made that all flares have the same luminosity. In our more realistic model we allow for the flare luminosity to depend on the occurrence radius. When we assume that the flare luminosity scales with radius as $\\propto r^{-3}$, most of the source X-ray luminosity comes from a few flares generated in the innermost part of the disk which enhances the variability. This trend is clearly seen from our analytical expression \\ref{eq:Nvar} for the normalized variance, which, for $\\beta_{rad} = 3$, and $R_{\\rm out}>> R_{\\rm in}$, reduces to Eq.~\\ref{eq:limit}. Therefore, the total number of flares, $N_{\\rm{}mean}$ can still be quite high for a source with a moderate variance.   The energy dependence of $F_{\\rm{}var}$ is generally weaker in the model than in data. Observed variations show trends with energy in 1--10~keV band. For example, in Markowitz et al. (2003) Akn 564 varies at the level 18--22\\%, IC 4329A at the level 11--14\\% and MCG--6-30-15 at 14--21\\%, depending on the considered energy. In our models the trends in $F_{\\rm{}var}$ with energy never exceed 2\\%.  This discrepancy can be possibly solved by relaxing several simplifications: \\begin{itemize}  \\item the local spectrum was computed in detail only at a single radius. In reality, the shape of the spot spectrum is expected to show significant trends with the disk radius, as for example emphasized by \\. Zycki \\& R\\'o\\.za\\'nska (2001). A grid of spot spectra should be computed (parameters space of the model is rather rich, and so the task is numerically very time consuming and we postpone it for future work);  \\item hydrostatic equilibrium was assumed to compute the irradiated disk structure. However, timescales of the flares and of restoring the hydrostatic equilibrium are roughly comparable (see the discussion by Nayakshin \\& Kazanas 2002 and Collin et al.\\ 2003). Therefore, neither the assumption of hydrostatic equilibrium nor the assumption of unperturbed disk are satisfactory. Actual disk evolution should be followed;  \\item we assumed a unique value of flare duration since we aimed at modeling the variations at the dominant timescale, e.g.\\ at the knee of the power spectrum. A distribution of flare timescales, a coupling between the flare occurrence (avalanches), and an exact profile of an individual flare are needed if we want the model to reproduce the entire power spectrum (e.g.\\ Lehto 1989, Abramowicz et al.\\ 1991; Xiong et al.\\ 2000, Merloni \\& Fabian 2001);  \\item we neglected the possible effect of the variable warm absorber which may be important in the soft X-ray band for some sources, as argued by Inoue \\& Matsumoto (2003);  \\item we neglected the possible contribution of the radiation reprocessed by some distant reflector, like an outer disk or dusty/molecular torus (e.g.\\ Krolik et al.\\ 1994).  \\end{itemize}  It is to be seen whether elimination of these assumptions would lead to better agreement of the predicted energy dependence of $F_{\\rm{}var}$ with the data. Particularly difficult seems to be the explanation of both apparently lower observed variability in the iron line region as well as the enhancement of the variability towards low energies seen in Fig.~5 of Markowitz et al. (2003).  We considered just a few special cases along this line.  A change of the life time of a flare from radius-independent to scaled with Keplerian timescale, without a change of other parameters, resulted in a fractional variability amplitude even less dependent on the energy than previously. It is simply caused by the fact that in the case of such a scaling relatively more energy is dissipated in outer region, so the most relativistically broadened and variable inner region contributes less to the total lightcurve. The overall normalization of the $F_{\\rm var}$ depends on the proportionality constant between the life times and the Keplerian timescale (with other parameters fixed).  If we assume that the overall radial dependence of the dissipation should not be modified, we can consider two representative examples.  In the first case we assume that the life time of a flare scales with the Keplerian timescale but the probability of a flare to appear at a given radius scales inversely with the local Keplerian timescale. We have therefore less long-living flares in the outer region and more short-living flares in the inner region, with scaling of a single flare luminosity with radius unchanged. In this case  again the fractional variability amplitude is less dependent on the energy than in our basic model. Having more flares localized in the inner region lead to reduction in the fluctuations in the relativistically smeared red wing.  In the second case we again assume that the life time of a flare scales with the Keplerian timescale, we still adopt the uniform distribution of the flares across the disk surface but this time we assume that a flare luminosity decreases with radius even more strongly ($\\beta_{\\rm rad} = 4.5$ instead of usually adopted $\\beta_{\\rm rad} = 3$) to compensate for an increase of the flare life time. Such a solution leads to slightly enhanced dependence of the fractional variability amplitude on the energy but the effect is not strong.  Complex dependence of the spectral shape of the reflected component on the disk radius may introduce significant modification to the predicted energy dependence. Clear suggestion of what is needed can be seen from the recent analysis of MCG--6-30-15  by Vaughan \\& Fabian (2003). A constant component with complex energy dependence is apparently needed in order to formally model the fractional variability amplitude in this source (see their figs. 11 and 16, top panel). This component may perhaps be understood as a reflection component (above $\\sim 1 $ keV) and a contribution from emission/scattering by some extended medium (below $ \\sim 1 $ keV). However, it is not clear whether there is any possibility to find a parameter range which would satisfy two 'opposing' trends seen in these data: we need strong reflection from distant disk region in order to reproduce the constant component but we need strong reflection from innermost region in order to explain the strongly relativistically broadened iron line profile also seen in these data. Attempts by \\Agata \\& \\. Zycki (2001) and Ballantyne et al. (2003) were not successful.   Flares are not the only possibility to explain the X-ray emission of AGN. Other scenarios include lamp-post (standing shock) model (e.g. Henri \\& Pelletier 1991, Malzac et al. 1998), model of gradual or rapid disk evaporation and its replacement by the hot flow (e.g. Narayan \\& Yi 1994, Liu et al. 2002), possibly with an outflow (e.g. Blandford \\& Begelman 1999), and the cloud model (e.g. Collin et al. 1996, Karas et al. 2000). Further work, taking into account variability issues, is needed to determine whether the flare/spot model is the most satisfactory."
    },
    "0402/astro-ph0402677_arXiv.txt": {
        "abstract": "s{ Evidence is mounting that some Ultra-luminous X-ray sources (ULXs) may contain accreting intermediate-mass black holes (IMBHs). We review the current observational evidence for IMBH-ULXs. While low-luminosity ULXs with L$_X \\lapprox$ 10$^{39.5}$ erg~s$^{-1}$ (assuming isotropic emission) are consistent with mildly X-ray beamed high-mass X-ray binaries, there are a considerable number of ULXs with larger X-ray luminosities that are not easily explained by these models.  Recent high-S/N XMM X-ray spectra are showing an increasing number of ULXs with ``cool disks'' -- accretion disks with multi-color blackbody inner disk temperatures kT$_{in} \\sim$ 0.1$-$0.2 keV, consistent with accreting IMBHs. Optical emission-line studies of ULX nebulae provide useful measurements of X-ray energetics, and can thus determine if the X-rays are emitted isotropically. Analysis of an optical spectrum of the Ho~II ULX nebulae implies an X-ray energy source with $\\sim$10$^{40}$ erg~s$^{-1}$ is present, suggesting an isotropically-emitting IMBH. The spatial coincidence of ULXs with dense star clusters (young clusters and globular clusters) suggests that IMBHS formed in these clusters could be the compact objects in the associated ULXs.  Quasi-periodic oscillations and frequency breaks in XMM power-density spectra of ULXs also suggest that the black hole masses are more consistent with IMBHs than stellar-mass black holes.  Since {\\it all of these ULXs with evidence for IMBHs are high-luminosity ULXs, i.e., L$_X \\gapprox$ 10$^{40}$ erg~s$^{-1},$} we suggest that this class of ULXs is generally powered by accreting IMBHs. } ",
        "introduction": "It has long been suspected that black holes of masses  $\\sim 10^2-10^4\\,M_\\odot$ may form in, for example, the centers of dense stellar clusters (e.g., Wyller 1970; Bahcall \\& Ostriker 1975; Frank \\& Rees 1976; Lightman \\& Shapiro 1977; Marchant \\& Shapiro 1980; Quinlan \\& Shapiro 1987; Portegies Zwart et al. 1999; Ebisuzaki et al. 2001). However, for many years there was no observational evidence for such a mass range.  In roughly the last decade, X-ray and optical observations have revived this possibility.  If such black holes exist, especially in dense stellar clusters, they have a host of implications, especially for cluster dynamical evolution and the generation of gravitational waves.  In this article, we discuss the evidence for intermediate-mass black holes (IMBHs) in Ultra-Luminous X-ray sources (ULXs).  ULXs are extra-nuclear point sources that have X-ray fluxes many times the angle-averaged flux of a $M \\gapprox 20\\,M_\\odot$ black hole accreting at the Eddington limit. Evidence for IMBHs in globular clusters and other astrophysical objects has been discussed in several recent review articles, such as van~der~Marel (2003), and Miller \\& Colbert (2004). We focus here on the ULXs, which we regard as having the most convincing evidence for IMBHs. ",
        "conclusions": "In summary, while many of the low-luminosity ULXs with L$_X \\lapprox$ 5 $\\times$ 10$^{39}$ erg~s$^{-1}$ are consistent with mild-beaming HMXB models (e.g., King et al. 2001, King 2003), there are a significant number of ULXs that are not. The ULXs that do show evidence for isotropically-emitting sub-Eddington IMBHs -- in the form of ``cool disks,'' powerful and narrow QPOs, or suggestive breaks in their PDS -- are all high-luminosity ULXs, with L$_X \\gapprox$ 10$^{40}$ erg~s$^{-1}$, precisely those that are not well explained by mild beaming.  We emphasize that X-ray or optical/NIR observational diagnostics are not yet able to sytematically determine the mass, emission anisotropy, or fuel source of ULXs.  Since the formation mechanism for IMBHs and IMBH-ULXs is not well understood, it is not absolutely certain what fraction of either the ``low-luminosity'' ULXs or the ``high-luminosity'' ULXs have IMBHs.  This holds for the ULXs correlated with star-formation in spiral and starburst galaxies, as well as the ``type-II'' (?) ULXs in elliptical galaxies."
    },
    "0402/astro-ph0402327_arXiv.txt": {
        "abstract": "We report on \\xmm{} spectroscopy of the low-luminosity active galaxies (LLAGN) M81 and NGC~4579 both of which have known black hole masses and well-sampled spectral energy distributions (SED). The iron K$\\alpha$ line profiles from both the LLAGN can be described in terms of two components -- a narrow line at $6.4\\kev$ and a moderately broad line (FWHM $\\sim 2 \\times 10^{4}\\kms$) arising from highly ionized, He-like or H-like species ($E \\sim 6.8\\kev$). We interpret the broad lines arising from an accretion disk the inner edge of which is restricted to large radii ($r_{in} \\sim 100 r_g$). However, the Eddington ratio, ${L}/{L_{Edd}}$, of these sources, is 3--4 orders of magnitude lower than that required to photo-ionize a cold disk to He-like iron. We suggest that the lines can be explained as collisionally ionized X-ray lines arising from the transition region between a hot (radiatively inefficient) flow in the inner regions and a cold disk outside $r \\sim 100r_g$. The accretion flow geometry probed by our {\\it XMM-Newton} observations is consistent with the truncated disk models proposed to explain the SED of LLAGNs. ",
        "introduction": "Active galactic nuclei (AGN) typically display a `big blue bump' in their optical/UV spectra interpreted as blackbody emission arising from optically thick, geometrically thin accretion disks. Their hard X-ray power-law spectrum is thought to be produced by Comptonization of the soft disk photons in a hot corona above the accretion disk (e.g.; Haardt \\& Maraschi 1991). Low luminosity AGN (LLAGN), emitting well below their Eddington limit (${L}/{{L}_{Edd}} \\le 0.01$), typically lack the `big blue bump' (e.g., Ho 1999; Di Matteo et al.~2003 for the case of M87) and for this reason it has been proposed that in these objects the accretion disk consists of of two zones: an outer thin disk that extends from some large radius down to a transition radius and an inner low-radiative efficiency accretion flow close to the black hole.  Fluorescence Fe K$\\alpha$ line emission can be a very powerful probe for the accretion flow geometry around black holes. Such lines have been interpreted as being produced through X-ray irradiation of the accretion disk; therefore the physical width of the line can serve as a trace of the inner accretion disk extent.  Observations of relativistically broadened iron K$\\alpha$ lines from some Seyfert galaxies (e.g., MCG-6-30-15 -- Wilms et al 2001; Fabian et al. 2002), which require that the disk must be irradiated from an X-ray source within $6r_g$, suggests that at least in this case the accretion disk extends to the last stable orbit. It is interesting to note that MCG-6-30-15 (FWHM(H$\\alpha$)$\\sim 2200\\kms$; Sulentic et al. 1998) is like narrow-line Seyfert~1 galaxies (NLS1), which are thought to be accreting at relatively high accretion rates.  In contrast, Lasota et al. (1996) suggested that the accretion disk in most LLAGN must be truncated. Gammie, Narayan, and Blandford (1999) have shown the broadband spectrum of NGC~4258 can be explained by emission from a geometrically thin accretion disk + ADAF model with a transition radius at $\\sim 10-100$ Schwarzschild radii ($r_S$). On the basis of their optical/UV continuum spectra, Quataert et al. (1999), have provided further evidence for such accretion flow geometry and showed that the spectral energy distributions of M~81 and NGC~4579 can be explained by an ADAF + disk with a transition radius at $\\simeq 100 r_S$.  One prediction from any such models, as discussed in Quataert et al (1999), is the lack of a relativistically broadened iron K$\\alpha$ line in the X-ray spectrum of these objects. The presence, or absence, of this feature in X-ray spectra from \\chandra{} and \\xmm{} can therefore be used to test the accretion flow geometry proposed for LLAGN. Specifically, any detection of relativistically broad iron K$\\alpha$ line would strongly argue against truncated disk models.  In this paper, we utilize \\xmm{} observations of two of the best studied LLAGN, NGC~4579 and M~81, to investigate whether the observed iron K$\\alpha$ profiles are consistent with truncated disk models. Both NGC~4579 and M~81 show broad H$\\alpha$ line and are classified as type 1 AGN (Barth et al. 2001; Filippenko \\& Sargent 1985). ",
        "conclusions": "Using the \\xmm{} EPIC observations, we detected complex iron line profiles from two LLAGN NGC~4579 and M~81. The line profiles are well described by a combination of a narrow neutral line and a broad Gaussian (FWHM $\\sim 20000\\kms$) or an accretion-disk line arising from highly ionized (He- or H-like) iron. The line profiles from both the AGN are remarkably similar except that the neutral iron K$\\alpha$ line from M~81 is much weaker than that that from NGC~4579.  The main result of this work is that the broad and highly ionized iron K$\\alpha$ lines from two LLAGNs NGC~4579 and M~81 are consistent with accretion disks whose inner edges are restricted to large radii $r_{in} \\sim 100r_g$."
    },
    "0402/astro-ph0402057_arXiv.txt": {
        "abstract": "Coalescing binary black holes experience an impulsive kick due to anisotropic emission of gravitational waves. We discuss the dynamical consequences of the recoil accompanying massive black hole mergers. Recoil velocities are sufficient to eject most coalescing black holes from dwarf galaxies and globular clusters, which may explain the apparent absence of massive black holes in these systems. Ejection from giant elliptical galaxies would be rare, but coalescing black holes are displaced from the center and fall back on a time scale of order the half-mass crossing time. Displacement of the black holes transfers energy to the stars in the nucleus and can convert a steep density cusp into a core. Radiation recoil calls into question models that grow supermassive black holes from hierarchical mergers of stellar-mass precursors. ",
        "introduction": "In a companion paper (\\citealt{Favata:04}; hereafter Paper I), the amplitude of the recoil velocity experienced by a binary black hole (BH) due to anisotropic emission of gravitational radiation during coalescence is computed. Here we explore some of the consequences of the kicks \\citep{Redmount:89}: the probability that BHs are ejected from galaxies and the implications for BH growth; the time scale for a kicked BH to return to the center of a galaxy; the effect of displacement on nuclear structure; and other observational signatures of the kicks. Unless otherwise indicated, notation is the same as in Paper I.  For inspiral from a circular orbit, the kick velocity is a function of the binary mass ratio $q=m_1/m_2\\le 1$, the BH spins ${\\tilde a}_1$ and ${\\tilde a}_2$, and the initial angle $\\iota$ between the spin of the larger BH and the orbital angular momentum of the binary. Following Paper I, the spin of the smaller BH is ignored. Although Paper I only considers the cases $\\iota=0$ and $\\iota=180$, the recoil for arbitrary inclination is likely to be bounded between these extreme values. Also, the detailed inclination dependence is unimportant in comparison with the large uncertainty already present in the contribution to the recoil from the final plunge and coalescence. We will therefore assume that the restriction to equatorial-prograde/retrograde orbits ($\\tilde{a}_2=[-1,1]$) considered in Paper I encompasses the characteristic range of recoil velocities.  Figure 2b of Paper I shows upper- and lower-limit estimates of the recoil velocity as a function of the effective spin parameter $\\tilde a$ for reduced mass ratio $\\eta=\\mu/M=q/(1+q)^2=0.1$. The {\\em upper limit} for $\\eta=0.1$ is well fit in the range $-0.9\\le\\tilde a\\le 0.8$ by the following fifth-order polynomial: \\begin{eqnarray} \\label{eq:upper} V_{\\rm upper}&=&465\\ {\\rm km\\ s}^{-1} {f(q)\\over f_{\\rm max}}(1- 0.281\\tilde a - 0.0361\\tilde a^2 \\nonumber \\\\ &-& 0.346\\tilde a^3 - 0.374\\tilde a^4 - 0.184\\tilde a^5) . \\end{eqnarray} Fitchett's (1983) scaling function $f(q)/f_{\\rm max}$, with $f(q)=q^2(1-q)/(1+q)^5$, equals $0.433$ for $\\eta=0.1$. The {\\em lower limit} curve of Paper I is well fit by \\begin{eqnarray} \\label{eq:lower} V_{\\rm lower}&=&54.4\\ {\\rm km\\ s}^{-1} {f(q)\\over f_{\\rm max}}(1+ 1.22\\tilde a + 1.04\\tilde a^2 \\nonumber \\\\ &+& 0.977\\tilde a^3 - 0.201\\tilde a^4 - 0.434\\tilde a^5) . \\end{eqnarray}  We convert these expressions into estimates of the bounds on $\\vk$ as follows. First, as discussed in Paper I, there is an ambiguity in how one translates the physical spin parameter ${\\tilde a}_2$ of the larger hole into the effective spin parameter ${\\tilde a}$ of equations (\\ref{eq:upper}) and (\\ref{eq:lower}). Here we adopt the \\citet{Damour:01} relation $\\tilde a=(1+3q/4)(1+q)^{-2}{\\tilde a}_2$. Second, Fitchett's scaling function assumes that both bodies are non-spinning, and vanishes when $q=1$. In fact, when $\\tilde a\\neq0$, significant recoil would occur even for $q=1$ due to spin-orbit coupling. We can guess the approximate form of a new scaling function by examining the spin-orbit corrections \\citep{kidder} to Fitchett's recoil formula. For equatorial orbits, equation (4) of Paper I suggests that $f(q)$ should be multiplied by the factor $|1 + (7/29)\\tilde{a}_2/(1-q)|/|1+ (7/29)\\tilde{a}_2/(1-q')|$, where $q'=0.127$ is the value used in defining $V_{\\rm upper}$ and $V_{\\rm lower}$ in equations (\\ref{eq:upper}) and (\\ref{eq:lower}).  Figure~\\ref{fig:vplunge} plots upper and lower limits to $\\vk$ as functions of ${\\tilde a}_2$ and $q$. The average over $\\tilde{a}_2$ of the upper limit estimates are $\\sim(138, 444, 154)$ km s$^{-1}$ for $q=(0.1,0.4,0.8)$; Figure~\\ref{fig:vplunge} suggests a weak dependence on $\\tilde{a}_2$. Lower limit estimates are more strongly spin-dependent; the averages over $\\tilde{a}_2$ are $\\sim(21.1,63.6,24.9)$ km s$^{-1}$ for the same values of $q$. For moderately large spins ($\\tilde{a}_2\\gap 0.8$) and prograde capture, the lower limit estimates exceed $100$ km s$^{-1}$ for $0.2\\lap q\\lap 0.6$. In what follows, we will assume that $\\sim 500$ km s$^{-1}$ is an absolute upper limit to $\\vk$. ",
        "conclusions": ""
    },
    "0402/astro-ph0402261_arXiv.txt": {
        "abstract": "{Low-frequency longitudinal oscillations of a flaring coronal loop are studied numerically. In the recent work of Nakariakov et al. A\\&A, 414, L25-L28 (2004) it has been shown that the time dependences of density and velocity in a flaring loop contain well-pronounced quasi-harmonic oscillations associated with a 2nd harmonic of a standing slow magnetoacoustic wave. In this work we investigate physical nature of these oscillations in greater detail, namely, we study their spectrum (using periodogram technique) and how does heat positioning affects the mode excitation. We found that excitation of such oscillations practically independent of the positioning of the heat deposition in the loop. Because of the change of the background temperature and density, the phase shift between the density and velocity perturbations is not exactly equal to the quarter of the period, it varies along the loop and is time dependent, especially in the case of one footpoint (asymmetric) heating. ",
        "introduction": "Magnetohydrodynamic (MHD) coronal seismology is the main reason for studying waves  in the solar corona. Also such studies are important in connection with coronal heating and solar wind acceleration problems. Observational evidence in the EUV coronal emission of coronal waves and oscillations is numerous (e.g. \\citep{o99,ow02}). Radio band observations also demonstrate various kinds of oscillations (e.g., the quasi-periodic pulsations, or QPP, see \\citep{a87} for a review), usually with the periods from a few seconds to tens of seconds.  Also, decimeter and microwave observations  show much longer periodicities, often in association with a flare. For example, \\citep{wx00} observed QPP with the periods of about 50~s at 1.42 and 2~GHz (in association with an M4.4 X-ray flare). Similar periodicities have been observed in the X-ray band  (e.g., \\citep{m97,terekhov02}) and in the white-light emission associated with the stellar flaring loops \\citep{mathio}. A possible interpretation of these medium period QPPs may be interpreted in terms of kink or torsional modes \\citep{zs89}. In our previous, preliminary  study \\citep{nt04}, we have outlined an alternative, {\\it simpler, thus more aesthetically attractive}, mechanism for generation of long-period QPPs. In this study we {\\it demonstrate in detail,} that in a coronal loop an impulsive (time-transient) energy release efficiently generates the second spatial harmonics of a slow magnetoacoustic mode. In particular, in the present study we study the spectrum (using periodogram technique) of these oscillations and how does heat positioning affects the mode excitation. ",
        "conclusions": "Initially we have used 1D radiative hydrodynamics loop model which incorporates the effects of gravitational stratification, heat conduction, radiative losses, added external heat input, presence of Helium, hydrodynamic non-linearity, and bulk Braginskii viscosity to simulate flares \\citep{t04}. As a byproduct of that study, in practically all our numerical runs we have detected quasi periodic oscillations in all physical quantities \\citep{nt04}. In fact, such oscillations are frequently seen during the solar flares observed in X-rays, 8-20 keV (e.g. \\citet{terekhov02}) as well as stellar flares observed in white-light (e.g. \\citet{mathio}). Our present analysis shows, {\\it in detail}, that quasi periodic oscillations seen in our numerical simulations bear many similar features as the observed ones. In summary \\citep{nt04} and the present study established the following features: \\begin{itemize} \\item We show that the time dependences of density and temperature in a flaring loop contain well-pronounced quasi-harmonic oscillations associated with standing slow magnetoacoustic modes of the loop. \\item For {\\it a wide range of physical parameters}, the dominant mode is the second spatial harmonic, with a velocity oscillation node and the density oscillation maximum at the loop apex. {\\it This result is practically independent of the positioning of the heat deposition in the loop}. \\item Because of the change of the background temperature and density, and the fact that density gradients in the transition region are not providing perfect reflecting boundary conditions for the formation of standing sound waves, the phase shift between the density and velocity perturbations is not exactly equal to the quarter of the period. \\item We conclude that the oscillations in the white light, the radio and X-ray light curves observed during solar and stellar flares may be produced by the slow standing mode, with the period determined by the loop temperature and length. \\item For a typical solar flaring loop the period of oscillations is shown to be about a few minutes, while amplitudes are typically few percent. \\end{itemize}  The novelties brought about by this study are that by studying the spectrum and phase shift of these oscillations we provide more definite proof that these oscillations are indeed 2nd harmonic of a standing sound wave, and that the single footpoint (asymmetric) heat positioning still produces 2nd spatial harmonic, though it is more complex than the apex (symmetric) heating (due to the presence of flows)."
    },
    "0402/astro-ph0402275_arXiv.txt": {
        "abstract": "I review the operational capabilities of the \\emph{Chandra} X-ray Observatory, including some of the spectacular results obtained by the general observer community.  A natural theme of this talk is that \\emph{Chandra} is revealing outflows of great quantities of energy that were not previously observable. I highlight the \\emph{Chandra} studies of powerful X-ray jets.  This subject is only possible due to the sub-arcsecond resolution of the X-ray telescope. ",
        "introduction": "X-ray Observatory}   The Advanced X-ray Astrophysics Facility (\\emph{AXAF}) was developed to serve the entire astronomical community for attacking a wide range of problems.\\cite{weisskopf87,weisskopf96,weisskopf00} The requirements to do this included large telescope area, access to the entire sky with  more than 85\\% available at any time, high observing efficiency and a long operational lifetime, instruments which could provide imaging and spectroscopy, including spatially resolved spectroscopy with at least modest (E/$\\Delta E \\approx$ 10 to 50) energy resolution, and the ability to locate the measured photons on the sky. However, by far the most stringent and crucial requirement was for the telescope to be able to image to better than 0.5\\arcsec\\ FWHM. More precisely, it was required that a 1\\arcsec\\ diameter circle about a point source would contain 70\\%, and 20\\%, of the photons imaged at 1 and at 8 keV, respectively. This allows high contrast for large dynamic range, and gives an imaging point spread function (PSF) which is not a strong function of energy. Prior to launch, \\emph{AXAF} was renamed the \\emph{Chandra} X-ray Observatory (CXO), in honor of the Indian-American astrophysicist Subramanyan Chandrasekhar.  The exquisite X-ray imaging capability of the High Resolution Mirror Assembly (HRMA) is the key to the scientific power of the \\emph{Chandra} observatory. Imaging to 0.5\\arcsec\\, in contrast to the $\\sim$5\\arcsec\\ imaging capability of the previous \\emph{Einstein} and \\emph{ROSAT} X-ray missions, is truly a 100-fold improvement in the imaging capability. This enables \\emph{Chandra} to study jets and outflows in quasars and active galaxies which had previously been considered ``point'' sources, to reveal structure and interactions within clusters of galaxies for which the distribution of hot gas had previously been considered smooth and symmetric, and to allow point source detection to fluxes 100 times fainter due to the reduction of the detector area which accumulates background.  One of two different imaging instruments can be moved into the telescope aim point at any time. The High Resolution Camera (HRC)  uses micro channel plates (MCP) to convert the X-ray, and to amplify the resulting electrons. The X-ray is located by determining the centroid of the resulting electron cloud. The Advanced Camera for Imaging Spectroscopy (ACIS, formerly the ``AXAF Camera for Imaging Spectroscopy''), uses charge coupled devices (CCD) to convert X-rays into a number of electrons directly proportional to the photon energy, and clocks these out in a fixed pattern which indicates the position at which the X-ray was imaged. A low energy transmission grating (LETG) gives dispersive resolution optimized for the 0.07 to 2 keV range, while a high energy transmission grating (HETG) consists of two separate sets of  gratings optimized for the 0.4 to 5 keV range (medium energy grating) and the 0.8 to 10 keV range (high energy grating). If desired, either the LETG or HETG (but not both!) can be inserted into the optical path.    A unique feature of celestial X-ray astronomy is that it is based on single photon counting.  The energies are large enough, from 0.1 to 10 keV in the case of \\emph{Chandra}, that this is possible, and from most celestial sources the fluxes are so low that it is imperative. At the limiting sensitivity in the deepest fields, 10 photons in a two week integration represents a significant detection.  Operationally this means that the satellite need not point rigidly to within a very small fraction of the desired resolution for the duration of the exposure.  In fact, it is desireable to force the telescope to dither over many pixels in order to use an averaged calibration, since it is not feasible to calibrate each resolution element to the accuracy ultimately desired. An optical CCD camera which simultaneously measures stars on the sky and fiducial lights on the instruments enables post facto ground reconstruction of the image to better than 0.1\\arcsec\\ . ",
        "conclusions": "\\emph{Chandra} observations have discovered the evidence of large outflows of energy and material previously  not seen directly. In the case of pulsar wind nebulae observations prove that jets are ejected along the spin axis, and that material is also ejected in the equatorial plane. Theoretically these would be oppositely charged plasmas.\\cite{michel00} In clusters of galaxies, the images prove a close X-ray radio connection. In quasars we have seen that the kinetic fluxes carried by jets represent very large energy contents. We might hypothesize that jets from the radio galaxies at the center of cooling flow clusters may carry similarly large energy fluxes -- clearly adequate to compensate for the radiative cooling of the gas, even if we allow that the radio source only has a $\\sim$10\\% duty cycle. To actually compensate for a cooling flow probably requires a feedback mechanism involving the infalling gas and the central AGN.\\cite{nulsen03} Observations of the Perseus cluster of galaxies led to the suggestion that sound waves provide the mechanism for dissipation of the jet energy quasi-isotropically through the cluster.\\cite{fabian03}  If the X-ray emission from jets does arise from the IC/CMB, then the same objects which we have already seen would be easily detectable at any arbitrarily larger redshift at which they might occur.\\cite{schwartz02} This is because the emission is proportional to the energy density of the target photons, which will increase as T$^4\\propto$T$_{0}^4$ (1+z)$^4$, compensating exactly for the cosmological diminution of surface brightness.  This review has been based primarily on the first three years of \\emph{Chandra} operation. We are just finishing the fourth year and starting the approved fifth year program. At the \\emph{Chandra} X-ray Center we have performed a study indicating the expected lifetime of the observatory will be at least 15 years. We expect these future observations to unveil the physical processes in all astronomical systems, including the few topics included in this review, in vastly greater breadth and detail."
    },
    "0402/astro-ph0402569_arXiv.txt": {
        "abstract": "We present the results of a radial velocity survey of a sample of Hyades stars, and discuss the effects of stellar activity on radial velocity measurements. The level of radial velocity scatter due to rotational modulation of stellar surface features for the Hyades is in agreement with the predictions of \\citet{SaDo97}- the maximum radial velocity rms of up to $\\sim$50~m~s$^{-1}$, with an average rms of $\\sim$16~m~s$^{-1}$. In this sample of 94 stars, we find 1 new binary, 2 stars with linear trends indicative of binary companions, and no close-in giant planets. We discuss the limits on extrasolar planet detection in the Hyades and the constraints imposed on radial velocity surveys of young stars. ",
        "introduction": "Radial velocity ($v_{\\rm r}$) surveys for extrasolar planets have been extremely successful (e.g. \\nocite{BuMaWi96} Butler et al. 1996). These surveys have, however, largely excluded young, active stars \\citep{VoBuMa02, CuMaBu99, SaDo97}. The reason given was that the activity levels of young stars is significant enough to cause large variations in the measured $v_{\\rm r}$. Although it does not introduce a true $v_{\\rm r}$ shift \\citep[e.g. Saar \\& Donahue 1997, hereafter SD97,][]{Ha02}, the apparent shift is caused by a change in the line shape of the absorption features. SD97 quantify the predicted amplitude of this phenomenon, and it has been observationally confirmed by several groups \\citep[e.g.][]{QuHeSi01,HeDoBa02, PaSaCo04, SaFi00, SaBuMa98}. While detection of extrasolar planets around young stars will be complicated by these spectral line profile variations, there is much to learn about the frequency of planets and their orbital characteristics at all stellar ages. So, it is necessary to learn the limitations of the techniques employed in planet detection and then proceed (if possible) with planet searches.  We present here the results of the radial velocity search for extrasolar planets in a sample of Hyades dwarfs. Primarily, we discuss the mean level of radial velocity noise caused by stellar magnetic activity and the possibilities of detecting planets in Hyades-aged stars. ",
        "conclusions": "We can determine significant period in data by various techniques, including those discussed in this paper. But, it is most useful to understand when significant periods are real or simply artifacts of sampling. The analysis of the periodogram produces periods with FAPs~$\\sim$10\\% for several stars. Phasing the data to these periods produces periodic curves (by-eye inspection). This is inadequate. Therefore, we have employed the method of \\citet{NeAn98} to explore the significance of detections. All short-periods detected turn out to be artifacts of the sampling and of the quality of the data.  The detection of planets around young stars is complicated by the rotational modulation of stellar active regions. The activity not only causes high levels of $v_{\\rm r}$ noise but can also yield periodic variations in the measured $v_{\\rm r}$ causing false detections. The procedure we adopted \\citep{NeAn98} picks out all significant signals given the quantity and quality of data, so we must be careful in the identification of the source of variability. In our data, we find no evidence for short-period massive planets or brown dwarfs.  Finally, of the 94 stars in this sample, 6 are either suspected or identified binaries and 1 has a velocity rms which is somewhat arge but further observations are required to say anything more concrete- it is still within possible ``jitter\" from high activity levels.  Future detection of extrasolar planets around young stars via the radial velocity method will be limited to high-mass planets and in particular, those with short orbital periods. Constraints on telescope time needed for these surveys becomes clear. In order to increase the odds of planet detection, as current planet searches have determined that only $\\sim$1\\% of stars do have ``hot Jupiters\", data must be sampled several times a month which requires a great deal of allocated telescope time."
    },
    "0402/astro-ph0402043_arXiv.txt": {
        "abstract": "We determined C, N and $\\alpha$-element relative abundances in the gas surrounding six QSOs at an average redshift of $< z > \\simeq 2.4$, by studying six narrow associated absorption systems in UVES high-resolution spectra. We found five systems with a metallicity (measured by C/H) consistent or above the solar value. The ionization structure observed in the associated systems is clearly different from that of the intervening ones, indicating that the associated systems are influenced by the strong UV flux from the QSO. There is a possible correlation (anticorrelation) between [N/C] ([Si/C]) and [C/H]  of the studied associated systems, and [N/C]~$\\ge 0$ when [C/H]~$\\ge 0$. We have compared  these observational results with the predictions of a model simulating the joint evolution of QSOs and their spheroidal hosts. The agreement turns out to be very good, in particular, the case envisaging massive haloes and high star-formation rates recovers both the correlation between [N/C] and [C/H] and the anticorrelation for [Si/C] vs. [C/H]. Narrow associated absorption systems prove to be powerful tracers of the chemical abundances in gas belonging to high redshift spheroidal galaxies. The outflow of this same gas, triggered by the QSO feedback, is probably going to contribute to the early enrichment of the surrounding intergalactic medium. A larger statistics, possibly increasing the number of ionisation stages, chemical elements and the redshift range, would allow us to put firm constraints on detailed chemical evolution models of galaxies at high redshifts. ",
        "introduction": "In this work, we want to address the star formation history and the evolution of massive early-type galaxies at high redshifts by measuring in a reliable way the metallicity and the chemical abundances of gas belonging to host galaxies and environments of QSOs.  Once considered rare and exotic objects, QSOs could instead represent a necessary phase in the evolution of massive early-type galaxies. This interpretation is supported by several pieces of evidence: Massive Dark Objects (MDOs, generally interpreted as dormant black holes) with masses in the range $\\sim 10^6$~-~$3\\times10^9$ M$_{\\sun}$ are present in essentially all local galaxies with a substantial spheroidal component \\citep[see ][ for a review]{korm:geb}, on the other hand the host galaxies of low redshift powerful AGN (radio-loud and radio-quiet QSOs and radio galaxies) are, in all the studied cases, luminous elliptical galaxies with $L>L^{\\ast}$ \\citep{dunlop03}.  The observational properties inferred for cluster and field elliptical galaxies up to redshift $z \\sim 1$ imply a high uniformity and synchronization in the galaxy formation process \\citep[e.g.][]{ellis97,bernardi98}.  The evolution with redshift of their optical-IR colours \\citep{stanford98} is consistent with the passive evolution of an old stellar population formed at $z \\ge 2-3$ and the measured positive [Mg/Fe] elemental ratio can be explained by a short and intense star formation burst \\citep[e.g. ][]{wortheyeal,matteucci94}.  In the standard framework of the hierarchical evolution of structures in a cold dark matter (CDM) universe large objects form by a sequence of mergers of smaller proto-galaxies. In particular, massive ellipticals are generated at low redshifts ($z \\le 2$) from the merger of two large disk galaxies which formed stars at a constant moderate rate up to that moment \\citep[e.g. ][]{baugheal,KC98}. In the merging event, the black holes (BHs) pre-existing in the progenitor galaxies coalesce and a fraction of the cold gas is accreted by the new BH which activates as a QSO, the rest of the cold gas is transformed into stars in a sudden burst \\citep{WB98,kauff:haeh,marta03,mencieal03}.  A different prescription in the framework of the hierarchical scenario is the {\\sl anti-hierarchical baryon collapse} where the formation of stars and of the central BH takes place on shorter time-scales within more massive dark matter haloes \\citep*{monaco00,granato01,archi02,granato04}. Supernova heating and QSO feedback are the physical processes that reverse the order of formation of galaxies compared to that of DM haloes because they slow down star formation most effectively in shallow potential wells. In the more massive  DM haloes star formation goes on rapidly causing at the same time the growth of the central BH which accretes the cold gas slowed down by the radiation drag. When the QSO activates, strong winds originate sweeping the interstellar medium and halting both the star formation and the BH growth. The time delay between the star formation onset and the peak of the QSO activity is again shorter for larger haloes. For the most massive galaxies ($M_{\\rm halo} \\ga  10^{12}\\ M_{\\sun}$) virializing at $3 \\le z \\le 6$, this time is $< 1$ Gyr, implying that the bulk of star formation may be completed before type Ia supernovae have the time to significantly enrich the interstellar medium with iron. A detailed analysis of the chemical evolution expected for this model is reported in  \\citet{romano02}.  The two above  described scenarii predict different chemical abundances, in particular at redshifts larger than $\\sim 2$. The metallicity and the elemental abundances of high redshift galaxies are hard to measure; on the other hand, high and intermediate resolution spectra of QSOs at redshifts as large as $z \\sim 6$ can be easily obtained with the present instrumentation.  We studied associated narrow\\footnote{The adjective ``narrow'' is used to distinguish this class of absorptions from the Broad Absorption Lines characterised by FWHM~$> 2000$ \\kms\\ and arising in gas ejected by the QSO at large velocities (see also Sections~2 and 7)} absorption lines exploiting high resolution, high signal-to-noise ratio spectra of $2 < z< 3$ QSOs obtained with the UVES spectrograph and a model for the photoionisation of the gas to derive chemical abundances in the QSO environments. Our results suggest that at these redshifts the gas associated with the QSO and with its host galaxy has already been enriched by the products of an intense star formation episode.  Section~2 introduces  the diagnostics that we used to determine the chemical abundances in QSO environments and reports previous results. In Section~3 we describe the selection criteria and the characteristics of our sample of associated narrow absorption line systems; the adopted photoionisation model and the methodology are reported in Section 4. Section~5 is devoted to the description of our results, which are compared with model predictions in Section~6. In Section~7, we summarise the results on QSO chemical abundances obtained using other methods. We draw our conclusions in Section~8. ",
        "conclusions": "Up to now the main approach to study the chemical abundances in QSO environments has been the analysis of BELs observed in their spectra. Metallicities determined from BELs are consistent with solar or slightly supersolar values without a significant evolution in redshift. Other elemental abundances are very difficult to measure, in particular determinations of the ratio $\\alpha$/Fe are very uncertain.  Associated narrow absorptions are complementary probes of the physical status of QSO-elliptical systems with respect to BELs. In general, they can be due to gas belonging to the interstellar medium of the galaxy, outflowing under the effect of the QSO or re-infalling on the QSO itself. Furthermore, it is more straightforward to derive chemical abundances from absorption lines than from emission lines. We need only to determine and apply the proper ionisation corrections to convert the measured ionic column densities into relative abundances.  In this paper, we selected six narrow absorption systems lying within 5000 \\kms\\ from a QSO emission redshift and determined the abundances of C, N and $\\alpha$-elements in the gas they originate from. We used high resolution, high signal-to-noise UVES QSO spectra and applied a procedure based on the photoionisation code Cloudy to compute the chemical abundances starting from the measured column densities.  We found that all systems but one in our sample have metallicities (measured by carbon) consistent with or larger than solar. We found also a possible correlation of [N/C] and an anticorrelation of [Si/C] with [C/H] with supersolar values of [Si/C]. These results are suggestive of rapid enrichment due to a short star formation burst, of duration $t_{\\rm burst}\\sim$ 1 Gyr (see Section~5). Since the very high luminosity QSOs in our sample should have $M_{\\rm BH}\\geq 10^9$ $M_{\\sun}$ , assuming $M_{\\rm sph}/M_{\\rm BH}\\sim 1000$ \\citep{mclure:dunlop} we expect SFR$\\geq 1000$ $M_{\\sun}$ yr$^{-1}$ in their hosts.  The predictions of the model of chemical evolution for a spheroidal galaxy where the star formation depends on stellar and QSO feedback are in good agreement with the observations. In particular, the agreement improves when taking into account a clumping factor \\citep{granato04}, which allows the gas to be efficiently converted into stars also in very massive dark haloes with  SFR$\\geq 1000$ $M_{\\sun}$ yr$^{-1}$.  In this way, narrow associated QSO absorption systems proved to be extremely useful in the study of the QSO environment, in particular when there is evidence of their intrinsicness. They can be used as estimators of the chemical abundances in high redshift spheroidal galaxies which are not easily determined otherwise. The probed gas will probably  be ejected from the galaxy due to the QSO feedback, thus we are also observing a potential source of enrichment of the intergalactic medium at high redshift.  In order to obtain a deeper insight in the evolution of QSO host-galaxies and environments it is essential to enlarge the data sample. In particular, obtaining high signal-to-noise spectra in the UV to reliably measure the doubly-ionised lines of C and N and increasing the redshift range especially at large values. Indeed, the five $z\\sim 4$ AALs analised up to now \\citep{savaglioeal97} seems to indicate a slightly lower average metallicity, [C/H]~$\\sim-0.5$, than for the bulk of the sample at redshift $z \\sim 2 - 2.5$.  More data will be fundamental to verify the observed correlations and to constrain the predictions of theoretical models."
    },
    "0402/astro-ph0402333_arXiv.txt": {
        "abstract": "{ In the present work, we propose a new method aiming at extracting the kinetic Sunyaev-Zel'dovich (KSZ) temperature fluctuations embedded in the primary anisotropies of the cosmic microwave background (CMB). We base our study on simulated maps without noise and we consider very simple and minimal assumptions. Our method essentially takes benefit from the spatial correlation between KSZ and the Compton parameter distribution associated with the thermal Sunyaev-Zel'dovich (TSZ) effect of the galaxy clusters, the later can be obtained by means of multi-frequency based component separation techniques.  We reconstruct the KSZ signal by interpolating the CMB fluctuations without making any hypothesis besides the CMB fluctuations are Gaussian distributed.  We present two ways of estimating the KSZ fluctuations, after the interpolation step. In the first way we use a blind technique based on canonical Principal Component Analysis, while the second uses a minimisation criterion based on the fact that KSZ dominates a small angular scales and that it follows a non-Gaussian distribution. We show  using the correlation between the input and reconstructed KSZ map that the latter can be reconstructed in a very satisfactory manner  (average correlation coefficient between 0.62 and 0.90), furthermore both the retrieved KSZ power spectrum and temperature fluctuation distribution are in quite good agreement with the original signal. The ratio between the input and reconstructed power spectrum is indeed very close to one up to a multipole $\\ell\\sim 200$ in the best case. The method presented here can be considered as a promising starting point to identify in CMB observations the temperature fluctuation associated with the KSZ effect. ",
        "introduction": "The Cosmic Microwave Background (CMB) temperature anisotropies contain the contribution of both the primary cosmological signal, directly related to the initial density fluctuations, and the foregrounds amongst which are the secondary anisotropies generated after matter-radiation decoupling.  They arise from the interaction of the CMB photons with the matter and can be of a gravitational type (e.g. Rees-Sciama effect (Rees \\& Sciama 1968)), or of a scattering type when the matter is ionised (e.g. Sunyaev-Zel'dovich (SZ) effect (Sunyaev \\& Zel'dovich 1972) or Ostriker-Vishniac effect (Ostriker \\& Vishniac 1986; Vishniac 1987)). Among all these secondary anisotropies, the dominant effect is the SZ effect.  It represents the inverse Compton scattering of the CMB photons by the free electrons of the ionised and hot intra-cluster gas. It results in the so-called thermal SZ (TSZ) effect whose amplitude is characterised by the Compton parameter $y$ (the integral of the pressure along the line of sight). The TSZ amplitude thus depends only on the cluster electron temperature and density distributions.  The inverse Compton effect moves the CMB photons from the lower to the higher frequencies of the spectrum. This results in a peculiar spectral signature with a decrement at long wavelengths and an increment at short wavelengths. When the galaxy cluster moves with respect to the CMB rest frame, {with a peculiar radial velocity $v_r$}, the Doppler shift induces an additional effect often called the kinetic SZ (KSZ) effect, which generates temperature anisotropies with the same spectral signature, at least in the non-relativistic approximation, as the primary CMB fluctuations.  The interest of the TSZ effect for cosmology has been recognised very early (see reviews by \\cite{rephaeli95}, \\cite{birkinshaw99} and \\cite{carlstrom02}). It is a powerful tool to detect high redshift galaxy clusters since it is redshift independent. In combination with X-ray observations it can be used to determine the Hubble constant and probe the intra-cluster gas distribution. Moreover, the KSZ effect may be the one of the best ways of measuring the cluster peculiar velocities by combining thermal and kinetic effects \\cite[]{suniaev80}.  The advantages of this method are: (i) it yields directly the peculiar velocities, bypassing the need to measure inaccurate distance indicators \\cite[]{faber76,tully77}; (ii) the method has a physical explanation and (iii) it is independent of distance.  The KSZ can be distinguished from the TSZ effect due to the different frequency dependence of their intensities.  The KSZ intensity reaches its maximum at a frequency of $\\sim 218 $ GHz, just where the TSZ intensity is zero .  Hence, this is the optimal frequency to the detect the KSZ signal. It has also been shown \\cite[]{hobson98,bouchet99,baccigalupi00,delab02,kuo02,maisinger03} that the TSZ signal can be extracted from the other astrophysical contribution by component separation techniques (Wiener filtering, Maximum Entropy, Independent Component Analysis, ...).  Despite the scientific interest of the KSZ effect as a probe of large scale matter distribution and structure formation theories, very few measurements of the peculiar velocities were achieved \\cite[]{holzapfel97,lamarre98,benson03}. As a consequence, very few methods have been proposed so far to address the specific underlying question of {\\it separating the secondary KSZ fluctuations from the primary anisotropies}. In an early work, \\cite{haehnelt96} used an optimal filtering (Wiener), with a spatial filter derived from X-ray observations of galaxy clusters, that minimises the confusion with CMB. However, this method implied the knowledge of the CMB power spectrum. \\cite{aghanim97} rather used a matched filter optimised on simulated data and independent of the underlying CMB model. Recently \\cite{hobson03} presented a Bayesian approach for detecting and characterising the signal from discrete objects embedded in a diffuse background. They showed that this approach is around twice as sensitive as the linear optimal filter approach proposed by \\cite{haehnelt96}.  In the present study, we propose a new method optimised to extract from the primary anisotropies, the temperature fluctuations, associated with the KSZ effect. The method is based on the fact that we have two sets of maps (provided, in a realistic case, by component separation techniques), the first set contains both CMB and KSZ temperature fluctuations and the second set consists of Compton parameter maps associated with the TSZ effect which is used as a spatial template.  In our study, we do not use real (observed) maps but we rather use two sets of simulated maps. We were able to retrieve, in the best possible way, the amplitude and the distribution of the temperature fluctuations associated with KSZ together with the associated power spectrum. ",
        "conclusions": "{ In this first attempt to extract a map of the KSZ temperature fluctuations from the CMB anisotropies we use a method which is based on very simple and minimal assumptions. We discuss the issue of noise and astrophysical contributions but we do not take them explicitly into account. Therefore, our results show the intrinsic limitations of the method in terms of reconstructing a KSZ map from a mixture of CMB and KSZ anisotropies. We demonstrate that the 15 KSZ reconstructed maps are in quite good agreement with the original input signal with a correlation coefficient between original and reconstructed maps of 0.78 on average, and an error on the standard deviation of the reconstructed KSZ map of only 5\\% on average.  To achieve these results, we use the hypothesis that a first step component separation provides us with: (i) a map of Compton parameters for the TSZ effect of galaxy clusters, and (ii) a map of temperature fluctuations for the primary CMB + KSZ cluster signal. Our method essentially takes benefit from the spatial correlation between KSZ and TSZ effects towards the same galaxy clusters. This correlation allows us to use the TSZ map as a spatial template in order to mask, in the CMB + KSZ map, the pixels where the clusters must have imprinted an SZ fluctuation. In practice a series of TSZ thresholds is defined and for each threshold, we estimate the corresponding KSZ signal by interpolating the CMB fluctuations on the masked pixels.  The series of estimated KSZ maps finally is used to reconstruct the KSZ map through the minimisation of a criterion taking into account two statistical properties of the KSZ signal (KSZ dominates over the primary anisotropies at small scales, KSZ fluctuations are non-Gaussian distributed).}"
    },
    "0402/astro-ph0402105_arXiv.txt": {
        "abstract": "{The high-frequency-peaked BL~Lac, MS\\,0205.7+3509 was observed twice with \\emph{XMM-Newton}.  Both X-ray spectra are synchrotron-dominated, with mean 0.2--10\\,keV fluxes of $2.80\\pm0.01$ and $3.34\\pm0.02\\times10^{-12}$\\,erg\\,cm$^{-2}$\\,s$^{-1}$. The X-ray spectra are well fit by a power-law with absorption above the Galactic value, however no absorption edges are detected, implying a low metallicity absorber ($Z_{\\sun} = 0.04^{+0.03}_{-0.01}$) or an absorber with redshift above one (best-fit $z=2.1$ for an absorber with solar abundances).  In either case the absorbing column density must be $\\sim9\\times10^{21}$\\,cm$^{-2}$.  A new optical spectrum is presented, with a \\ion{Mg}{ii} absorption doublet detected at $z=0.351$, but no other significant features.  The optical spectrum shows little reddening, implying a low dust to gas ratio in the absorber. MS\\,0205.7+3509 must therefore be viewed through a high column density, low-metallicity gas cloud, probably at $z=0.351$ and associated with the galaxy that has been shown to be within $\\sim2$\\arcsec\\ of the BL~Lac. ",
        "introduction": "Blazars are divided into BL~Lacs and quasars (either flat-spectrum radio-loud, optically violently variable, highly-polarised or core-dominated) based on the strength of emission lines in the optical spectrum \\citep[and references therein]{1997A&A...325..109S}.   It has been suggested that sources with BL Lac characteristics are actually gravitationally microlensed quasars \\citep{1986A&A...157..383N,1990Natur.344...45O}. In these cases, it is expected that stellar mass lenses in a foreground galaxy significantly amplifies the central QSO continuum source but not the emission from the line-emitting regions and that variations in the relative source-lens position could account for the rapid variability observed in many BL~Lacs. Sources of this kind should clearly have foreground galaxies, which would result in an apparent decentering of the AGN from the ``host\" and an excess of absorption in these sources. However, the suggestion that BL~Lacs are gravitationally microlensed quasars can be discounted for most BL~Lacs \\citep{1992MNRAS.257..404P} and only a few remain as possible or probable candidates, most notably AO\\,0235+164 which appears to have foreground absorption \\citep{1993ApJ...415..101A,1994ApJ...432..554M,1996ApJ...459..156M} and a companion AGN \\citep{1996AJ....112.2533B}. Other candidates include PKS\\,0537-441 \\citep[which shows rapid microvariability, but does not show evidence for a foreground object in optical imaging or spectroscopy,][]{1999A&AS..135..477R,2002A&A...392..407P}, and B2\\,1308+326 which has characteristics intermediate between BL~Lacs and quasars \\citep{1993ApJ...410...39G,2000A&A...364...43W}, but where high resolution imaging of the BL~Lac with the HST WFPC2 \\citep{1999ApJ...512...88U} was consistent with a point source.  MS\\,0205.7+3509 is another such rare candidate, and while deep imaging has revealed that the BL~Lac is centred on a host galaxy that is likely an elliptical and not offset in a spiral host as had been inferred from previous observations \\citep{1995ApJ...454...55S}, a companion galaxy was also detected very close to the BL~Lac line of sight \\citep{1997A&A...321..374F} which had caused the previous inference of decentering in a spiral host to be made. X-ray observations with \\emph{ROSAT} \\citep{1995ApJ...454...55S} and \\emph{ASCA} \\citep{1999A&A...345..414W} showed the existence of absorption well above the Galactic level, indeed at a level second only among BL~Lacs to PKS\\,1413+135 \\citep{2002AJ....124.2401P}.  These results, on MS\\,0205.7+3509, imply that the X-ray absorber is in the companion galaxy which is foreground to the AGN.  It has been suggested that stars in the halo of this companion galaxy could be responsible for microlensing of the BL~Lac \\citep{1997A&A...321..374F,1999A&A...345..414W}.  In spite of the relatively good spectral resolution of \\emph{ASCA}, the redshift of the absorber could not be constrained from those observations \\citep{1999A&A...345..414W}.  A redshift of $z=0.318$ was proposed based on the tentative detection of a \\ion{Ca}{ii} absorption system reported in the optical spectrum of this source \\citep{1991ApJ...380...49M} indicating the possible redshift either of the host galaxy or of a foreground absorber and to date this redshift had been used as the best available \\citep{2000AJ....120.1626R,1995ApJ...454...55S}.  Though somewhat lower in terms of column density, MS\\,0205.7+3509 is six times brighter in X-rays than PKS\\,1413+135, making it one of the best available cases in which to study absorption in the hot phase of the ISM in a galaxy that is not at low redshift. \\emph{XMM-Newton} observations were performed in an attempt to determine the nature of the X-ray absorber in MS\\,0205.7+3509, in particular in the context of a foreground lensing galaxy.  Results from these data are presented in this paper. Sect.~2 deals with the observations and the data reduction procedures; results from the X-ray and optical data are in Sect.~3.  A discussion of these results and a summary of our conclusions are given in Sect.~4. Uncertainties given are 90\\% confidence limits unless otherwise stated. A flat universe with $H_{0}=75$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$ and $\\Omega_\\Lambda=0.7$ are assumed throughout. ",
        "conclusions": "Observations with \\emph{ASCA} indicated the peak of the synchrotron component in MS\\,0205.7+3509 to be between the UV and soft X-rays at the time of those observations \\citep{1999A&A...345..414W}, consistent with results from the first \\emph{XMM-Newton} observation.  The harder spectrum and higher flux during the second observation imply a shift in the synchrotron peak frequency to higher energies.  The featureless nature of the absorbed spectrum implies either that the absorber is metal poor or that the absorbing material is at a redshift much higher than the $z=0.351$ \\ion{Mg}{ii} absorption system detected in the optical spectrum. In either case the absorbing gas column density is similarly high ($\\sim10^{22}$\\,cm$^{-2}$), implying an optical extinction that is not observed \\citep[A$_{\\rm V} \\simeq 5$ at $z=0$ where the gas to dust ratio is similar to the Galactic value,][]{1978ApJ...224..132B}, a fact that supports the low-redshift, low-metallicity model for the absorber.  It is possible that the absorber is associated with the BL~Lac host galaxy. But the lack of gas and dust in elliptical galaxies \\citep{1999sfet.conf..119K}, the typical hosts of BL~Lacs, and the proximity of the companion galaxy \\citep[2.3\\arcsec,\\ corresponding to an apparent linear distance of $\\sim11$\\,kpc,][]{1997A&A...321..374F} strongly suggest a link between the companion and the absorbing gas, as proposed by \\citet{1997A&A...321..374F} and \\citet{1999A&A...345..414W}.  Furthermore, this is, as far as we are aware, the highest column density so far observed in the spectrum of a BL~Lac object \\citep[after PKS\\,1413+135 which appears to have large absorption in its host galaxy, but the host in that case is an edge-on spiral,][]{2002AJ....124.2401P}, implying that the absorption seen here is not associated with the BL~Lac host.  In this case the $z=0.351$ \\ion{Mg}{ii} belongs to the X-ray absorbing gas which must be poor in dust and metals, implying that the BL~Lac is illuminating a fairly pristine gas cloud that is associated with the galaxy close to the line of sight at redshift $z=0.351$."
    },
    "0402/astro-ph0402663_arXiv.txt": {
        "abstract": "{We report the analysis of an {\\it XMM-Newton} observation of the close binary HD\\,159176 (O7\\,V + O7\\,V). The observed L$_\\mathrm{X}$/L$_\\mathrm{bol}$ ratio reveals an X-ray luminosity exceeding by a factor $\\sim$7 the expected value for X-ray emission from single O-stars, therefore suggesting a wind-wind interaction scenario. EPIC and RGS spectra are fitted consistently with a two temperature {\\tt mekal} optically thin thermal plasma model, with temperatures ranging from $\\sim$2 to 6\\,10$^6$ K. At first sight, these rather low temperatures are consistent with the expectations for a close binary system where the winds collide well before reaching their terminal velocities. We also investigate the variability of the X-ray light curve of HD\\,159176 on various short time scales. No significant variability is found and we conclude that if  hydrodynamical instabilities exist in the wind interaction region of HD\\,159176, they are not sufficient to produce an observable signature in the X-ray emission. Hydrodynamic simulations using wind parameters from the literature reveal some puzzling discrepancies. The most striking one concerns the predicted X-ray luminosity which is one or more orders of magnitude larger than the observed one. A significant reduction of the mass loss rate of the components compared to the values quoted in the literature alleviates the discrepancy but is not sufficient to fully account for the observed luminosity. Because hydrodynamical models are best for the adiabatic case whereas the colliding winds in HD\\,159176 are most likely highly radiative, a totally new approach has been envisaged, using a geometrical steady-state colliding wind model suitable for the case of radiative winds. This model successfully reproduces the spectral shape of the EPIC spectrum, but further developments are still needed to alleviate the disagreement between theoretical and observed X-ray luminosities. ",
        "introduction": "HD\\,159176 is a relatively bright (m$_v = 5.7$) double-lined spectroscopic binary (Trumpler \\cite{Tr}) in the young open cluster NGC\\,6383. The system has been well studied in the visible and UV wavelengths. Conti et al.\\ (\\cite{CCJ}) derived an orbital solution that was subsequently improved by Seggewiss \\& de Groot (\\cite{SdG}), Lloyd Evans (\\cite{LE}) and most recently by Stickland et al.\\ (\\cite{SKPP}). The binary has an orbital period of 3.367\\,days and consists of two nearly identical O-stars in a circular orbit.  Conti et al.\\ suggested that both stars are O7 stars that have evolved off the main-sequence and nearly fill up their Roche lobes. However, Stickland et al.\\ argued that the stars were probably not evolved. Although the system does not display photometric eclipses, Thomas \\& Pachoulakis (\\cite{TP}) reported ellipsoidal variability with an amplitude of about 0.05\\,mag in the optical and UV wavebands. These light curves were analysed by Pachoulakis (\\cite{Pach}) who inferred radii of both stars of order $0.25\\,a$, where $a$ is the orbital separation. According to these results, the stars do not fill up their critical volume and are thus not very deformed.\\\\  Several observations point towards the existence of a wind interaction process in HD\\,159176. For instance, the optical spectrum of the system displays the so-called Struve-Sahade effect, i.e.\\ the absorption lines of the approaching star appear stronger (Conti et al.\\ \\cite{CCJ}, Seggewiss \\& de Groot \\cite{SdG}, Lloyd Evans \\cite{LE}), although the reverse effect is seen in the UV (Stickland et al.\\ \\cite{SKPP}). Though the origin of this effect is as yet not established, a commonly proposed scenario involves the existence of an interaction process within the binary system (Gies et al.\\ \\cite{gies}).  Analysing UV resonance line profiles of HD\\,159176 as observed with {\\it IUE}, Pachoulakis (\\cite{Pach}) derived a mass-loss rate of about $3\\,10^{-6}$\\,M$_{\\odot}$\\,yr$^{-1}$ for each star (note that this value is a factor five larger than the one derived by Howarth \\& Prinja \\cite{HP}). The stellar winds of the components of HD\\,159176 are therefore probably sufficiently energetic to interact. The interaction region is expected to be located roughly mid-way between the stars and it prevents the wind of each star from deploying into the direction towards the other star. This situation has an impact on the UV resonance line profiles and the UV light curve (Pachoulakis \\cite{Pach}, Pfeiffer et al.\\ \\cite{PPKS}). Pfeiffer et al.\\ analysed the variations of the UV resonance lines and concluded that there exists a source of extra emission between the stars. These authors associated this extra emission with resonant scattering of photospheric light by the material inside a colliding wind region. According to their analysis, this shock region is slightly wrapped around the secondary star.  HD\\,159176 was detected as a rather bright X-ray source with {\\it EINSTEIN} ($0.157 \\pm 0.008$\\,cts\\,s$^{-1}$ with the IPC, Chlebowski et al.\\ \\cite{Chle}) and {\\it ROSAT} ($0.291 \\pm 0.034$\\,cts\\,s$^{-1}$ during the All Sky Survey, Bergh\\\"ofer et al.\\ \\cite{BSC}). Chlebowski \\& Garmany (\\cite{CG}) suggested that the excess X-ray emission observed in many O-type binaries compared to the expected intrinsic contribution of the individual components is produced by the collision of the stellar winds (see also e.g.\\ Stevens et al.\\ \\cite{SBP}). Therefore, it seems likely that at least part of the X-ray flux of HD\\,159176 may originate in the wind interaction region. In contrast with this picture, Pfeiffer et al.\\ (\\cite{PPKS}) suggest that the bulk of the X-ray emission arises primarily from the intrinsic emission of the individual components, rather than from a colliding wind interaction.\\\\  To clarify this situation, we obtained an AO1 {\\it XMM-Newton} observation of HD\\,159176. Since the system is X-ray bright, it is well suited to investigate the X-ray properties of a short-period early-type binary.  \\begin{table}[h] \\caption{\\label{tbl-1a} Relevant parameters of the HD\\,159176 binary system adopted throughout this paper unless otherwise stated. The numbers are taken from Pachoulakis (\\cite{Pach}, P96) and Diplas \\& Savage (\\cite{DS}, DS).} \\begin{center} \\begin{tabular}{l c c c} \\hline Parameter & Prim. & Sec. & Ref. \\\\ \\hline $a\\,\\sin{i}$\\,(R$_{\\odot}$) & \\multicolumn{2}{c}{28.9} & P96 \\\\ $M$\\,(M$_{\\odot}$) & 31.9 & 31.6 & P96 \\\\ $i\\,(^{\\circ})$ & \\multicolumn{2}{c}{$\\sim 50$} & P96 \\\\ $T_{\\rm eff}$\\,(K) & 42500 & 35000 & P96 \\\\ $R_*$\\,(R$_{\\odot}$) & 9.8 & 9.3 & P96 \\\\ $v_{\\infty}$\\,(km\\,s$^{-1}$) & \\multicolumn{2}{c}{2850} & P96 \\\\ ${\\dot{\\it {\\rm M}}}$\\,(M$_{\\odot}$\\,yr$^{-1}$) & $3.2\\,10^{-6}$ & $2.6\\,10^{-6}$ & P96 \\\\ $\\log{\\it {\\rm N_{\\rm H, ISM}}}$\\,(cm$^{-2}$) & \\multicolumn{2}{c}{21.23} & DS\\\\ \\hline \\end{tabular} \\end{center} \\end{table} ",
        "conclusions": "} The main results from our study of the {\\it XMM-Newton} data of HD\\,159176 are the following: \\begin{enumerate} \\item[-] The study of EPIC and RGS data reveals a soft spectrum which is consistently fitted with a two temperature thermal model. The hottest component is at about 0.6 keV for EPIC-pn and RGS, and about 1.0 keV for EPIC-MOS data. The X-ray luminosity is rather low, with a value of L$_\\mathrm{X}$/L$_\\mathrm{bol}$ showing a moderate excess ($\\sim$ 7) compared to what is expected for isolated stars with the same bolometric luminosity but without any wind-wind interaction.\\\\ \\item[-] Our analysis of RGS data reveals that lines are significantly broadened. For instance, the \\ion{O}{viii} Ly $\\alpha$ line at 19.0 \\AA\\, has a full width at half maximum of about 2500 km\\,s$^{-1}$, in agreement with X-ray lines originating either from shocks distributed throughout the wind or from a colliding wind zone.\\\\ \\item[-] The study of EPIC light curves failed to reveal any significant variability on time scales of 100 -- 2000\\,s. This indicates that, at least in a system like HD\\,159176, the hydrodynamic instabilities that might exist in the region of the shocked winds are not able to produce a clear variability of the X-ray emission.\\\\ \\item[-] The EPIC spectrum reveals a high energy tail which can not be fitted by thermal models ({\\tt mekal} or colliding wind). This hard X-ray emission component was fitted with a power law with a photon index of about 2.5. Its presence could reflect a non-thermal process such as inverse Compton scattering (Pollock \\cite{Pol}, Chen \\& White \\cite{CW}). The photon index is not too far from the 1.5 value expected for an X-ray spectrum arising from a population of relativistic electrons generated through an acceleration mechanism involving strong shocks. The possibility of a non-thermal X-ray component was already suggested for instance for a system like WR\\,110 by Skinner et al. (\\cite{Sk}). So far, the most prominent indication of relativistic electrons in stellar winds of early-type stars have been found in the radio domain where a significant fraction of the stars were found to display a non-thermal, probably synchrotron, emission (e.g. Bieging et al.\\,\\cite{BAC}). In the case of HD\\,159176, radio observations failed to reveal such a non-thermal component, providing only an upper limit on the radio flux at 6 cm (Bieging et al.\\,\\cite{BAC}). However, this does not rule out the possibility that relativistic electrons could be accelerated at the wind collision shock. In fact, the acceleration site would be buried so deeply within the radio photosphere that no synchrotron emission could escape and we would therefore observe HD\\,159176 as a thermal radio emitter. \\\\ \\item[-] Besides the non-thermal tail, the observed spectral shape can be consistently reproduced using a steady-state colliding wind model with a mass loss rate value of about 1.7 -- 2.6\\,10$^{-7}$ M$_\\odot$\\,yr$^{-1}$ and a terminal velocity ranging between 1850 and 1950 km\\,s$^{-1}$ (or between 2140 and 2240 km\\,s$^{-1}$ for EPIC-MOS data). However, this model is unable to predict X-ray luminosities compatible with the observed spectrum of HD\\,159176. Theoretical values are systematically higher than the observed X-ray luminosities. This disagreement between theory and observation is discussed hereafter.\\\\ \\end{enumerate}  The disagreement between our {\\it XMM-Newton} observation and the theoretical predictions could possibly be explained by several factors. First, let us recall that the kinetic power of the collision should be considered as an upper limit on the X-ray luminosity even in a highly radiative system. For instance, some of the collision energy might be taken away by the shocked gas which ends up with negligible velocity and pressure near the line of centers. In order not to pile up at the `stagnation point', this gas must be advected from the system, and the work needed to lift it out of the gravitational potential of the system could take away energy at the expense of X-ray emission. Second, higher values of the parameter $\\eta$ (see the unequal wind case, Sect.\\,\\ref{uneq}) should be considered. Although the current version of the model is unable to deal with the case where the primary wind crashes onto the secondary photosphere, this scenario should be envisaged. Let us emphasize however that the optical spectrum of HD\\,159176 does not provide support for very large values of $\\eta$. Third, diffusive mixing between hot and cool material is likely to exist due to the instability of the shock front. As a consequence, the material tends to emit a much softer spectrum (i.e. EUV, or even UV), at the expense of X-rays. Unfortunately, current simulations lack the needed resolution to accurately deal with this mixing. Note that thermal conduction (Myasnikov \\& Zhekov \\cite{MZ}) is also expected to produce a softer spectrum. Finally, further developments of the current model are needed to address this issue. An improvement of the current steady-state model would be for instance to consider the effect of mechanisms able to lower the predicted luminosities like sudden radiative braking (Gayley et al.\\,\\cite{GOC}), which can potentially occur every time unequal winds interact in a close binary system. In addition, orbital effects should also be included to study such systems.\\\\  HD\\,159176 is the first short period colliding O + O binary studied with the high sensitivity of the instruments on board the {\\it XMM-Newton} satellite. The major point of this system is that the shocks associated with the wind collision are radiative, making them very difficult to simulate with current hydrodynamic models. For the first time, an alternative model has been used to address this case, following a steady-state geometrical approach leading to promising results. To achieve a better understanding of the special case of radiative colliding wind shocks, other close binary systems should be observed in conjunction with further developments of the new theoretical approach followed in this study.   \\acknowledgement{Our thanks go to Mathias Ehle ({\\it XMM}-SOC) for his help in processing the EPIC data and to Alain Detal (Li\\`ege) for his help in installing the {\\sc sas}. We wish to thank Andy Pollock (ESA) for discussion. The Li\\`ege team acknowledges support from the Fonds National de la Recherche Scientifique (Belgium) and through the PRODEX XMM-OM and Integral Projects. This research is also supported in part by contracts P4/05 and P5/36 ``P\\^ole d'Attraction Interuniversitaire'' (SSTC-Belgium). JMP gratefully acknowledges funding from PPARC for a PDRA position. IIA acknowledges support from the Russian Foundation for Basic Research (grant 02-02-17524). This research has made use of the SIMBAD database, operated at CDS, Strasbourg, France and of the NASA's Astrophysics Data System Abstract Service.}"
    },
    "0402/astro-ph0402380_arXiv.txt": {
        "abstract": "We study X-ray and variability and distance of \\source. We derive the distance of $\\ga 7$ kpc, based on recent determination of the binary parameters. We study data from the ASM aboard \\ginga, the BATSE aboard \\gro, and the ASM, PCA and HEXTE aboard \\xte. From 1987 to 2004, \\source\\ underwent $\\sim$15 outbursts and went through all known states of black-hole binaries. For the first time, we present the PCA data from the initial hard state of the outburst of 2004. We then study colour-colour and colour-flux correlations. In the hard state, there is a strong anticorrelation between the 1.5--5 and 3--12 keV spectral slopes, which we explain by thermal Comptonization of disc photons. There is also a softening of the spectrum above 3 keV with the increasing flux that becomes stronger with increasing energy up to $\\sim$200 keV. This indicates an anticorrelation between the electron temperature and luminosity, explained by hot accretion models. In addition, we see a variable broad-band slope with a pivot at $\\sim$200 keV. In the soft state, there is a high energy tail with varying amplitude beyond a strong and variable blackbody component. We confirm the presence of pronounced hysteresis, with the hard-to-soft state transitions occurring at much higher (and variable) luminosities than the soft-to-hard transitions. We fit the \\xte/ASM data with a model consisting of an outer accretion disc and a hot inner flow. State transitions are associated then with variations in the disc truncation radius, which we fit as $\\sim 6GM/c^2$ in the soft state and several times that in the hard state.  The disappearence of the inner disc takes place at a lower accretion rate than its initial appearance due to the dependence of the transitions on the source history. We provide further evidence against the X-ray emission in the hard state being nonthermal synchrotron, and explain the observed radio-X-ray correlation by the jet power being correlated with the accretion power. ",
        "introduction": "\\label{s:intro}  The Galactic X-ray binary \\source\\ was discovered $>$30 years ago (Markert et al.\\ 1973), and it has been extensively studied since then. Its optical companion is undetectable and most of the observed optical emission originates in an accretion disc. The system is classified as a low-mass X-ray binary from upper limits on the luminosity of the companion star (e.g., Shahbaz, Fender \\& Charles 2001).  \\source\\ is a transient source exhibiting state transitions between soft, hard and off states. Based on its temporal and spectral characteristics, it is considered to be a black hole binary (Sunyaev \\& Revnivtsev 2000; Zdziarski et al.\\ 1998, hereafter Z98), which is in agreement with an estimate of its mass function $\\ga 2\\msun$ at the 95 per cent confidence (Hynes et al.\\ 2003).  Most of its X$\\gamma$-ray spectra in luminous states are remarkably similar to those of Cyg X-1, and well fitted by Comptonization (Z98; Wardzi\\'nski et al.\\ 2002, hereafter W02). However, while the two sources share common properties on time scales $\\la$ a few days, their evolution patterns are rather different on longer time scales (e.g., Smith, Heindl \\& Swank 2002b). In particular, the spectral evolution of Cyg X-1 is directly correlated with the luminosity changes, whereas the spectral changes in \\source\\ have been observed to lag the variations of the luminosity. These lags constitute a hysteresis in the lightcurve of the source; the hard-to-soft state transitions during the rise phase occur at higher luminosities than the soft-to-hard ones during the decline phase.  A major difference between \\source\\ and Cyg X-1 is that the latter is a persistent source in a high-mass X-ray binary. Thus, accretion in Cyg X-1 is via a focused wind from the high-mass companion, while in \\source\\ it is due to Roche-lobe overflow from the low-mass companion. The different behaviour of low and high-mass black-hole binaries is likely to be due to different disc sizes (e.g., Smith et al.\\ 2002b). In a low-mass X-ray binary, the disc is expected to be large. This then leads to delays in the propagation of the accretion rate from the disc outer regions to the hot inner flow. Another consequence of the different types of accretion flows in the two systems is that whereas Cyg X-1 is a persistent source, in which the bolometric luminosity (and presumably the accretion rate) varies only by $\\sim$5 or so (Zdziarski et al.\\ 2002, hereafter Z02), \\source\\ is a transient, in which accretion is unstable and the accretion rate in the vicinity of the black hole varies by at least several orders of magnitude. Consequently, \\source\\ goes through a wider range of spectral states than Cyg X-1.  In this paper, we first estimate the distance to the source based on recent optical/IR data. Then we study the long-term variability of \\source\\ during 1987--2004 and its theoretical interpretation. We use data from the All-Sky Monitors (ASM) on board \\ginga\\/ and {\\it Rossi X-ray Timing Explorer (RXTE)\\/} satellites, the Burst and Transient Source Experiment (BATSE) on board {\\it Compton Gamma Ray Observatory (CGRO)}, and the currrently available data from the Proportional Counter Array (PCA) and High-Energy Transient X-ray Experiment (HEXTE) on board \\xte. ",
        "conclusions": "We have found that the current optical/IR data imply that the distance to \\source\\ is $\\ga 7$ kpc. Its most likely location is in the Galactic bulge. We also argue against the recently proposed distance of $\\sim$15 kpc, in which case \\source\\ would be a halo object.  Then, we have performed a comprehensive analysis of 16 years of the long-term variability of \\source. We found the long-term $\\langle L\\rangle$ from this transient source has increased about threefold since 1991, corresponding to marked changes in the character of the outbursts. The epoch MJD 48400--49500 was characterized by a $\\sim$500 d cycle of strong hard outbursts separated by deep quiescence and $\\langle L\\rangle\\sim 0.015\\ledd (d/8\\,{\\rm kpc})^2 (10\\msun/M)$. The epoch MJD 49500--50800 was characterized by a persistent hard state and $\\langle L\\rangle$ about 1.5 times the above value. This increase lead then to a long soft state, a $\\sim$1000 d queiescence, and the most recent $\\sim$400 d very strong outburst. During this epoch, MJD 50800--53000, $\\langle L\\rangle\\sim 0.05\\ledd (d/8\\,{\\rm kpc})^2 (10\\msun/M)$.  We have studied in detail hysteresis in this source. Especially interesting is the strong dependence of the flux of the hard-to-soft transition on the preceding behaviour of the hard/off state. The 1998 transition, which followed the long persistent hard-state period, took place at the flux about three times lower than the 2002 one, which followed the long queiescent epoch. We model the hysteretic variations by independent variability of the accretion rate and the inner radius of the cold disc. The maximum possible $\\dot M$ in the hard state is higher than the minimum possible one in the soft state, but transitions from the hard to soft state can occur below that maximum.  We have found a new X-ray spectral correlation in \\source, between the relative strength of the soft excess (due to blackbody disc emision) in the hard state and the hardness of the main power law component. The effect is explained by cooling of the hot Comptonizing plasma by a variable fraction of the disc emission, probably related to a variable overlap of the disc by the hot flow. Interestingly, the correlation forms an upper horizontal branch in the colour-colour diagram, previously thought to be specific to neutron-star binaries. In contrast to Cyg X-1, we do not find any orbital modulation of the X-ray flux.  We have also found a pivoting pattern of hard-state X-ray variability, with the pivot energy at $\\sim$200 keV, and confirmed the existence of an anticorrelation between the bolometric flux and the high-energy cutoff, presumably related to the temperature of the Comptonizing electrons. The weak correlation of the X-ray slope and the flux (in contrast to the case in Cyg X-1) is explained by the wide range of the fluxes at which the the hard state appears, related to the transient nature of \\source.  Our results provide strong support to the thermal-Comptonization nature of X-rays in the hard state. We explain that the correlation between the X-ray and radio fluxes by the total power of the accretion emission being correlated with the power in outflowing particles forming the radio jet, which is, in turn, correlated with the radio flux."
    },
    "0402/astro-ph0402349_arXiv.txt": {
        "abstract": "{The identification of black holes is one of the most important tasks of modern astrophysics. Candidates have been selected among binary stars based on a high mass function, and seriously considered when the lower mass limit exceeds $\\sim \\, 3 \\, M_{\\odot}$. More recently the absence of (Type I) thermonuclear bursts has been advanced as an additional criterion in favor of the black hole interpretation, since the absence of a solid surface naturally precludes the accumulation and ignition of accreting material. We discuss in this {\\it Letter} the possibility that self-bound stars made of CFL-paired quarks mimic the behavior of at least the low-mass end black holes as a result of: a) higher maximum masses than ordinary neutron stars, b) low steady luminosities due to the bare surface properties, and c) impossibility of generating Type I bursts because of the complete absence of normal matter crusts at their surfaces. These features caution against a positive identification of event horizons based on the lack of bursts. ",
        "introduction": "The strong gravitational field regime of General Relativity is widely believed to be physically realized in black holes, uniquely possessing event horizons. The positive identification of black holes among existing astrophysical objects has the potential of providing a glimpse into that regime which captures the imagination of scientists and public alike. This is why the study of black hole candidates has consumed so much time of research efforts since the early days of X-ray astronomy, when the first strong sources were identified and some of them tentatively associated with astrophysical black holes.   Unfortunately, the task of a positive identification is still plagued with caveats, and even though the advances in the last decade or so have been quite impressive (\\cite{Orosz}), the burden of proof is still with observational astrophysicists. Meanwhile, the discussion continues and new proposals arise as potential unique signatures of event horizons possessed by black holes only. In a recent series of papers (\\cite{Nar1,Nar2} and references therein), Narayan and coworkers have proposed a class of X-ray binaries (the soft-X transients, or SXT) as attractive candidates to black holes. The argument is that those sources show, in addition to a mass function $f(M)$ larger than the expected $\\sim \\, 3 \\, M_{\\odot}$ maximum limit for NS models (the mass function is an absolute lower limit to the mass of the compact object in the SXT), an absence of type I bursts tentatively interpreted as evidence for an event horizon. In fact, a systematic study undertaken (\\cite{Nar2}) to pin down the physical condition for type I thermonuclear bursts has shown that SXT sources should burst if they possess a normal matter crust, therefore the absence of bursts could be interpreted as indicating the presence of event horizons. In a recent paper, Yuan, Narayan and Rees (2004) made a general attempt to show that {\\it any} neutron star (composed by matter described by a more or less general equation of state) should experience thermonuclear Type I bursts at appropriate mass accretion rates. The question is whether an ``abnormal'' surface also allows such physical behavior. We shall argue below that general models do not necessarily possess the anomalous high-density surface properties of self-bound quark stars, which is one of the key ingredients that entangle their discrimination within the black hole candidates.   Since the proof of the existence of event horizons is so important for modern astrophysics, careful examinations of the possible loopholes and alternatives to the signatures are needed. As an example of the latter, a critical analysis by Abramowicz, Klu\\'zniak and Lasota (2003) concluded, based on a series of persuasive arguments, that a positive proof of the event horizon will be impossible in principle. Specifically, they argue that the case of SXTs as black holes based on the absence of type I bursts would be weakened by the finding of any kind of exotic star without a ``normal'' crust. We have reexamined this objection with the aim of addressing the recently proposed models of self-bound CFL stars and report our findings below. ",
        "conclusions": "The photon emission properties of CFL strange stars are expected to resemble those of bare SS (\\cite{ISPE}). Since pairing effects should appear in the plasma frequency $\\omega_p$ through the baryon number density as a correction of order  $\\mu \\Delta^2$, the plasma frequency $\\omega_p$ will not be very different from (typical) $20$ MeV of unpaired quark matter, and thus the equilibrium photon radiation will be suppressed. A tiny luminosity makes CFL strange stars very difficult to detect directly. The thermal emission of photons from the bare quark surface of an (unpaired) strange star (due mainly by electron-positron pair production) shown to be much higher than the Eddington limit by Page \\& Usov (2002) may not operate for CFL strange stars since no electrons will be present at the surface and the emissivities due to other processes are negligible at low temperatures.  These features  (high maximum masses,  absence of normal matter crusts and lack of surface emission) show that ``exotic'' stellar models may be constructed in which Type I thermonuclear bursts can not occur, but which are not black holes either. If physically realized in nature, some of the SXT systems observed to possess a relatively low mass function (e.g. SXT A0620-00, with $f(M) \\geq 3 \\, M_{\\odot}$; \\cite{McR}) may harbor self-bound CFL stars.  In these models, the lack of Type I thermonuclear bursts is not interpreted as a signature of an event horizon, but rather as a consequence of the impossibility of a normal crust with $\\rho \\leq 10^{6} g cm^{-3}$ where accreted matter could accumulate and eventually ignite. Further work is needed to elaborate or rule out this tentative association, but it is already clear that CFL stars provide a definite counterexample of the event horizon proof as a representative of a class of exotic alternatives not easily discarded."
    },
    "0402/astro-ph0402513_arXiv.txt": {
        "abstract": "The recent observation (Ferraro et al. 2003b) of the blue straggler population in 47 Tucanae  gives the first detailed characterization of their spatial distribution in the cluster over its entire volume. Relative to the light distribution, blue stragglers  appear to be overabundant in the core and at large radii. The observed surface density profile shows a central peak, a zone of avoidance and a rise beyond twenty core radii. In light of these findings we explored the evolution of blue stragglers mimicking their dynamics in a multi-mass King model for 47 Tucanae. We find that the observed spatial distribution  can not be explained within a purely collisional scenario in which  blue stragglers are generated exclusively in the core through direct mergers. An excellent fit is obtained if we require  that a sizable fraction of blue stragglers is generated in the peripheral regions of the cluster inside primordial binaries that evolve in isolation experiencing mass-transfer. ",
        "introduction": "Blue Stragglers (BSs), first discovered by Sandage in the globular cluster M3,  are stars lying above and blue-ward of the turn-off point in a cluster color-magnitude diagram.  In recent years the high angular resolution and the UV imaging capabilities of the {\\it Hubble Space Telescope} (HST) have made possible the search of BSs even in the cores of highly concentrated globular clusters (GCs). Observations indicate that BSs in GCs are preferentially found in the core (Ferraro et al. 1999a); but in at least two GCs, M3 (Ferraro et al. 1993; Ferraro et al. 1997) and M55 (Zaggia et al. 1997) BSs are seen also in the external region of the GC. Ferraro et al. (2003b) have recently found that the radial distribution of BSs in 47~Tuc also appears bimodal, i.e. highly peaked in the core, decreasing at intermediate radii and rising  again at larger radii. A long standing problem is the formation mechanism of BSs. Many scenarios have been proposed to explain the BS origin (Fusi Pecci et al. 1992; Bailyn 1995; Bailyn and Pinsonneault 1995; Procter Sills, Bailyn \\& Demarque 1995; Sills \\&{} Bailyn 1999; Sills et al. 2000; Hurley et al. 2001) and two seem to be the most likely ones. The first, i.e., the {\\it collisional scenario}, indicates that BSs are the end-product of a prompt merger between two main sequence stars in a direct collision that involves a (resonant) three or four body encounter (Davies, Benz \\& Hills 1994; Lombardi et al. 2002); these BSs acquire kicks generated by dynamical recoil.  The second, or {\\it mass-transfer scenario}, suggests that BSs are generated in  primordial binaries (hereafter PB) that evolve mainly in isolation or harden gently by long-distance gravitational encounters until they reach contact, leading to (unstable) mass-transfer and final coalescence (Carney et al. 2001). In both these mechanisms BSs are formed with a mass exceeding the turn-off mass of the cluster and can stay on the main sequence through the mixing of the hydrogen-rich surface layers of its two progenitor stars. These two scenarios do not necessarily exclude each other and may coexist (Leonard 1989; Fusi Pecci et al. 1992; Bailyn \\& Pinsonneault 1995; Ferraro et al. 1997; Sills \\& Bailyn 1999; Hurley et al. 2001). Indeed, bimodal BS radial distributions seem to invite the invocation of two mechanisms: the BSs in the core are principally created by star collisions, while external BSs are formed from mass-transfer in PBs. Still Sigurdsson, Davies \\& Bolte (1994) suggested that the bimodal population in M3 might be entirely explained within a collisional model where external BSs are formed in the core and ejected into the outer regions by recoil. However, further studies of the BS luminosity function and their comparison with theoretical models (Bailyn \\& Pinsonneault 1995; Sills \\& Bailyn 1999; Sills et al. 2000; Ferraro et al. 2003b), highlight the difficulty, if not the impossibility, to obtain a good fit of the observations assuming that all the BSs are formed dynamically in the core. In this paper we study, in the light of the most recent data (Ferraro et al. 2003b), the case of 47~Tuc, comparing the observed  bimodal distribution of BSs with a series of simulations carried on using a new version of the dynamical code described in Sigurdsson \\& Phinney~(1995). ",
        "conclusions": "The comparison of the observed BS distribution in 47~Tuc with the simulated distributions supports  the hypothesis that internal BSs principally result from stellar collisions, while external BSs (outside of 20 $r_c$) are exclusively generated by mass-transfer in PBs. The best fit to the observational data of 47~Tuc is obtained when a sizable fraction (25\\%) of BSs is generated  from PBs in peripheral regions ([30,60] $r_c$). The internal BSs contributing up to 75\\% of the observed are all born inside 0.5 $r_c$ with a natal kick of 1 $\\sigma{}$, and do not pollute the external regions.  External BSs contribute little to the core population. A scenario in which all BSs are generated over the entire cluster by mass-transfer in PBs can not be ruled out. We believe however that a blending between collision induced evolution and internal evolution is at play in the GC core to explain internal BSs. Our main finding is the need of a population of external PBs in order to generate the bimodal distribution of the BSs observed in 47~Tuc. The required fraction (10\\%) to fit the data does not seem unreasonable. More accurate counts of BSs in GCs with widely different properties are about to be collected  from high resolution photometry (Ferraro et al. 2004 in preparation). These observations may shed light into the nature of BSs and the importance of PBs in GC evolution. In light of these upcoming observations, in paper II (Mapelli et al. in preparation),  we plan to continue our analysis with  our simulations over a wide sample of GCs. Theoretical studies using N-body (Baumgardt \\& Makino 2003) or Monte Carlo techniques (Ivanova et al. 2004) will eventually become necessary tool for exploring the formation and evolution of BSs in GCs."
    },
    "0402/astro-ph0402039_arXiv.txt": {
        "abstract": "{ L1521E seems unique among starless cores. It stands out in a distribution of a ratio $(R)$ that we define to asses core evolution, and which compares the emission of the easily-depleted C$^{18}$O molecule with that of the hard to deplete, late-time species N$_2$H$^+$. While all cores we have studied so far have $R$ ratio lower than 1, L1521E has an $R$ value of 3.4, which is 8 times the mean of the other cores. To understand this difference, we have modeled the C$^{18}$O and N$_2$H$^+$ abundance profiles in L1521E using a density distribution derived from 1.2mm continuum data. Our model shows that the C$^{18}$O emission in this core is consistent with constant abundance, and this makes L1521E the first core with no C$^{18}$O depletion. Our model also derives an unusually low N$_2$H$^+$ abundance. These two chemical peculiarities suggest that L1521E has contracted to its present density very recently, and it is therefore an extremely young starless core. Comparing our derived abundances with a chemical model, we estimate a tentative age of $\\le 1.5 \\times 10^5$ yr, which is too short for ambipolar diffusion models. ",
        "introduction": "Observations of starless dense cores over the last few years have shown that molecular depletion onto cold dust grains is a common feature in low-mass star-forming regions (e.g., \\citealt{cas99,taf02}). As a core contracts, the amount of depletion of certain molecules like CS and CO increases with time (e.g., \\citealt{ber97,aik03}), so we can expect that depletion will provide an accurate clock for core evolution when its details are well understood. Even before this is a reality, depletion can be used as a qualitative time marker, in the sense that evolved cores should be more depleted of certain species than younger cores. When we attempt to build a sequence of cores having different amounts of depletion, however, we find that the population of undepleted (or lightly) depleted cores is missing. These undepleted cores should represent the earliest stages of core contraction, and probably they have not been identified so far because of an observational bias. Given the interest of young cores for studies of core contraction, finding a member of this population is an urgent challenge.  To identify cores with low degree of depletion, we have started a systematic study of cores with weak NH$_3$ lines in the survey of \\citet{ben89}, as previous  observations of cores bright in NH$_3$ have only found cases of strong depletion (e.g., \\citealt{cas99,taf02}). Among our targets, the L1521E core in Taurus has revealed itself as a core with negligible depletion, and here we report a preliminary analysis of its properties. This core has already been identified as a very young core by \\citet{hir02}, who found it very prominent in carbon-chain molecules, comparable to TMC-1 (see also \\citealt{aik03}). In this paper we show that L1521E has the lowest measured level of C$^{18}$O depletion, and in a forthcoming article we will present the results of an ongoing chemical survey of this core. ",
        "conclusions": ""
    },
    "0402/hep-ph0402282_arXiv.txt": {
        "abstract": "It is shown that in a smooth hybrid inflation model in supergravity adiabatic fluctuations with a running spectral index with $\\ns >1$ on a large scale and $\\ns <1$ on a smaller scale can be naturally generated, as favored by the first-year data of WMAP. It is due to the balance between the nonrenormalizable term in the superpotential and the supergravity effect. However, since smooth hybrid inflation does not last long enough to reproduce the central value of observation, we invoke new inflation after the first inflation.  Its initial condition is set dynamically during smooth hybrid inflation and the spectrum of fluctuations generated in this regime can have an appropriate shape to realize early star formation as found by WMAP.  Hence two new features of WMAP observations are theoretically explained in a unified manner. ",
        "introduction": "\\label{sec:intro}  The observation of cosmic microwave background anisotropy by the Wilkinson Microwave Anisotropy Probe (WMAP) has determined cosmological parameters with a good accuracy. They have shown that the geometry of our universe is practically flat and that the energy density of the universe is dominated by dark energy and compensated by dark matter and a small amount of baryons \\cite{Bennett,Spergel}. Furthermore, primordial density fluctuations are shown to be adiabatic, Gaussian, and nearly scale invariant, which suggests that they were produced during inflation \\cite{Spergel,Peiris}. Thus, the so-called concordance model was confirmed.  Going into the details, however, we find several interesting features that may not be reconciled with a simple scale-invariant spectrum : namely, an early period of re-ionization (see, e.g., \\cite{YY} for a possible explanation), a lack of fluctuation power on the largest scale,\\footnote{This feature was already seen in COBE observations and a possible explanation was proposed before WMAP data were released \\cite{JY99}.} the running spectral index of fluctuations, $\\ns$ \\cite{YY,FLZZ,KS,KYY}, and the oscillatory behavior of multipoles which may suggest oscillations in the primordial power spectrum \\cite{kogo}. Of course, these properties may disappear eventually when the observations are improved. But it is still important to consider a model to explain them at the present stage.  In this paper, we discuss the running of spectral index from $\\ns>1$ on a large scale to $\\ns<1$ on a smaller scale. More concretely, it is shown that $\\ns=1.13\\pm 0.08$ and $d\\ns/d\\ln k=-0.055^{+0.028}_{-0.029}$ on the scale $k_0=0.002\\,\\Mpc^{-1}$ \\cite{Peiris}.  Among the three types of slow-roll inflation--namely, new \\cite{newinf}, chaotic \\cite{chaoinf}, and hybrid \\cite{hybrid} inflation--the first two scenarios predict fluctuations with $\\ns < 1$ while hybrid inflation can realize both those with $\\ns >1$ and $\\ns < 1$.  Although it is fairly easy to construct a model whose spectral index runs from $\\ns <1$ to $\\ns > 1$ for decreasing length scales, it is quite nontrivial to realize the opposite running.  The hybrid inflation model in supergravity proposed by Linde and Riotto \\cite{LR} is an exception in which the desired feature is realized due to the contributions to the potential from both one-loop effects and supergravity effects.  Based on this observation, some models have been discussed in an attempt to reproduce the results of the WMAP, but it turned out that the large enough running is incompatible with long enough inflation \\cite{KS,KYY,YY}.  This is because the Yukawa coupling constant must be relatively small for sufficient inflation while it must be large for large running.  This problem was first solved in \\cite{KYY} by introducing another inflaton whose appropriate initial condition is automatically prepared during the hybrid inflation regime \\cite{IY,Kawasaki:1998vx,Kanazawa:1999ag}.  This model, however, could not reproduce the central value of the running spectral index obtained by WMAP data but could realize the feature only at the one-sigma level because small-scale fluctuations tend to be too large and we have to hide the corresponding scales in an unobservable region.  More serious is the problem of the initial condition common to other hybrid inflation models that only a very limited initial configuration can lead to inflation \\cite{hybridinit}.  Both these problems have been solved in the chaotic hybrid new inflation model in supergravity proposed by us \\cite{YY}, which can also realize mildly large fluctuations in the appropriate scales to realize early star formation to help early re-ionization.  In the present paper we present another possible mechanism to realize running spectral index from $\\ns >1$ to $\\ns < 1$ for decreasing length scale: namely, smooth hybrid inflation in supergravity.  This scenario was originally proposed by Lazarides and Panagiotakopoulos \\cite{smooth,shift,SS}, in which nonrenormalizable terms are introduced and gauge symmetry remains broken even during hybrid inflation. Thus, topological defects are not produced at the end of inflation. In this paper, we discuss smooth hybrid inflation in supergravity and investigate whether the running spectral index is obtained with the desired property.  As will be shown shortly, the spectral index runs from $\\ns>1$ on a large scale to $\\ns<1$ on a smaller scale without resorting to the one-loop effects unlike our previous models \\cite{KYY,YY}.  Another merit of this scenario is that it can be realized with natural initial conditions in minimal supergravity \\cite{smoothini}. We find, however, that we cannot yield large enough $e$-folds of inflation with large enough running of the spectral index whichever power of the nonrenormalizable term we may choose.  So another inflation is required after smooth hybrid inflation as with the case with our previous models \\cite{YY,KYY}.  Adopting new inflation as the second inflation, we make a specific model of smooth hybrid new inflation in supergravity.  Generally speaking, density fluctuations produced at the onset of new inflation become large.  Actually, as shown in \\cite{KYY}, if we consider usual hybrid inflation before new inflation, the density fluctuations produced during new inflation are too large, which may cause an overproduction problem of dark halos.  However, in the case of smooth hybrid inflation, they can be adequately large, which may be helpful for early star formation.  Thus this scenario can also solve the two problems of the hybrid new inflation model of \\cite{KYY} just as the chaotic hybrid new inflation model \\cite{YY} does.  Which of these two remaining models is more appropriate may be judged by future observations from the presence or the absence of cosmic strings, because the latter model predicts cosmic strings with their energy scale close to the current observational upper bound imposed by cosmic microwave background radiation.  On the other hand, it has been claimed that long cosmic strings lose their energy directly into particles instead of string loops \\cite{Vincent:1996rb}.  Although we understand this issue is still in dispute \\cite{moore}, if it turned out to be true, it would rule out our previous model because too many high-energy cosmic rays would be produced \\cite{Wichoski:1998kh}.  The rest of the paper is organized as follows. In Sec. II we consider smooth hybrid inflation in supergravity and investigate the spectral nature of produced density fluctuations. Then in Sec. III, after reviewing new inflation, we introduce smooth hybrid new inflation, and investigate their dynamics and density fluctuations. Section IV is devoted to a discussion and future outlook. In this paper, we set $\\mg=2.4\\times 10^{18}$ GeV to be unity otherwise stated. ",
        "conclusions": "In this paper we proposed a new model of inflation in supergravity, in which the two new features discovered by the recent precision measurements of cosmic microwave background anisotropy can be explained simultaneously and naturally: namely, the running of spectral index of density fluctuations on large scale as preferred by the first-year WMAP data and a large enough amplitude of fluctuation on small scale relevant to first star formation to realize early re-ionization as discovered by WMAP.  The desired running feature--that is, the spectral index with $\\ns >1$ on a large scale and $\\ns <1$ on a smaller scale--is naturally generated by the balance between the nonrenormalizable term in the superpotential and supergravity effects without resorting to the one-loop effect contrary to our previous models \\cite{YY,KYY}.  Compared with the chaotic hybrid new inflation model in supergravity \\cite{YY}, the present model has somewhat simpler symmetry structures, although we have been unable to explain the hierarchy of the energy scales of two inflations here unlike in the previous model \\cite{YY}. Because our previous model induces string formation with a fairly large energy scale, if forthcoming analysis could rule out such topological defects, the present model would be the only surviving model among the two."
    },
    "0402/astro-ph0402062_arXiv.txt": {
        "abstract": "E+A galaxies, whose spectra have deep Balmer absorption lines but no significant [OII] emission, are the best candidates for an evolutionary link between star-forming, gas-rich galaxies and quiescent, gas-poor galaxies. Yet their current \\emph{morphologies} are not well known. We present \\emph{HST}/WFPC2 observations of the five bluest E+A galaxies ($z\\sim0.1$) in the Zabludoff et al. sample to study whether their detailed morphologies are consistent with late-to-early type evolution and to determine what drives that evolution. The morphologies of four galaxies are disturbed, indicating that a galaxy-galaxy merger is at least one mechanism that leads to the E+A phase.  Two-dimensional image fitting shows that the E+As are generally bulge-dominated systems, even though at least two E+As may have underlying disks. In the Fundamental Plane, E+As stand apart from the E/S0s mainly due to their high effective surface brightness. Fading of the young stellar population and the corresponding increase in their effective radii will cause these galaxies to migrate toward the locus of E/S0s. E+As have profiles qualitatively like those of normal power-law early-type galaxies, but have higher surface brightnesses. This result provides the first direct evidence supporting the hypothesis that power-law ellipticals form via gas-rich mergers. In total, at least four E+As are morphologically consistent with early-type galaxies.  We detect compact sources, possibly young star clusters, associated with the galaxies. These sources are much brighter ($M_R\\sim-13$) than Galactic globular clusters, have luminosities consistent with the brightest clusters in nearby starburst galaxies, and have blue colors consistent with the ages estimated from the E+A galaxy spectra (several $10^8$ yr). Further study of such young star cluster candidates might provide the elusive chronometer needed to break the age/burst-strength degeneracy for these post-merger galaxies. ",
        "introduction": "If galaxies evolve morphologically from late to early types, then some may be now changing from star-forming, gas-rich, disk-dominated objects into quiescent, gas-poor spheroidals.  Spectroscopic surveys have identified at least one set of candidates for such a transformation: ``E+A'' galaxies\\footnote{ Because their spectra are a superposition of a young stellar population (represented by A stars) and an old population (characterized by K stars), these galaxies became known as ``E (for elliptical) + A'' galaxies \\citep{Dressler} or, more straightforwardly, ``K+A'' or ``k+a'' galaxies \\citep{Franx,Dressler99,Poggianti}.}, whose spectra have deep Balmer absorption lines but no significant [OII] emission, indicating that star formation ceased abruptly in these galaxies within the last $\\sim$ Gyr.  In general, E+A galaxies lack significant amounts of HI gas \\citep{chang} and have hot, pressure-supported kinematics \\citep{norton}, suggesting that these galaxies are indeed evolving --- somehow --- from late to early types.  However, we do not yet know whether their current {\\it morphologies} are consistent with late-to-early type evolution or what drives E+A evolution.  While the mechanism (or mechanisms) that causes galaxies to pass through an E+A phase is not understood, there are several clues. First, E+A spectra suggest a recent burst of star formation that required the rapid consumption or dispersal of a gas reservoir.  Second, although they were first studied in distant clusters \\citep{Dressler}, E+As --- at least at low redshifts ($z\\sim0.1$) --- lie mostly in low density environments \\citep{zab96,hogg}.  Third, in low-resolution POSS images, some E+As have features suggestive of tidal tails \\citep{zab96}.  Could E+As be the result of disk galaxy mergers, which are both common in the field and known to enhance star formation?  In the merger hypothesis, E+As are further along the ``Toomre sequence'' \\citep{toomre} and thus more relaxed than systems like the Antennae, whose morphology and kinematics are in such disarray that it is nearly impossible to constrain its endproduct.  E+As may thus teach us considerably more about the endpoints of galaxy-galaxy mergers.  We cannot test this picture of E+A formation, or whether the E+A phase is a bona fide late-to-early type transition, without detailed morphological information.  Simulations predict that well-evolved major mergers have a hybrid morphology, including fading, low surface brightness tidal tails at large radii, a more relaxed spheroid-dominated core, and a population of young star clusters \\citep{Barnes,Barnes_Hernquist,Ashman_Zepf,Mihos}. Identifying such low surface brightness or small scale features, even at low redshifts, requires spatial resolution on the order of 100 pc and low sky background levels.  Therefore, Hubble Space Telescope imaging of nearby E+As is required.   In this paper, we present the detailed \\emph{HST}/WFPC2 morphologies of the five bluest E+A galaxies in the \\citet{zab96} sample.  We review the sample and the data reduction methods in \\S\\ref{sec:reduction}. We describe the qualitative morphologies of these galaxies in \\S\\ref{sec:impression}, discussing the observed tidal features and the implications for E+A origin.  We address the question of whether E+As are consistent with evolution into early types by fitting two-dimensional, surface brightness models to each image and deriving structural parameters such as bulge-to-disk ratio, effective radius, and central surface density (\\S\\ref{sec:fit}). In \\S\\ref{sec:color}, we examine the color gradients in the E+As and compare them with the expectations from disk merger models. We compare the results with the fundamental plane for early type galaxies and with the surface brightness profiles of the nearby elliptical galaxies in \\S\\ref{sec:FP} and \\S\\ref{sec:nuker}, respectively. In \\S\\ref{sec:compact}, we search for star clusters in the E+As and ask whether their properties are consistent with late-to-early type galaxy evolution. We discuss the implications of our results for higher redshift galaxy surveys in \\S\\ref{sec:highz}, cautioning that bulge-to-disk decompositions, quantitative measures of asymmetry, and tests to uncover tidal features may mislead.  Section \\ref{summary} summarizes our results. ",
        "conclusions": "The \\emph{HST} images provide a wealth of information on the small and large scale structure of these galaxies. With the goal of understanding the origin of the E+A phenomenon and into what these systems will evolve, we investigate the morphologies of these systems, their color profiles, their location on the Fundamental Plane \\citep{Jor96} of elliptical galaxies, and their relationship to ``core'' and ``power-law'' ellipticals \\citep{f97}. We also discover a population of associated point sources (possibly young star clusters). Finally, we review the implications of our results, obtained for low-redshift E+As, for the identification and study of such systems at higher redshifts. The reader is referred to Tables 2-5 for a summary of the quantitative results discussed in this section.  \\subsection{Morphologies: First Impressions} \\label{sec:impression} Figure \\ref{fig:tile} shows the WFPC2 mosaic and PC images of our five E+A galaxies at different contrast levels. The full mosaic images ($80''\\times80''$) are in the left column. The center of each E+A is located in the PC, which is in the upper right corner of each mosaic. Tidal features that extend into the other CCDs are evident in EA1-3. The middle and right columns contain the F702W ($24''\\times24''$) and F439W ($12''\\times12''$) PC images, respectively, on a logarithmic flux scale. \\emph{HST}/WFPC2 observations are relatively insensitive in the bluer band so that the signal in the F439W images typically extends out only to $\\sim3$ kpc, 25\\% of the red coverage, and even there it is of low signal-to-noise.  These five E+As exhibit a variety of morphologies ranging from a highly complex system (EA1) to what could visually be classified as a barred S0 galaxy (EA5), even though they have been uniformly selected using spectroscopic criteria, i.e., ``k+a'' type spectra from the LCRS.  EA1 stands apart from the other four E+As. It is composed of two components that are separated spatially by $\\sim$ 3 kpc and another companion with a projected separation of  14 kpc (assuming the companion is at the redshift of EA1). The association is supported by an asymmetric feature emanating from the companion that could be tidal material and a faint bridge that appears to connect it to EA1.  EA2 and EA3 also exhibit highly disturbed morphologies, although EA3 could be visually classified as a normal face-on spiral galaxy in the low contrast PC image. This ambiguity in visual classification is discussed in more detail in \\S\\ref{sec:highz}. EA2 has tidal tail that extends to at least 50 kpc.  EA4 and EA5 appear less disturbed, although EA4 has somewhat irregular outer isophotes, some lopsidedness (in the F439W filter image), and shell-like structures closer to the center that are visible in the PC image. The mechanism or mechanisms responsible for the spectral E+A phenomenon produce a variety of morphologies. Whether all of these systems will evolve into a somewhat more homogeneous population --- for example, early-type galaxies --- is yet unclear.   \\subsection{Morphologies : Bulge-Disk Decompositions} \\label{sec:fit} While EA2-5 appear to have significant spheroidal components, EA3 and EA4, at least, also seem to have a flattened, or perhaps disk-like, morphology. Understanding the fate of these systems requires a quantitative estimate of the relative importance of the dynamically hot and cold stellar components.   Measuring the surface brightness profile for asymmetric, disturbed systems is challenging. To mitigate potential systematic problems, we use two different algorithms. First, to obtain photometric parameters, $r_e$ and $\\mu_e$, we use the two-dimensional image fitting algorithm GALFIT \\citep{chien02} designed to extract structural parameters directly from the galaxy image. GALFIT assumes a two-dimensional model profile for the galaxy. The functional form of the models we choose to fit include combinations of an $r^{1/4}$-law, a S\\'ersic $r^{1/n}$-law, an exponential disk profile, and  a spatially constant sky background. We fit the following: the $(x,y)$ position of the center, $\\mathrm{M}_{tot}$ (the total magnitude of the component), $r_e$ (the effective radius), $n$ (the S\\'ersic index), $q$ (the axis ratio defined as $b/a$), the major axis position angle, and $c$ (the diskiness/boxiness index, where $c > 0$ indicates boxy). This index $c$ plays the same role as the $\\cos 4\\theta$ Fourier coefficient term used often in isophote analysis \\citep{RZ}. As GALFIT explores parameter space, it convolves the model image with a point-spread function (PSF) and compares it to the data for each parameter set. The model PSFs are generated for each galaxy by the TinyTim \\citep{Krist99} software for the WFPC2. Although convolution is computer intensive, the advantage of the convolution process is that it preserves the noise characteristics of the images and can be applied to low signal-to-noise images.   Because GALFIT begins with a very specific, smooth model, which may be a poor representation of such distorted galaxies, we also measure surface brightness profiles using the IRAF/ELLIPSE algorithm. This approach allows the center, major axis position angle,  and ellipticity of each ellipse to change, but does not enforce a model radial profile. To accurately recover the surface brightness profiles without recourse to {\\it ad hoc} models, we applied 20 iterations of Richardson-Lucy deconvolution \\citep{rich, lucy}. \\citet{l98} showed that the WFPC2 PSF can depress the brightness profile as far out as $0\\asec5$ from the galaxy center.  Richardson-Lucy deconvolution allows the intrinsic brightness profile to be recovered to the few percent level down to $r\\sim0\\asec05,$ with adequate exposure levels ($S/N\\sim50$ in the galaxy center).  With reduced $S/N$ and only 20 deconvolution cycles, the central ($r=0$) point in the profile may remain slightly-depressed, dependent on the (unknown) intrinsic structure of the galaxy center.   Because EA1 is too disturbed to be reasonably modeled by a simple disk+bulge model, we restrict our analysis to EA2-5. For each galaxy, we fit  three different light distributions: $r^{1/4}$ law, $r^{1/n}$ S\\'ersic law, and $r^{1/4}$ + exponential disk law. For EA2, we do not fit the $r^{1/4}$ + exponential disk law model because we might be seeing this galaxy close to edge-on (see the linear residuals in Figure \\ref{fig:residual}), and it is hard for GALFIT to fit an edge-on disk with an extended tail. The structural parameters and the reduced $\\chi_\\nu^2$'s of these three GALFIT models are listed in Tables \\ref{tab:param} and \\ref{tab:decomp}. With the exception of one case, $1 < \\chi_\\nu^2 < 2$. These values of $\\chi_\\nu^2$ are somewhat larger than statistically acceptable, due presumably to the presence of asymmetric components, as can be seen in Figure \\ref{fig:residual}.   Of the three profiles we consider, only the S\\'ersic profile has the flexibility to model either a spheroidal or disk-like system by varying the parameter $n$. Therefore the best-fit value of $n$ can guide our conclusions about the nature of the galaxy. An exponential disk corresponds to a value of $n = 1$, while the classic de Vaucouleurs profile corresponds to $n = 4$. However, the correspondence between disk system, spheroid, and $n$ is not quite this simple --- fitting S\\'ersic profiles to SDSS galaxies, Blanton et al. (2003) show a peak at $n = 1$ corresponding to disky systems, but no peak at $n = 4$. Instead, spheroidal systems show a range of $n$ values. This result is further complicated when one factors in differences in radial ranges fit --- for example, fitting the inner slope of cuspy power-law ellipticals (e.g., Lauer et al. 1995) will give a much higher $n$ value than will fits at larger radii.   With these caveats in mind, we find that a single S\\'ersic profile fit yields $n > 5$ for all our galaxies, demonstrating that the light is dominated by a spheroidal component. Indeed, the high values for $n$ indicate a very high concentration of the light, even more than expected for a classic de Vaucouleurs profile. Such high concentrations are consistent with the idea that central starbursts have raised the central luminosity density (e.g., Mihos \\& Hernquist 1994). For example, in the case of EA4, masking the inner kpc and refitting the Sersic law results in a value of $n = 3.6$, much more typical of a normal elliptical. This is not always the case, however --- in EA3, the high S\\'ersic value persists even when the nucleus is masked out. For EA3 the fitted value ($n = 8.7$) is unusually high compared to normal ellipticals (e.g., Kelson et al. 2000; Graham et al. 2001; Graham 2002).  An additional complication to the interpretation of these fits is that, while the light appears to be dominated by a concentrated spheroid, some of the galaxies appear to contain an additional disk-like component that would affect any dynamical model of a merger and its aftermath.   To determine whether these galaxies do indeed contain a disk component, we also fit models with two components. The resulting radial profiles for EA3 and 4 (Figure \\ref{fig:profile}) and the significant decrease in $\\chi^2_\\nu$ (an improvement in the fit at the 99\\% confidence level) demonstrate that a pure spheroid model is not the preferred model for these two systems. To avoid the degeneracies present in fitting disk and bulge simultaneously, we also fit single S\\'ersic profiles just to the outer parts of the galaxies. Using the effective radii of the bulges calculated from the two-component fit, we mask pixels inside a chosen radius, vary that radius to be $\\sim$ $3-5 r_e$ and refit a single component. We mark the effective radii and disk scale lengths with circles in Figure  \\ref{fig:residual}. For EA3, the best-fit S\\'ersic indicies are $n=$ 2.5, 1.6 and 1.3  for masks corresponding to $3 r_e$, $4 r_e$, and $5 r_e$, respectively. For EA4, we measure  $n=$ 0.93 and 0.86 when we apply $3 r_e$ and  $4 r_e$ masks, respectively. In both of these cases, the S\\'ersic index beyond several $r_e$ is as expected for an exponential disk and the fit spans 4 to 5 disk scale lengths. Although we cannot discriminate tidal material from a possible underlying disk, we conclude that in EA3 and EA4 there is material beyond that described by a spheroid and that it is consistent with an underlying disk. EA3 and EA4 appear to be sufficiently relaxed that no significant dynamical evolution is expected, so they may become S0's. We also hypothesize that their progenitors may have included a disk that was significantly heated but not completely destroyed during an intermediate mass ratio merger (e.g., Naab 2000; Bendo \\& Barnes 2000).   Even though no disturbed tidal structure is apparent in EA5, the modeling is complicated by the presence of a strong bar-like structure. The presence of a bar-like feature suggests an underlying disk. When viewed at the different contrasts in Figure \\ref{fig:ea5}, EA5 is composed of at least three distinct components, a extended light distribution in outer part (axis ratio $q \\sim 0.8$), a compact and elliptical bar structure ($q \\sim 0.4 - 0.5$), and a very bright blue central nucleus. The three-component fit gives the S\\'ersic index $n=1.1$ for the central nucleus, $n=0.5$ (Gaussian) for the bar, and $n=1.5$ for the outer disk-like region. This three-component S\\'ersic profile fit (Figure \\ref{fig:ea5}) is acceptable and suggests the presence of disk. For mask sizes $2.5 r_e$, $3 r_e$ and $3.5 r_e$, the best fit S\\'ersic indicies $n$ are 2.0, 1.8, and 1.8, respectively. However, unlike for EA3 and EA4, we do this fit in a limited region and cannot conclude that EA5 has a distinct exponential disk component. The presence of a bar-like feature also suggests an underlying disk.   For the E+As that may contain a disk component, we calculate a bulge-to-total light ratio (B/T) to quantify the relative importance of the bulge and disk-like components. B/T for EA3 and EA4 is 0.56 and 0.62, respectively. These values are larger than the typical B/T for Sa galaxies (0.45) and comparable to the median for S0's (0.63; Kent 1985). Despite the complications of fitting EA5, various modes of fitting the galaxy produce B/T $\\sim 0.7$. Unless the bulge and any underlying disk-like component fade at dramatically different rates, which is unlikely given the relatively weak large-scale population gradients in these galaxies \\citep{norton}, the descendants of these galaxies must be early type (S0 or E, if the disk-like material is tidal debris that disperses or collapses onto the central component).   Given the asymmetric features, how reliable are these fits? There are several ways to check the results for possible systematic errors. First, we compare the fitted analytical profiles to the radial surface brightness profiles obtained from the isophote fitting procedure. In Figure \\ref{fig:profile}, we plot the radial surface brightness profiles of a chosen model for each galaxy: $r^{1/4}$ profile for EA2 and EA5, $r^{1/4}$ + exponential disk profile for EA3 and EA4, and the profiles obtained from the isophote fitting. The differences between the data (ELLIPSE) and models (GALFIT) range mostly between $\\pm 0.5$ mag/arcsec$^2$, are not global, and reflect local asymmetric components. Second, we examine the residual images obtained by subtracting the smooth and symmetric  models from the data (see Figure \\ref{fig:residual}). In all cases we see evidence for components beyond the bulge + disk model. We then calculate how much light remains in the residual images to quantify the goodness of the fit. The relative asymmetric light --- excess (deficit) --- within a 10$''$ radius is 16(8)\\%, 6(5)\\%, 8(9)\\%, 9(8)\\% of the symmetric model components for EA2-5, respectively. Most (50\\% to 80\\%) of the under(over)subtracted light comes from the central region within 0.5$''$, where the even a small amount of fractional deviation from the data can dominate the residual flux over the outer faint parts. Except for EA2, the global under(over)subtractions are roughly the same and localized fluctuations dominate the residuals, so we conclude that our global fits are reliable. In the case with the most residual light (EA2), 24\\% of the light cannot be explained by a symmetric model and the positive residuals dominate all over the galaxy. Therefore, we cannot exclude the possibility that we are looking right along the interaction plane (the tidal debris are quite linear). The point sources near the E+A bulges in the residual images are discussed in \\S\\ref{sec:compact}.   \\subsection{Morphologies: Asymmetric Components} \\label{sec:asymmetric} So far we have fit symmetric smooth models with moderate success, but have found that asymmetric features are quite common in our sample. Asymmetry, in particular lopsidedness, has been used to measure disturbances in local ``normal'' disk galaxies \\citep{RZ, ZR} and correlates with recent star formation \\citep{ZR, greg}. There are multiple ways in which one can quantify asymmetry, but here we choose to follow what was done for local spirals by \\cite{RZ}. This measurement is based on the azimuthal Fourier decomposition of the surface brightness along elliptical isophotes.  For the two most disk-like of the E+As (EA3 and 4), we calculate the Fourier decomposition of the F702W band surface brightness distribution. We use a grid with 24 azimuthal and 36 radial bins from semi-major axes of 4 to 200 pixels. The center of the azimuthal grid is identified as the brightest central point in the galaxy image. Figure \\ref{fig:fourier} shows the amplitudes of the various first Fourier terms as a function of radius. In field spirals, $A_1 > 0.2$ is identified as strong lopsidedness, found in $\\sim 20$\\% of the cases,  and interpreted as the result of a recent interaction (other explanations, such as halo-induced disk sloshing have since been suggested; e.g., Levine \\& Sparke 1998; Kornreich et al 2002). EA3 has $A_1 \\ll 0.2$  at all radii except in the transition region between the inner spiral arms and the tidal tails at $\\sim 3-6$ kpc and at large radii where the uncertainties are large. Although EA4 has large $A_1$ within 3 kpc, which is consistent with the appearance of the residual image (Figure \\ref{fig:residual}), $A_1 < 0.2 $ for $1.5 < r_d < 2.5$, the range used in the study of field spirals. The lopsidedness of EA3 and EA4 is consistent with that of normal spiral galaxies despite the cataclysmic event that occurred $\\sim$ Gyr ago indicated by the spectra. This result has two possible interpretations:  either these galaxies have had sufficient time to relax and ``smooth out'' interaction-induced asymmetries (e.g., Mihos 1995), or they are the result of encounters not strong enough to cause major dynamical damage. Of course, the latter possibility runs into the problem of how to trigger such a massive starburst without dynamically disturbing the galaxy.  \\subsection{Color Gradients} \\label{sec:color} The color gradients of E+As are constraints on merger models and clues as to what these galaxies will ultimately become. In Figure \\ref{fig:color}, we show the F702W$-$F439W color profiles  of EA2-5, which are obtained by using the results from the ELLIPSE task and include the A-type K-correction. Because of the shallow exposure in F439W band, the color profiles are limited to $r \\lesssim 2 - 3$ kpc, which is only 25\\% of the radial coverage available in the red. To derive the color profiles within the most central region ($<0.5''$), we use the deconvolved images. However, because deconvolution can produce large artificial fluctuations in low signal-to-noise data, we cannot use the deconvolved F439W profiles in the outer regions of the galaxy. We compromise by using the deconvolved images for $r < 0.5''$ and the non-deconvolved images for $r > 0.5''$.  The overall colors of E+As are relatively blue globally due to the recent star formation (bottom panel of Figure \\ref{fig:color}). The radial extent of the blue colors confirms previous observations \\citep{Franx, Caldwell, norton} that the recent star formation region is not confined within the innermost regions. However, the color gradients, especially within 1 kpc of the centers, are as diverse as the overall morphologies. While, EA3 and EA5 have blue nuclei and become redder going outward, EA2 becomes bluer with radius, and EA4 shows a relatively flat profile.   The colors are the result of the complicated interplay between age, metallicity and dust. The lack of HI in these systems \\citep{chang,miller} and of any patchiness in the images of EA2-5 argue against high levels of dust (but it is still possible that high density pockets of dust are present, particularly toward the nucleus of some of these systems). With the exception of EA1, none of the E+As show the irregular, filamentary structures expected from strong dust lanes. We thus conclude that the variety of color gradients within the inner few hundred parsecs reflects variations of the spatial distribution of the young population, which in some systems appears to be preferentially located near the center of the galaxy and in others appears to avoid the center. Perhaps this reflects differences in the types of encounter involved and its ability to drive true nuclear starbursts --- e.g., differences between prograde and retrograde encounters \\citep{Barnes_Hernquist}, major versus minor mergers \\citep{Hernquist_Mihos95}, or differences in the structural properties of the progenitor galaxies \\citep{Mihos_Hernquist96}.   \\subsection{Relationship to Fundamental Plane} \\label{sec:FP} To investigate whether E+As can evolve into E/S0 galaxies, we compare the stellar kinematics and structural parameters of E+As with ``normal'' early type galaxies. \\citet{norton} found that the old component of E+A galaxies is offset (brighter by $\\sim0.6$ mag) from the the local Faber-Jackson relation. Using the structural parameters that can only be measured using \\emph{HST} imaging, we extend this comparison to various projections of the Fundamental Plane (hereafter FP) in Figure \\ref{fig:FP}. To compare our results with the FP of \\citet{Jor96}, we correct these observables to our adopted cosmology ($H_0 = 70 \\mathrm{ \\ km \\ s^{-1}\\ Mpc^{-1}}$, $(\\Omega_M,\\Omega_{\\Lambda}) =(0.3, 0.7)$). Changes in the cosmological parameters only affect zero points in the FP equation. We use velocity dispersions from \\cite{norton} for the K-star component and the structural parameters from a single $r^{1/4}$ model \\footnote{There is an issue as to which structural parameters (for the entire galaxy or the bulge only) one should use to construct the FP for S0 or disk galaxies. However, we opt to use the ($r_e$, $\\mu_e$) from a single $r^{1/4}$ profile, because we are comparing with \\citet{Jor96}, who also used single de Vaucouleurs profiles for S0's.}. We transform the F702W magnitude to a Gunn $r$ magnitude using the average (F702W $-$ Gunn $r$) color for galaxies of various Hubble type \\citep{fuku}. The average ($r-\\mathrm{F702W}$) colors range from 0.56 for elliptical to 0.51 for Scd  galaxies. Even for the extreme case of irregular (Im) galaxies, ($r-\\mathrm{F702W}$) is different by only $\\sim$ 0.1 magnitude from the average value of 0.54.  We show various projections of  the FP in Figure \\ref{fig:FP}. Figure \\ref{fig:FP}(a) shows the face-on view of the FP given by $x = (2.21~log~r_e - 0.82~log\\lan I \\ran_e + 1.24~log~\\sigma)/2.66$, $y = (1.24~log \\lan I \\ran_e + 0.82~log~\\sigma)/1.49$ (see \\citet{Jor96}). The dashed line indicates the bound set by the limiting magnitude, but the upper dotted boundary is not caused by a selection effect. Figure \\ref{fig:FP}(b) shows the edge-on view of the FP along the long axis of the distribution, given by $y=1.24~log~\\sigma - 0.82~log\\lan I \\ran_e$. Figure \\ref{fig:FP}(c) shows the Faber-Jackson relation. Figure \\ref{fig:FP}(d) shows the correlations between $r_e$ and $\\lan \\mu_e \\ran$.  The four E+As stand apart from the E/S0s in the edge-on view of the FP, but otherwise populate the same general region of the 3-D volume. The most striking deviation of the E+As among the scaling laws lies in the $\\mu_e-r_e$ correlation. EA2-5 are more than a half magnitude brighter than the median E/S0 galaxies with the same effective radii. Especially, EA3 and EA4 have a large excess surface brightness over the dotted boundary of E/S0 galaxies in the FP. We measure the excess brightness to be 0.86 and 0.54 mag relative to the dotted upper boundary of the $\\mu_e-r_e$ projection for EA3 and 4, respectively. As seen Figure \\ref{fig:FP}(c), although EA2-5 have intermediate luminosities ($-22<M_r<-20$) in comparison to E and S0 galaxies, they have a large mean surface brightness within the effective radius.  Will E+As fade onto the locus of E/S0's after a few Gyr? We estimate the amount of fading within an effective radius using STARBURST99 models \\citep{starburst99}. Assuming solar metallicity, a Salpeter IMF, an underlying 5 Gyr-old population (not evolving), and a single instantaneous starburst with a mass fraction of 50\\%, the E+As will fade $\\sim 0.5$ magnitudes during the first Gyr after the burst. This degree of fading places EA2, EA4, and EA5 on the E/S0 locus, but not EA3. This model is oversimplified, however. The evolutionary tracks in the $\\mu_e-r_e$ projection will proceed from upper-left to lower-right, because of the decline of the mean surface brightness within $r_e$ both due to fading and the inclusion of the larger fraction of the galaxy. But, calculating these effects requires a detailed knowledge of the spatial distribution of the young and old populations.    \\subsection{Comparison with to ``Core'' and ``Power-law'' Ellipticals} \\label{sec:nuker} Does the consistency in global surface brightness profiles between E+As and normal early-type galaxies (as discussed in \\S \\ref{sec:FP}) extend to the innermost radii? The recently discovered dichotomy in the central surface brightness profiles (``core'' vs. ``power-law''; Lauer et al. 1995; Faber et al. 1997) of normal early-type provides a framework for this comparison. In Figure \\ref{fig:nuker}, we plot the F702W surface brightness profiles for our E+As and the F555W profiles of \\citet{f97} sample (assuming $V - R = 0.5$ and adjusting for differences in adopted cosmology). The apparent lack of any central ``break'' and associated transition to a shallow cusp in the E+A profiles leads us to classify the E+As as ``power-law'' galaxies. Not only are the profiles shapes consistent, but the total luminosities of these E+As lie within the range spanned by normal power-law early-types, and any fading will place them even more securely in that range.  Although the profile shapes and total luminosities of E+As are consistent with those of power-law early-types, Figure \\ref{fig:nuker} shows that the E+A surface brightness are generally higher. EA2 is consistent with the highest surface brightness normal power-law galaxies, and the other three E+As have even higher surface brightness. If fading of the young stellar population does not significantly alter the profile shape, then these galaxies will retain their power-law profiles. While we cannot be certain that the profiles will remain unchanged, the current inner profiles and total luminosities are consistent with E+A evolution into normal power-law early-type galaxies.    \\subsection{Discovery of Young Star Clusters} \\label{sec:compact} A specific type of asymmetric small-scale feature is the compact sources most visible in the residual images. We identify a set of these sources using SExtractor \\citep{bertin} to analyze both the combined and residual images. Because we expect stellar clusters to be unresolved at these distances (one pixel corresponds to $\\sim 65-95$ pc at $z = 0.07 - 0.12$), confusion between real sources and hot pixels or background fluctuations is a serious problem. After testing various detection criteria and comparing visually, we decided to consider an object real only if it has at least three adjacent pixels that are each at a flux level 3 times above the local background rms. To estimate the local background level, we use a background mesh size for SExtractor of $\\sim 4 - 8$ pixels, depending on the image. The residual images are used to detect compact sources close to the galaxy center, where the rapidly varying galaxy light complicates detections in the original images.  Using these criteria, we detect 35, 9, 29, 10 and 1 point-like sources within the PC images of EA1-5, respectively. These objects may be foreground stars, background compact sources, giant HII regions,  or star clusters. Unfortunately, due to the insensitivity of the F439W band, we are able to obtain colors for only a few of the objects. The lack of  emission lines in the spectra of  these E+A galaxies \\citep{zab96,norton} suggests that these  galaxies are not littered with giant HII regions. The spatial distribution of the compact objects suggests that the majority of these sources are associated with the E+As. For example, in EA1 almost all of the point sources are located around EA1 and its companion, and three of them lie in the bridge-like structure connecting the two galaxies. In EA2, some objects lie in the elongated tidal tails. In EA3, five to six unresolved objects surround the nucleus of EA3 and form a concentric circle. We apply a Kolmogorov-Smirnoff (K-S) one-sample test to find the probability that the radial distributions of the compact sources are drawn from a random distribution. We show the cumulative distribution of the angular distances from the center of each galaxy in Figure \\ref{fig:kstest}. There is almost zero probability that the sources follow the random distribution, indicating that the detected compact sources are truly associated with the E+A galaxies. Because of the post-starburst nature of E+A galaxies and the discovery of numerous young clusters in ongoing mergers (e.g., Holtzman et al. 1992; Whitmore et al. 1993), it is likely that the majority of these objects are relatively young star clusters.  Like the morphological diversity of the E+As themselves, the number of cluster candidates varies widely among EA1-5. In particular, the difference between EA3 and EA5 is remarkable because these otherwise appear to have comparable masses (the velocity dispersions of their old populations are the same and the luminosity ratio is $\\sim 0.5$). The principal difference between these systems is that EA5 is more distant. To determine whether the distance difference is responsible for the difference in detected sources, we model the detectability of the point sources as a function of distance. For a given size (half-light radius) and brightness, we generate two dimensional surface brightness distributions of the model clusters for a King profile \\citep{king} with concentration $c = 0.5 - 2.5$. We convolve the model profiles with the PC PSF generated from the TinyTim software \\citep{Krist99}. All models and PSFs are constructed at a resolution of 1/10th of a PC pixel, re-binned to the original image size, and smeared using the pixel smearing kernel. After adding the Poisson noise and sky background, we attempt to detect the cluster using the same detection criteria that we apply to the original images. The position on the PC and the spectral variation of the PSF do not affect our detection threshold.  The number of detected pixels above 3$\\sigma$ is nearly insensitive to the size of the sources and the shape of the profile in the subpixel regime. In other words, the images of point sources give us little or no information about their spatial structures, but they can be detected if they are sufficiently bright. We adopt the detection limit as the magnitude at which the number of detected pixels above 3$\\sigma$ drops below three pixels. The number of compact sources above EA5's detection limit ($M_R<-12.5\\pm0.25$) and well outside the E+A are 6 to 16 for EA3 and is 1 for EA5. The uncertainty in the detection limit arises from uncertainties in the aperture correction and photometry. Assuming EA5 has the same number of the compact sources as EA3, we would expect to detect 3 to 8 sources in EA5 after rough scaling of the projected area occupied by the detections around each galaxy. The statistics of our sample thus preclude us from distinguishing between the cluster populations of EA3 and EA5. If the difference is significant, it might reflect a difference in the starburst (strength or progenitors) or an age difference. The latter possibility is consistent with the observation that EA5 has the reddest $D_{4000}$ and is the most symmetric among our sample.  Matching the observed colors of the cluster candidates with the predictions of stellar population synthesis models can help break the degeneracy between the age and starburst strength of our post-merger E+As. We show the evolution of the color and brightness of a simple stellar population based on Starburst99 models \\citep{starburst99} in Figure \\ref{fig:evolution}. Because of the unknown SEDs of our sources, the K correction of the observed colors is problematic. To avoid applying an uncertain K correction, we calculate the redshifted (F439W-F702W) colors from the model spectra provided from Starburst99. Superposed on the various models, we show the range of the colors observed and the corresponding range of ages for each of the five point-like sources in EA1 with $\\sigma(\\mathrm{F439W-F702W}) < 0.25$ and sufficient S/N. They are identified by eye in the F439W band and are not located in the crowded region. Because the absence of the O, B stellar signatures \\citep{zab96,norton} suggest that the ages of the clusters are older than $10^7$ yr, we rule out the solutions in the double trough of the model. Assuming a Salpeter IMF that extends from 1 to 100 $M_\\odot$, solar metallicity ($Z=0.020$), and an instantaneous starburst, we estimate the ages of the candidate clusters to be between $\\sim$ 30 Myr and 500 Myr.  We applied the following statistical test to constrain the time since the starburst under the assumption that all the candidate clusters formed at the same time and have a spread of estimated age due to photometric errors. For each model age, we draw a thousand samples of five clusters each with colors scattered according to a Gaussian of dispersion 0.2 mag. We calculate the percentage of simulations that have a median color as far from the true color as the observed median color. If that number is 5\\% or less, the model is rejected. Because of the slow evolution of colors in the relevant age range, we can constrain the age of clusters (time elapsed since the starburst) to only $\\sim$ 35 $-$ 450 Myr at the 95\\% confidence level (see the horizontal line in Figure \\ref{fig:evolution}). While our constraint on the age is drawn from only a few cluster candidates, this range is as good as the typical age ranges that \\citet{Leonardi_Rose} derived using CaII and $H\\delta/\\lambda4045$ indices from the integrated spectra of E+A galaxies. Because any reddening by the host galaxy will make the sources look older, our estimated ages are upper limits on the true ages. However, it is unlikely that the reddening affects our results seriously because the five sources are at large radii, where the extinction should be minor compared to the galaxy center. To improve upon the age estimate, in particular to reduce contamination, to correct the reddening, and to get more precise ages for the sources, ideally we would need deep multicolor photometry to produce a two-color diagram like Fig. 6 in \\citet{jharris}. However, we can see from the present analysis that we would benefit from just more clusters with colors. We conclude that the derived ages are consistent with what we expect on the basis of the integrated galaxy spectra, again supporting the claim that these are stellar clusters associated with the E+A phenomenon, and suggesting that this is part of a population of clusters formed during the E+A phase.  We compare the R band cluster luminosities of EA1, EA3, the Milky Way globular clusters \\citep{harris} and the clusters in starbursting, merging systems (e.g., NGC3597, Carlson et al 1999). If the compact sources in the E+As are clusters, they are much brighter ($M_R \\sim -13$) than Galactic globular clusters ($M_R > -11$) and similar to clusters in galaxies with on-going starburst. The latter agreement supports the interpretation that we have identified the bright end of a population (e.g., $M_R \\sim -14$ in NGC3597) of star clusters formed during a starburst that occurred $< 1$ Gyr ago. A fading of several magnitudes is required for these systems to resemble the massive end of the Milky Way cluster population.   \\subsection{Implications for High Redshift Interacting Galaxies} \\label{sec:highz} The difficulty we have experienced in determining whether these E+As have a disk component, have tidal components, and are asymmetric suggests that at high redshift these galaxies might be classified as ``normal'' morphologically. For example, the tidal tails that connect onto what may be a disk would naturally be interpreted as spiral arms if one is not able to trace the tails out to large radii. To better understand this effect, we rebin the F702W band image of EA3 and fade it artificially according to our adopted cosmology to mimic its appearance at higher redshift (Figure \\ref{fig:highz}). The rebined images are convolved with the PSF, and the sky background and noise proportional to the exposure time are added to the redshifted images. For each image, we assume the same exposure time (2100s) and average sky brightness. The tidal features are lost at a redshift over $\\sim 0.5$ (see also Mihos 1995; Hibbard \\& Vacca 1997), and EA3 appears to be a normal spiral or S0 galaxy. The mean surface brightness of the tidal tails of EA3 is $\\sim 24 - 25$ $\\mathrm{mag~arcsec^{-2}}$. This exercise suggests that some of the disky E+A galaxies found frequently in distant clusters might be misclassified as non-interacting disk galaxies and nevertheless have tidal features like EA2-4. The lack of apparent tidal features and of asymmetry at high redshift should not be interpreted directly as an absence of interactions."
    },
    "0402/astro-ph0402581_arXiv.txt": {
        "abstract": "Chandra \\acis-S3 observations of the nearby S0 galaxy \\src\\ resolve much of the X-ray emission into 73 point-sources, of which 37 lie within the \\dtwentyfive\\ isophote. The remaining galaxy emission comprises hot, diffuse gas and unresolved sources and is discussed in two companion papers. The point-source luminosity function (XLF) shows the characteristic break seen in other early-type galaxies at $\\sim 2\\times 10^{38}$ \\ergps. After applying corrections for detection incompleteness at low luminosities due to source confusion and contamination from diffuse galactic emission, the break vanishes and the data are well-described as a single power law. This result casts further doubt on there being a ``universal'' XLF break in early-type galaxies marking the division between neutron-star and black-hole systems. The logarithmic slope of the differential XLF (dN/dL), $\\beta=2.7\\pm0.5$, is marginally ($\\sim2.5\\sigma$) steeper than has been found for analogous completeness-corrected fits of  other early-type galaxies but closely matches the behaviour seen at high luminosities in these systems. Two of the sources within \\dtwentyfive\\ are Ultra-luminous X-ray sources (ULX), although neither have  \\lx$>2\\times 10^{39}$ \\ergps. The absence of very luminous ULX in  early-type galaxies suggests a break in the XLF slope at $\\sim$1--2$\\times 10^{39}$ \\ergps, although the data were not of sufficient quality to constrain such a feature in \\src. The sources have a spatial  distribution consistent with the optical light and  display a range of characteristics that are consistent with an LMXB population. The general spectral characteristics of the individual sources, as well as the  composite source spectra, are in good agreement with observations of other early-type galaxies, although a small number of highly-absorbed sources are seen. Two sources have very soft spectra, two show strong variability, indicating compact binary nature and one source shows evidence of an extended radial profile. We do not detect a central source in \\src, but we find a faint (\\lx$=2\\pm1 \\times 10^{38}$ \\ergps) point-source coincident with the centre of the companion dwarf galaxy \\srctwo. ",
        "introduction": "Prior to the launch of \\chandra\\ only a small number of the very brightest point-sources in early type galaxies could be resolved from the diffuse galactic emission \\citep[][]{fabbiano89,colbert99,roberts00}. The study of their point-source populations was therefore restricted to the average, composite properties inferred by decomposing the emission into a number of spectral or spatial model components \\citep[\\eg][]{matsushita94,brown01}. The advent of \\chandra\\ has revolutionized this field, however, allowing a large fraction of the sources to be resolved and studied directly \\citep[for a recent review, see][]{fabbiano03}. Extragalactic point-sources most probably represent an heterogeneous mixture of different source types, although the old stellar populations characteristic of early-type galaxies suggest a predominantly low-mass X-ray binary (LMXB) nature. As endpoints of stellar evolution, the properties of X-ray binaries are a crucial diagnostic for the evolution of the stars as a whole within the galaxy. Although the data are generally of insufficient quality to investigate fully for each source the rich, diagnostic phenomenology of Galactic LMXB \\citep{white95}, in some cases it has been possible to identify variability \\citep[\\eg][]{sarazin01,kraft01} and the spectral signatures of black-hole binaries \\citep[\\eg][]{makishima00,humphrey03a}.  In order to study the properties of the X-ray point-source population as a whole, it is common practice to consider the X-ray luminosity function (XLF) \\citep[\\eg][]{sarazin01,blanton01b,zezas02d}, the shape of which is a strong function of the age of the stellar population \\citep{kilgard02,belczynski03}. Between early-type galaxies, there is remarkable similarity in the shape of the XLF \\citep{kim03a}, which is typically found to be a steep power law with a break occurring around 2--4 $\\times 10^{38}$ \\ergps\\ and a high-luminosity slope, $\\beta \\simeq$ 2--3 \\citep{sarazin01,blanton01b,kraft01,colbert03a}. These slopes are also broadly consistent with old stellar populations in M31 \\citep{kong03}. It has been suggested that the presence of breaks at luminosities at around the Eddington limit of a 1.4${\\rm M_\\odot}$ neutron star (${\\rm L_{EDD}=}$2--4$\\times 10^{38}$ \\ergps, depending on the composition of the accreting matter and the neutron-star equation of state: \\citealt{paczynski83}) may arise from  a division between neutron-star and black-hole systems \\citep[\\eg][]{sarazin01}. However in the elliptical galaxy NGC\\th 720, a substantially higher luminosity break, at ${\\rm \\sim 1\\times 10^{39}}$ \\ergps\\ was found instead by \\citet{jeltema03}. \\citet{kim03b} pointed out that in the elliptical galaxy NGC\\th 1316, a break is arises only if no correction is made for  point-source detection incompleteness. Applying this correction to a sample of early-type galaxies, \\citet{kim03a} found a similar result, although they still found marginal evidence of a break at $\\sim 5 \\times 10^{38}$ \\ergps. This luminosity is sufficiently high, however, to cast further doubt upon the idea that it marks the division between neutron-star and black-hole binaries.  At the extreme end of the XLF are the  so-called ``Ultra Luminous X-ray sources'' (ULX), which are non-nuclear objects with luminosities exceeding $10^{39}$ \\ergps, sometimes reaching as high as $10^{41}$ \\ergps\\ \\citep[\\eg][]{fabbiano89,makishima00,zezas02b,davis03}. They have been found in galaxies of all morphological types \\citep{colbert02}, although there is a strong association between ULX and star-formation so that their specific frequency is far lower in early-type galaxies \\citep{kilgard02,humphrey03a,irwin03b}. However, in the otherwise normal elliptical NGC\\th 720, \\citet{jeltema03} found a remarkable population of 9 ULXes. The nature of ULX is somewhat enigmatic. Although the less extreme objects (\\lx$\\sim$1--2$\\times 10^{39}$ \\ergps) are consistent with isotropic emission from Eddington-limited black-hole binaries with compact-object masses $\\sim$10--20${\\rm M_\\odot}$, there has been much debate over whether the most luminous objects represent ``Intermediate mass black holes'' with mases \\gtsim ${\\rm 100 M_\\odot}$, super-Eddington emission from sources in the thermal-timescale mass-transfer phase, or beamed emission from lower-luminosity sources \\citep{colbert99,king01,king02}.  Another feature of the LMXB population of particular interest is its possible association with globular clusters (GCs), since it provides valuable insight into X-ray binary formation. The fraction of the extragalactic LMXB identified with GCs varies considerably between galaxies, from $\\sim$20--70\\% \\citep{angelini01,kundu02a,sarazin01}. Nonetheless, there is evidence that the frequency of at least part of the LMXB population correlates with the numbers of GCs so that $\\sim$4\\% of all GCs contain an LMXB, consistent with the Milky Way \\citep{kundu03,sarazin03}. These authors also found evidence that LMXB favour brighter, redder GCs.  In this paper, we analyse with \\chandra\\ the hitherto unresolved point-source population of the lenticular galaxy \\src. Using \\rosat\\ \\pspc\\ data of this galaxy, \\citet[][hereafter BC]{buote96a} found elongation of the X-ray isophotes and demonstrated that the galaxy mass profile was marginally inconsistent with the optical light, indicating the presence of a substantial dark matter halo. Unfortunately, most of the point-source population could not be resolved, so that it was necessary to estimate limits for the unresolved point-source contribution to the X-ray emission, and tight constraints on the shape of the dark matter halo could not be obtained. Using \\asca\\ data \\citet{buote97a} found a hard spectral component in the integrated stellar emission, indicating a significant contribution from unresolved LMXB, although this did not allow these sources to be studied in detail.\\chandra\\ allows us for the first time to separate a large fraction of the LMXB from the diffuse emission of \\src, and to study them directly. The properties of the  diffuse gas and the mass distribution are discussed separately in two companion papers \\citep[][hereafter Papers I and II, respectively]{humphrey04b,buote04a}.  In order to determine the properties of the stellar population accurately, it is important to adopt a reliable distance estimate (\\cf\\ NGC\\th 4038/9, for which a recently revised distance estimate would reduce the reported ULX population threefold; \\citealt{saviane03}). For NGC\\th 1332, distances determined both by surface brightness fluctuations (SBF) and the globular cluster luminosity function (GCLF) are in excellent agreement, being $23\\pm3$\\th Mpc and $22\\pm6$\\th Mpc, respectively \\citep{tonry01,kundu01}.  All errors quoted here, and subsequently, are 90\\% confidence limits unless otherwise stated. ",
        "conclusions": "We have found a total of 73 point-sources within the \\chandra\\ S3 field of \\src, including 37 which lie within the \\dtwentyfive\\ isophote and none of which is coincident with the galaxy centroid. We expect a total of $1.9^{+1.5}_{-1.2}$ unrelated background or foreground sources to be within \\dtwentyfive, so that the vast majority of these sources are associated with the galaxy. The spatial distribution of the sources was consistent with the optical isophotes, as may be expected from a population of sources of stellar origin which has not been recently disturbed by merger activity \\citep{zezas03a}. In total 30\\% of the X-ray sources were associated with globular clusters, and 4.4\\% of GCs contained an X-ray source, both of which results are consistent with other early-type galaxies. The phenomenology of the individual sources was consistent with expectations for an LMXB population. Two sources had soft spectra, consistent with high-state black-hole binaries and two sources were variable during the observation. Intriguingly one source was extended, although its radial profile was consistent with two unresolved point-sources.  Without taking into account the impact of source detection incompleteness at the lowest luminosities (where source confusion and increased background due to diffuse gas reduce the probability of detecting a source), we found that the XLF of the point-sources within \\dtwentyfive\\ can be fitted by a power law with a break at \\lx$\\sim 2\\times10^{38}$ \\ergps. Similar results have been reported in other early-type galaxies \\citep[\\eg][]{sarazin01,blanton01b}. However, when we made corrections for the expected level of incompleteness, we found that the data were extremely well-fitted by a single power law, with a differential slope, $\\beta=$ $2.7\\pm0.5$. Assuming that the XLF remains unbroken down to $10^{37}$ \\ergps\\ (as seen for sources in M\\th 31: \\citealt{kong02}), we estimate a total luminosity of ${\\rm 180\\pm70 \\times 10^{38}}$ \\ergps\\ due to {\\em unresolved} point-sources in the \\dtwentyfive\\ region. Comparing this to the B-band luminosity of the galaxy, obtained from the face-on, extinction-corrected B-band luminosity listed in the \\leda\\footnote{\\href{http://leda.univ-lyon1.fr/}{http://leda.univ-lyon1.fr/}} catalogue, we obtained  ${\\rm L_{LMXB}/L_B=0.8\\pm0.3\\times 10^{30}}$~${\\rm erg\\ s^{-1} L_{B\\odot}^{-1} }$. In addition, we compared with the K-band luminosity (adopting the ${\\rm K_{20}}$ magnitude obtained by 2MASS), and obtained ${\\rm L_{LMXB}/L_K=0.13\\pm0.05\\times 10^{30}}$~${\\rm erg\\ s^{-1} L_{K\\odot}^{-1} }$. Both of these values were in excellent agreement with observations of other early-type galaxies \\citep{kim03a}.  The slope of our XLF fit closely resembles that found at high luminosities in a variety of other early-type galaxies \\citep{sarazin01,blanton01b,colbert03a}. A slope of $\\beta \\simeq 2.7$ is also strikingly similar to observations of a primarily old stellar population in M31, but which does not extend to such high luminosities \\citep{kong03}. The slope is, however, marginally ($\\sim2.5\\sigma$) steeper than the incompleteness-corrected {\\em single} power law fits found in a recent survey of early-type galaxies \\citep{kim03a}. Intriguingly, these authors reported some evidence of a break at $\\sim 5\\times 10^{38}$ \\ergps, above which $\\beta \\simeq 2.7$ is observed. A similar result has also been found for M84 \\citep{finoguenov02}. A break at this luminosity should be observable within our data, and its apparent absence is intriguing. Whilst it is certainly true that {\\em statistically} acceptable single power law fits were found for most galaxies in the sample of \\citet{kim03a}, the resulting slope ($\\beta\\simeq2.0$) may simply be an ``average'' of the slopes above and below the break. The steep slope seen in \\src\\ might therefore suggest that such a break is absent in this galaxy. It must be stressed, however, that there are some differences in the analysis procedures adopted by these authors and those used in the present work, which may simply weight our fit somewhat more heavily towards the high luminosity sources (especially our inclusion of a high-luminosity, unoccupied bin). Due to the small numbers of sources, however, the error-bars on the best-fit slope in \\src\\ are large, so that the disagreement is only at $\\sim 2.5\\sigma$. It remains to be seen whether comparable results are observed in any other galaxies containing significantly more X-ray point sources. Nonetheless, if the absence of the break in our XLF is confirmed, it is an interesting result, which casts further doubt on there being a ``universal'' break in early-type galaxies corresponding to the division between neutron-star and black-hole binary systems.  The origin of the XLF break is still somewhat unclear. Although it was originally suggested that a break might mark the division between neutron-star and black-hole binary systems \\citep{sarazin01}, the disparity between the reported break luminosities in several galaxies and the Eddington limit for a 1.4${\\rm M_\\odot}$ (canonical) neutron-star suggests another origin, perhaps in the star-formation history of the galaxy \\citep[\\eg][]{jeltema03,wu01a}. In that case, however, we might expect to see a strong correlation between the age of the system (since the last major merger) and the XLF shape; comparing the slopes found by \\citet{kim03a}, including that for NGC\\th 720, and galaxy ages given by \\citet{terlevich02} we find no evidence of a correlation. However, this includes only 8 sources and will need to be confirmed with a larger sample.    We detected  two ULX within the \\dtwentyfive\\ region of \\src, although neither of these sources have \\lx$>2\\times 10^{39}$ \\ergps. Such numbers are consistent with other  galaxies of similar morphology \\citep{humphrey03a}. The unbroken XLF of late-type galaxies suggests that the ULX associated with star-formation can be attributed to the high-luminosity tail of the HMXB XLF, which is flatter than the LMXB XLF \\citep{grimm03}. It may be that, in a similar way, the ULX in early-type galaxies simply represent the high-\\lx\\ tail of the LMXB XLF. \\citet{irwin03b} have argued that sources with \\lx$>2\\times 10^{39}$ \\ergps\\ are extremely rare in  early-type galaxies, suggesting that the most luminous ULX may therefore represent a population almost entirely alien to such galaxies. However luminous ULX have been reported in NGC\\th 1399 and NGC\\th 720  which are associated with GCs \\citep[hence unlikely to be background objects;][]{angelini01,jeltema03}, so they are clearly not entirely absent and therefore it is of interest to determine if their paucity simply reflects the steepness of the XLF. In \\src, since our fitting took account of the absence of sources at high luminosities, the absence of very luminous ULX is consistent with the high-luminosity tail of the XLF. However, we also compared the ratio of the number of sources measured by \\citet{irwin03b} in the 1--2 ${\\times 10^{39}}$ \\ergps\\ and the $>2\\times 10^{39}$ band, to that expected assuming our best-fit XLF slope ($\\beta=2.7$). We found a likelihood of $\\sim 1$\\% of reproducing the data from the model, although allowing $\\beta$ to steepen within its error-bars increased the likelihood to $\\sim 5$\\%. Nonetheless this suggests the possibility of a change in the XLF at $\\sim$1--2$\\times 10^{39}$ \\ergps. Intriguingly, an XLF break was reported by \\citet{jeltema03} around this luminosity."
    },
    "0402/astro-ph0402548_arXiv.txt": {
        "abstract": "{Using the {\\it Hubble space Telescope Fine Guidance Sensor}, we have measured the trigonometric parallax of the bright cataclysmic variable 1223\\,Sgr. The absolute parallax is $\\pi_{\\rm abs}= 1.96\\pm 0.18$\\,mas, making V1223\\,Sgr the most distant CV with a well-determined trigonometric parallax. This distance, a Lutz-Kelker correction, and the previously measured extinction yield an absolute visual high-state magnitude $M_{\\rm V} =4.0\\pm0.2$.  We outline a model, which is consistent with the observed spin-down of the white dwarf and provides for much of the UV/optical emission by reverberation of X-rays. From previous X-ray {and UV/optical} data, we derive {an accretion luminosity \\lacc\\ = $\\ten{(2.6\\pm0.8)}{34}$\\,\\erg}, a white dwarf mass $M_1 = 0.93\\pm0.12$\\,\\msun, and an accretion rate $\\dot M = \\ten{(1.4\\pm0.3)}{17}$\\,\\gs. ",
        "introduction": "\\label{sec-intro}  V1223\\,Sgr (4U1849-31) is one of the brightest confirmed intermediate polars, both optically and in X-rays. The only previous distance estimate of 600\\,pc was based on the \\mbox{observed $E_{\\rm B-V}=0.15$ and a mean reddening} of 0.25\\,mag\\,kpc$^{-1}$ (Bonnet-Bidaud et al. 1982). Since V1223 Sgr is not included in ground-based parallax programs, we have obtained a high-precision parallax using the {\\it Hubble Space Telescope Fine Guidance Sensor} ({\\it FGS}). This observation increases the number of CV parallaxes measured with the {\\it HST FGS} to eight (Harrison et al. 1999, 2000, 2003, McArthur et al. 1999, 2001, Beuermann et al. 2003).  V1223\\,Sgr shows modulations at the orbital period \\porb=3.366\\,hrs, the spin period of the white dwarf \\pspin=745.5\\,s (Osborne et al. 1985), and the beat period \\pbeat=794.38\\,s (Steiner et al. 1981). The UV/optical spectral energy distribution was interpreted by Bonnet-Bidaud et al. (1982) and Mouchet (1983) in terms of an accretion disk which reaches close to the white dwarf and should cause it to be spun up. Unexpectedly, however, the white dwarf was found to spin down (Jablonski \\& Steiner 1987, van Amerongen et al. 1987) suggesting that the disk stays much further away from the white dwarf. Our accurate distance allows us to determine the luminosity of V1223\\,Sgr, estimate its accretion rate, and to show that a standard intermediate-polar model is consistent with the available observations. ",
        "conclusions": "We have presented an accurate parallax of V1223 Sgr which allows us to derive the luminosities in the different wavelength bands. Based on this result, we have tested the hypothesis that much of the UV/optical emission is produced by the reverberation of X-rays and not by the release of {gravitational} energy in a luminous disk. {Our analysis is based on the assumption that we see only the X-rays from one pole. This assumption and the high X-ray temperature reported by Beardmore et al. (2000) lead to a total X-ray luminosity $L_{\\rm x} = \\ten{(2.2\\pm0.7)}{34}$, which is much higher than previously thought and can power a large fraction of the observed UV/optical radiation by the reprocessing of X-rays. Interestingly, this is possible without the presence of a substantial source of XUV radiation.  Likely reprocessing sites are the pole cap of the white dwarf, the magnetically guided accretion flow, the disk, and the irradiated face of the secondary star. While the latter two sites may be responsible for the observed flux modulated at the sideband frequency (Steiner et al. 1981), reprocessing in the former locations may produce a so far undiscovered component, which is photometrically modulated at the spin period of the white dwarf. The accretion rate derived from the X-ray luminosity and X-ray temperature is $\\dot M = \\ten{(1.4\\pm0.3)}{17}$\\,\\gs, which includes a small correction for the contribution by cyclotron emission. This result is independent of the white dwarf mass.  The observed UV/optical/IR spectral energy distribution can equally well be fitted by a truncated disk plus reprocessed component (model~B) and by a luminous disk (model~A). The observed spin-down of the white dwarf (Jablonski \\& Steiner 1987, van Amerongen et al. 1987), however, requires that the the inner edge of the disk is not too far inside the corotation radius, a condition which is met only by the truncated-disk model~B.  Strong internal absorption and the reprocessing of a major fraction of the emitted X-rays is a general feature observed also in other intermediate polars, some of which} are even more strongly internally absorbed than V1223\\,Sgr (e.g. Norton \\& Watson 1989).  Simultaneous X-ray/UV studies could shed light on the physical processes acting in these systems. It is surprising that no such study is yet available."
    },
    "0402/astro-ph0402124_arXiv.txt": {
        "abstract": "We investigated the effect on the $U,~B,~V$ magnitudes of the removal of the emission lines from the spectra of some symbiotic stars and novae during their nebular phases. We approached this problem by a precise reconstruction of the composite UV/optical continuum and the line spectrum. Corrections $\\Delta U$, $\\Delta B$, $\\Delta V$ are determined from the ratio of fluxes with and without emission lines. Here we demonstrate this effect on the case of the symbiotic nova V1016~Cyg during its nebular phase. We find that about 68\\%, 78\\% and 66\\% of the observed flux in the $U$, $B$ and $V$ filters is radiated in the emission lines. The effect should be taken into account before using the observed color indices of emission-line objects for diagnosis of their radiation in the continumm. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0402/astro-ph0402312_arXiv.txt": {
        "abstract": "We discuss the Bose-Einstein condensation of relativistic vector charged particles in a strong external magnetic field in very dense matter, as may be paired spin-up electrons. We show that for electrons such systems may maintain self-consistently magnetic field strengths in the interval $10^{10}-10^{13}$ Gauss. This could be the origin of large magnetic fields in some white dwarfs, but may also impose bounds due to the arising of strong anisotropy in the pressures, which may produce a transverse collapse of the star. ",
        "introduction": "In the paper \\cite{Chaichian} was shown that a relativistic gas under the influence of a magnetic field $B$ of order of the quantum electrodynamics  limit of $m_e^2/e \\sim 10^{13}$ Gauss and for densities $N \\leq 10^{30}$ cm$^{-3}$ becomes unstable and collapses, since the pressure perpendicular to the field vanishes. Physically, the system, infinitely degenerate with regard to the orbit's center quantum number,  becomes otherwise one dimensional, with all the electrons falling on the Landau ground state $n=0$. The magnetic ``Bohr radius\" being of order $\\sqrt{\\hbar c/eB}=10^{-2}$\\AA. Its spectrum looks equivalent to that of a free unidimensional particle moving along the external field. under these conditions the system is unable to exert any transverse pressure.  Looking at the problem from another side, in a recent paper \\cite{Chakrabarty} has been shown that an electron gas confined to the Landau ground state cannot be in $\\beta$-decay equilibrium in a neutron star due to an incompatibility among the spin orientation of the particles involved. The suggestion is given in \\cite{Chakrabarty1} that a bosonization of the electron gas may take place, which would mean a solution to the above-mentioned problem. We will discuss here the behavior of the bosonic gas whose constituent particles are relativistic bosons, which we suggest may represent the electron pairs, under such extreme conditions of confinement to the $n=0$ Landau ground state. As is known in normal superconductors, scalar pairing condensates occur in absence of a magnetic field (Cooper pairs). The pairing leading to bosonization is mediated by some interaction as is the Coulomb exchange interaction among electrons, for instance, in the ferromagnetic pairing in atoms. In addition to bosonization, another effect could be associated with a small change in the new particle mass, that is, the final particle would have a mass smaller than twice that of the initial unpaired particles since the whole relativistic energy is reduced due to inertial effects. In view of this, we shall assume that this binding energy is not significant.   When a magnetic field is applied, the superconductivity (condensate) is distroyed at some critical magnetic field, (the Schafroth critical field). But if the magnetic field increases largely enough  to have a significant fraction of its density in the Landau ground state, it has been suggested that the condensate reappears as a consequence of some interaction compatible with a spin-one vector pairing \\cite{Boebinger}. This would lead to a superconductive-ferromagnetic behavior.  Such particles would  carry twice the charge of the electron and an effective mass which in principle we take of order twice of that of the electron, although some corrections must be introduced however, due to effects coming from the large density and the magnetic field. Then the system may be treated by following the same formalism used in previous references \\cite{Chaichian,Hugo}. As a consequence of condensation the system would behave as ferromagnetic and under the action of an external field $H$, a magnetization ${\\cal M}$ arises, leading us to define a microscopic magnetic field $B=H+4\\pi {\\cal M}$. The interesting point here is that, due to the positive character  of ${\\cal M}$ it may occur that $B \\sim 4\\pi {\\cal M}$, or $H \\ll 4\\pi {\\cal M}$ i.e., the microscopic magnetic field be produced by self-magnetization. ",
        "conclusions": "From all the previous reasoning we conclude that an  electron system, as a white dwarf, can be hardly stable at fields of order $B_c$. In principle, such fields can be maintained self-consistently, but the possibility of a collapse is highly increased: the one-dimensional world created by the magnetic field is completely unstable. The previous results, if applied to the condensation of $\\rho$ and $\\omega$ mesons (for instance, in neutron stars), would be compatible with the self-magnetization condition for densities of order $10^{39}$~cm$^{-3}$ if the resulting magnetic fields are in the interval $10^{17}-10^{19}$~G. Instabilities would lead to a collapse. This point deserves further research, as it provides additional arguments to those given in \\cite{aurora} against the claimed existence of the so-called magnetars, because those objects seem to be unstable under such $\\rho$ and $B$ conditions."
    },
    "0402/astro-ph0402638_arXiv.txt": {
        "abstract": "I will review the constraints set by X-ray measurements of afterglows on several issues of GRB, with particular regard to the fireball model, the environment, the progenitor and  dark GRB. ",
        "introduction": "This conference took place few months after the switch-off of BeppoSAX (\\cite{psbsax95}), on April 30, 2002. Launched on April 30, 1996, this mission carried out  observations of all classes of X-ray sources during its  operative lifetime of 6 years. A total of 62 Msec of pointed observations with its Narrow Field Instruments (NFI) were carried out. A substantial fraction (about 50\\%) of the total observing programme was devoted to observations of compact galactic sources and AGN, i.e. the classes of sources mostly suited to the exploitation of the broad band spectral coverage of BeppoSAX NFI (0.1-200 keV).  The other strong asset  of the mission was the capability of discover and carry out deep observations of transient phenomena in the sky. This was assured by wide field X-ray and gamma-ray monitors (Wide Field Cameras, WFC, and Gamma-Ray Burst Monitor, GRBM) coupled with a high level of flexibility of  ground scientific operations in carrying out fast Target of Opportunity Observations (TOO) with NFI. In fact, a substantial part of the program was devoted to such observations: about 190 NFI observations (corresponding to a total of 7.2 Msec), out of which 2.2 Msec  on Gamma-Ray Bursts. Turning then to GRB, 56 GRB (including 8 X-ray rich GRB) were localized by wide field instruments and their position distributed within few hours. 38 GRB were observed with fast TOO observations (from 5 hrs to 1 day) with NFI. The first GRB observation took place on July 20, 1996, during the scientific verification phase (Piro et al, 1997), and the last one just the last days of operations. The most famous events were GRB970228, that led to the discovery of the first X-ray and optical afterglows (\\cite{cfh+97,vgg+97}),  GRB970508 whose precise and fast localization (\\cite{paa+98}) allowed the first determination of  distance and the discovery of the first radio afterglow and fireball observational evidence (\\cite{mdk+97,fkn+97}), and GBR980425 (\\cite{paa+00}), with its association with SN1998bw (\\cite{gvv+98}).  In recent years, most of the research activities in the field have focussed on 3 main topics.  \\begin{itemize}  \\item Progenitor and central engine.  \\item Origin of dark GRB, X-ray flashes and short GRB  \\item Cosmology with GRB  \\end{itemize}  Those areas of research are closely intertwined. The origin of dark GRB or X-ray flashes could have relevant implications both on the progenitor/central engine and on cosmological studies. The nature of the progenitor is then relevant to cosmological studies with GRB, because of their possible use as  tracers of star-formation in the Universe.  In this review I will focus on the impact of X-ray measurements on the fireball model, the environment and the origin of the progenitor,  and on dark GRB.   \\begin{figure} \\plottwo{all_grb_powerlaws.ps}{Flux_X_11hr_all_in_one_7_gaus.ps} \\caption{Light curves (best fit power laws: left panel) and distribution of F(1.6-10 keV) at 11 hrs (right panel) with BeppoSAX} \\end{figure} ",
        "conclusions": ""
    },
    "0402/astro-ph0402474_arXiv.txt": {
        "abstract": "We present a new method for measuring the CMB temperature quadrupole, using large scale CMB polarization. The method exploits the fact that CMB polarization is partially sourced by the local temperature quadrupole. We link the temperature with the polarization spectrum directly by relating the local quadrupole at the onset of reionization to both of them. The dominant contribution is at $l<30$ and since we use many $l$ values, we can reduce the error significantly below cosmic variance. In particular, for our fiducial model, the error on the temperature quadrupole is reduced to $24\\%$. This has the potential of reducing the probability of a low quadrupole by two orders of magnitude. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0402/astro-ph0402197_arXiv.txt": {
        "abstract": "We report on the \\cxo\\ ACIS-S3 imaging observation of the Galactic globular cluster M4 (NGC\\,6121). We detect 12 X-ray sources inside the core and 19 more within the cluster half-mass radius. The limiting luminosity of this observation is $\\Lx\\approx10^{29}$~\\ergsec\\ for sources associated with the cluster, the deepest X-ray observation of a globular cluster to date. We identify 6 X-ray sources with known objects and use \\rosat\\ observations to show that the brightest X-ray source is variable. Archival data from the \\emph{Hubble Space Telescope} allow us to identify optical counterparts to 16 X-ray sources. Based on the X-ray and optical properties of the identifications and the information from the literature, we classify two (possibly three) sources as cataclysmic variables, one X-ray source as a millisecond pulsar and 12 sources as chromospherically active binaries. Comparison of M4 with 47\\,Tuc and NGC\\,6397 suggests a scaling of the number of active binaries in these clusters with the cluster (core) mass. ",
        "introduction": "Faint X-ray sources ($\\Lx \\lapp 10^{35}$~\\ergsec) were first discovered in observations made with the \\textit{Einstein} and \\rosat\\ observatories (Hertz \\& Grindlay\\,1983\\nocite{hg83}; Verbunt\\,2001\\nocite{ver01}). It was only with \\chandra, however, that a large number of faint sources was identified: more than one hundred in 47\\,Tuc, and up to a few dozen each in e.g.\\ NGC\\,6397, NGC\\,6752, NGC\\,6440, and $\\omega$\\,Cen (Grindlay et al.\\,2001a,b\\nocite{ghem01,ghe+01}; Pooley et al.\\,2002a,b\\nocite{plh+02,plv+02}; Rutledge et al.\\,2002\\nocite{rbb+02}). These faint sources represent a mix of objects with different X-ray luminosities. Brightest are the neutron stars accreting at a low rate from a companion (quiescent low mass X-ray binaries or qLMXBs), followed by white dwarfs accreting from low-mass companions (cataclysmic variables or CVs). The active binaries (ABs) tend to be the faintest, while radio pulsars with short periods (millisecond or recycled pulsars, MSPs) have X-ray luminosities in a similar range as CVs. The three types of chromospherically or magnetically active binaries are detached binaries of two main sequence stars (BY~Dra systems), detached binaries of a main sequence star and a giant or a sub-giant (RS~CVn systems) and contact binaries (W~UMa systems).  Globular clusters contain many more neutron star binaries per unit mass than the galactic disk. Hence, if in a globular cluster one finds a binary with a neutron star (which is rather difficult to make via ordinary evolution of an initial binary), it is highly probable that this binary was formed via a close encounter between stars (Fabian et al.\\,1975\\nocite{fpr75}; Hills\\,1976\\nocite{hil76}). Isolated millisecond pulsars in globular clusters are thought to have been formed in such binaries, and thus are probably also a result of close stellar encounters.  On the other hand, binaries that are common in the field, are more likely to be of primordial origin in globular cluster, including the chromospherically active binaries. Cataclysmic variables are relatively common in the Galactic disk, and those in globular clusters could in principle originate from primordial binaries or stellar encounters. Pooley et al.\\,(2003\\nocite{pla+03}) showed that the number of faint sources above the threshold $\\Lx> 4\\times10^{30}$ \\ergsec\\ in a cluster scales with its collision number $\\Gamma$, which is a theoretical estimate for the number of close encounters. Since the majority of such sources are cataclysmic variables, this suggests that most cataclysmic variables in globular clusters are in fact formed via close stellar encounters.  Most clusters investigated with \\chandra\\ so far have relatively high collision numbers (see e.g.\\ Table~1 in Pooley et al.\\,2003\\nocite{pla+03}); because the limit to which X-ray sources in these clusters could be detected is relatively high, most X-ray sources known in these clusters are either neutron stars (accreting or radio pulsars), or cataclysmic variables. So far, a sizable number of X-ray sources has been identified as chromospherically active binaries only in 47\\,Tuc (Edmonds et al.\\,2003a,b\\nocite{eghg03a,eghg03b}). In this paper, we discuss a globular cluster with a relatively low collision number, M4 (NGC\\,6121). This cluster is a relatively nearby cluster with a moderate absorption ($d=1.73$~kpc, $A_V=1.32$, Richer et al.\\,1997\\nocite{rfi+97}). The core and half-mass radii of M4 are $49\\farcs8$ and $3\\farcm65$, respectively (Harris\\,1996\\nocite{har96}). We use these values throughout the paper.  A priori, we would thus expect that the majority of X-ray sources in this cluster to be chromospherically active binaries. That chromospherically active binaries exist in this cluster is evident from optical studies.  For example, Kaluzny et al.\\,(1997\\nocite{ktk97}) discovered a number of optical variables, including several contact binaries.  The presence of a recycled radio pulsar PSR~B1620$-$26 (Lyne et al.\\,1988\\nocite{lbb+88}) in M4 is remarkable, but can be explained by noting that the small collision numbers of many globular clusters still add up, so that at least some of the clusters with small $\\Gamma$ should contain a binary with a neutron star (Verbunt \\& Hut\\,1987\\nocite{vh87}). PSR~B1620$-$26 is in a 191 day orbit around a white dwarf of $\\sim\\!0.3$~\\Msun; the binary is accompanied by a third object of planetary mass in an orbit of $\\sim100$~years (Thorsett et al.\\,1999\\nocite{tacl99}; Sigurdsson et al.\\,2003\\nocite{srh+03}). ",
        "conclusions": "The X-ray luminosities of the \\chandra\\ sources in M4 are amongst the lowest ever observed in a globular cluster. More than half of the X-ray sources in this cluster have $\\Lx<10^{30}~\\ergsec$, compared to 10\\% of the sources in 47\\,Tuc (Grindlay et al.\\,2001a\\nocite{ghem01}) and 28\\% of the sources in NGC\\,6397 (Grindlay et al.\\,2001b\\nocite{ghe+01}).  It is therefore not surprising that only one of the X-ray sources detected with \\chandra\\ may have been detected before. The position of the marginal \\rosat\\ HRI source R9 (Verbunt 2001\\nocite{ver01}) is 7$\\farcs$5 from that of CX1. Even though this is further than expected from the error given for the \\rosat\\ HRI source, we think the identification is probable. The X-ray luminosity of the \\rosat\\ HRI source is only \\ee{1.3}{31}\\ \\ergsec, a factor six below the luminosity detected with \\chandra. The upper limit from the \\rosat\\ PSPC observation is  \\ee{1.5}{31}\\ \\ergsec. (These luminosities are re-computed for a distance of 1.73\\,kpc and for the spectrum of CX1, and thus differ from those given by Verbunt; mainly because he used a distance of 2.2\\,kpc.)  This large variability indicates that CX1 is a cataclysmic variable rather than a radio pulsar. Note that even if the \\rosat\\ source is not identical to CX1, it still provides an upper limit to the flux of CX1 during the \\rosat\\ observations, and thus proves that CX1 is highly variable.  Several other X-ray sources coincide with optical variables discovered by Kaluzny et al.\\,(1997\\nocite{ktk97}) and Mochejska et al.\\,(2002\\nocite{mktp02}). The sources CX13, CX15 and CX18 are coincident with W~UMa variables V49, V48 and V55, with orbital periods of  0.283, 0.298 and 0.311~days, respectively. The upper limit on the X-ray luminosity of the other W~UMa binaries found by Kaluzny et al.\\,(1997\\nocite{ktk97}), V44, V47, V50, V51, V53 and V54, is about \\ee{3}{29}\\,\\ergsec. Both detections and upper limits of these W UMa variables are in agreement with the range of X-ray luminosities of W~UMa binaries in the \\rosat\\ All Sky Survey (St\\c{e}pie\\'n, Schmitt \\& Voges\\,2001\\nocite{ssv01}).  CX8 coincides with the variable V52, which is classified as a BY~Dra system with a period of 0.777~days. The possible blend of X-ray sources CX5 and CX9 coincides with the variable V56. No period is known for this object, but its variability (a rise of 0.1 mag in 4 days) and its location on the giant branch (Mochejska et al.\\,2002\\nocite{mktp02}) are suggestive of an RS CVn variable.  If it is, it is the first RS CVn binary detected in X-rays in a globular cluster -- a marked contrast with old open clusters, where RS CVns dominate the X-ray sources (Belloni et al.\\ 1998\\nocite{bvm98}). The exposure of the eclipsing binary V54 is not as good as for the other sources, hence the upper limit on the X-ray luminosity of this source is somewhat higher. Detections and upper limits of these binaries are well within the range observed for magnetically active binaries, e.g.\\ in the \\rosat\\ All Sky Survey (Dempsey et al.\\ 1993\\nocite{dlfs93}).  \\begin{deluxetable}{lcccccc} \\footnotesize \\tablecaption{Scaling Parameters of M4, NGC\\,6397 and 47~Tuc\\label{tab:scaling}} \\tablehead{ \\colhead{Cluster} & \\colhead{log $\\rho_0$} & \\colhead{$r_\\mathrm{c}$} & \\colhead{$d$} & \\colhead{$M_V$} & \\colhead{$\\Gamma$} & \\colhead{$M_\\mathrm{c}$} \\\\ \\colhead{} & \\colhead{($L_\\sun\\ \\mathrm{pc}^{-3}$)} & \\colhead{(\\arcsec)} & \\colhead{(kpc)} & \\colhead{} & \\colhead{} & \\colhead{}} \\startdata M4        & 4.01 & 49.8 & 1.73 & $-$6.9 & 1.0 & 1.0 \\\\ NGC\\,6397 & 5.68 &  3.0 & 2.3 & $-$6.6 & 2.1 & 0.024 \\\\ 47~Tuc    & 4.81 & 24.0 & 4.5 & $-$9.4 & 24.9 & 12.4 \\enddata \\tablecomments{Values for central density ($\\rho_0$), core-radius ($r_\\mathrm{c}$), distance ($d$) and absolute visual magnitude ($M_V$) originate from Harris\\,1996 (version of February 2003). For M4, the values of $\\rho_0$ and $M_V$ are computed for the distance and reddening of Richer et al.\\,(1997). The collision number is computed from $\\Gamma \\propto \\rho_0^{1.5}\\ r_\\mathrm{c}^2$ and the core mass from $M_\\mathrm{c} \\propto \\rho_0\\ r_\\mathrm{c}^3$. Values for $\\Gamma$ and $M_\\mathrm{c}$ are normalized to the value of M4.} \\end{deluxetable}  In trying to determine the X-ray luminosity function of M4, we can use the number of detected sources as a function of luminosity. However, the number of excess counts, not allocated to individually detected sources, also contains information about the luminosity function. If we assume a luminosity function $dN\\propto{\\Lx}^{-\\gamma}d\\log \\Lx$ and a reference luminosity $L_\\mathrm{r}$, then the ratio of the contributions to the total luminosity by sources in the ranges (0.1--1)$L_\\mathrm{r}$ and (1--10)$L_\\mathrm{r}$ is given by \\begin{equation} R_L(L_\\mathrm{r}) \\equiv {\\int_{0.1L_\\mathrm{r}}^{L_\\mathrm{r}}\\Lx \\mathrm{d} N\\over\\int_{L_\\mathrm{r}}^{10L_\\mathrm{r}}\\Lx \\mathrm{d} N} = 10^{\\gamma-1} \\label{eqL} \\end{equation} and the ratio of the number of sources in the same ranges is \\begin{equation} R_N(L_\\mathrm{r}) = 10^{\\gamma} = 10 R_L(L_\\mathrm{r}). \\label{eqN} \\end{equation} In particular, we note that for $\\gamma<0.7$ and a luminosity function continuing to arbitrarily low luminosities the total luminosity of all sources with $\\Lx < L_\\mathrm{r}$ is less than the luminosity for the sources with (1--10)$L_\\mathrm{r}$.  In the following we assume that the ratio of counts is proportional to the ratio of luminosities, and we use as a reference luminosity $L_\\mathrm{r}\\equiv6\\times10^{29}$~\\ergsec.  Following the method described in Johnston \\& Verbunt\\,(1996\\nocite{jv96}) and Pooley et al.\\,(2002b\\nocite{plv+02}) we derive from the list of detected X-ray sources (Table~\\ref{tab:srcs}) that $\\gamma=0.71$ for X-ray luminosities above $\\Lx(0.5-2.5)=4.4\\times10^{29}$~\\ergsec. The K-S probability of this value is 93\\%. For K-S probabilities above 10\\% the slope of the luminosity function has $0.47<\\gamma<1.07$.  In addition to the detected sources, we have an excess in the core of M4 of about 150 counts, uncorrected for absorption, which corresponds to roughly 300 counts corrected for absorption in the 0.5--6.0~keV band (which is the addition of counts in \\xsoft\\ and in \\xhard). In this band, the five faintest sources in the core, CX20, CX21, CX22, CX26 and CX27 have between 14 and 10 counts. A minimum number of sources required to explain the excess counts is found by assuming that each source has 10--14 counts, which gives 22--30 sources.  The total number of $\\sim\\!350$ counts from sources {\\em in the core} of M4 with $\\Lx<L_\\mathrm{r}$ (five detected sources and the excess) is similar to that of the 6 sources in the range (1--10)$L_\\mathrm{r}$. From our remark following Eqs.\\,\\ref{eqL} and \\ref{eqN}, we see that this implies that $\\gamma\\ge0.7$; for lower values of $\\gamma$ the sources with $\\Lx<L_\\mathrm{r}$ do not contribute enough counts to explain the observed number. On the other hand, for $\\gamma=1$ the luminosity function must have a cutoff near 0.1$L_\\mathrm{r}$, because otherwise the sources with $\\Lx<L_\\mathrm{r}$ would produce more counts than is observed. With 6 sources in the range (1--10) $L_\\mathrm{r}$, and for $\\gamma=0.7-1.0$ we have 30--60 core sources in the range (0.1--1)$L_\\mathrm{r}$.  In the core of 47\\,Tuc there are 14 sources in the range (10--100) $L_\\mathrm{r}$, which for $\\gamma=0.7-0.8$ leads to predicted numbers of 70--90 (of which 27 are already detected individually) in the range (1--10) $L_\\mathrm{r}$, and 350--560 in the range (0.1--1) $L_\\mathrm{r}$.  These numbers are compatible with the total number of counts in the core of 47\\,Tuc from the observed sources plus the excess of about 500 counts, estimated by Grindlay et al.\\ (2001a\\nocite{ghem01}).  In Table~\\ref{tab:scaling} we compare the collision numbers and the masses of the cores of M4 and 47\\,Tuc.  We expect that the numbers of cataclysmic variables scales with the collision number, and the number of magnetically active binaries with the mass. In the range (1--10) $L_\\mathrm{r}$ the ratio of numbers of sources in the cores of M4 and 47\\,Tuc suggest that the ratio is set by mass rather than collision number. This in turn suggests that magnetically active binaries dominate the numbers at these luminosities, in agreement with our suggested identifications in M4.  Comparison with NGC\\,6397 is also interesting, as its collision number is somewhat higher, $\\Gamma=2.1$ (normalized to the value of M4) while its total mass (as measured from its absolute magnitude) is somewhat lower than that of M4. Within its half-mass radius, NGC\\,6397 contains 9 cataclysmic variables, 4 magnetically active binaries and 10 as yet unclassified sources (generally of low X-ray luminosity). Within the half mass radius of M4 we find 2 or 3 cataclysmic variables, 12 magnetically active binaries and 15 unclassified sources, of which an estimated 6--12 are background sources. The ratio for the numbers in NGC\\,6397 and M4 of magnetically active binaries is as expected for a scaling with mass. As already noted by Pooley et al.\\ (2003\\nocite{pla+03}), NGC\\,6397 contains rather more cataclysmic variables than expected on the basis of its collision number.  This is reflected in the flat slope of its X-ray luminosity function. If the high number of bright systems in NGC\\,6397 is explained by a higher mass of this cluster in the past (as suggested by Pooley et al.\\ 2003\\nocite{pla+03}), one would expect an accordingly higher number of magnetically active binaries, in contrast to what is observed (assuming that magnetically active binaries have a similar evaporation rate as cataclysmic variables, much lower than the evaporation rate of single stars).  In this respect, it is worth noting that the time scale on which a binary can be destroyed by a close encounter is short in NGC\\,6397 compared to most other clusters, including M4 and 47\\,Tuc (Verbunt 2003\\nocite{ver03}, in particular Figure 3b). Since magnetically active binaries tend to have longer orbital periods then cataclysmic variables, they would be easier destroyed by close encounters. This might explain the absence of large numbers of such binaries in NGC\\,6397."
    },
    "0402/astro-ph0402368_arXiv.txt": {
        "abstract": "We present magnesium abundances derived from high resolution spectra using several Mg I and two high excitation Mg II lines for 19 metal-poor stars with [Fe/H] values between -1.1 and +0.2. The main goal is to search for systematic differences in the derived abundances between the two ionisation state lines. Our analysis shows that the one-dimensional LTE and N-LTE study finds a very good agreement between these features. The [Mg/Fe] vs. [Fe/H] relationship derived, despite the small sample of stars, is also in agreement with the classical figure of increasing [Mg/Fe] with decreasing metallicity. We find a significant scatter however, in the [Mg/Fe] ratio at [Fe/H]$\\sim -0.6$ which is currently explained as a consequence of the overlap at this metallicity of thick and thin disk stars, which were probably formed from material with different nucleosynthesis histories. We speculate on the possible consequences of the agreement found between Mg I and Mg II lines on the very well known oxygen problem in metal-poor stars. We also study the [O/Mg] ratio in the sample stars using oxygen abundances from the literature and find that the current observations and nucleosynthetic predictions from type II supernovae disagree. We briefly discuss some alternatives to solve this discrepancy. ",
        "introduction": "It  is  currently  assumed   that  the  chemical  composition  of  the atmospheres of unevolved (dwarf)  stars reflects that of the material from which the stars are formed. Therefore, abundance studies in these stars covering a wide  range of metallicities is probably the best tool to study the different scenarios proposed for the formation of the Galaxy, its star formation rate history and chemical evolution.  During the  last decade there  has been a considerable  improvement in stellar abundance  analyses.  This fact,  together  with the  more accurate  derivation  of   the  basic  stellar  parameters  (effective temperature, gravity etc),  mainly after the general use  of the {\\it Hipparcos  Catalogue} data  (ESA 1997),  has considerably reduced the uncertainty in  the derivation of  stellar abundances; formal errors  of a few hundredths  of dex in the  abundance ratios are now being obtained!  (see  e.g. Reddy  et  al.  2003).  On  the  other hand,  the classical LTE  abundance analysis using 1D model  atmospheres is being progressively replaced by full N-LTE abundance analysis, in some cases using also  3D hydrodynamical  model atmospheres (Nordlund  \\& Dravins 1990; Asplund et  al. 1999; Asplund \\& Garc\\'\\i  a-P\\'erez 2001). With such strong  tools available,  recent analyses of the  abundance ratios ([X/Fe]) as a function of the stellar metallicity ([Fe/H]\\footnote{The usual bracket notation is used  throughout the paper for the abundance of  an  element  X  with  respect to  the  hydrogen  [X/H]$=  \\rm{log} {{(N_X/N_H)}_\\star\\over{(N_X/N_H)}_\\odot}$, where N represents number density.})  are  showing  a new, detailed structure in the [X/Fe] vs. [Fe/H] trends (on occasions, unexpected) that were hidden in previous analyses, in part due to lower quality data and larger abundance errors (e.g. Gratton  et al.  1996; Nissen \\&  Schuster 1997;  Fuhrmann 1998; Idiart  \\&  Th\\'evenin 2000;  Prochaska et al. 2000; Mashonkina  \\&  Gehren  2001; Gehren  et al.  2004; Gratton et al. 2003a,b). These new studies clarify theories of the Galaxy's formation (see e.g. Weimberg et al. 2002; Samland \\& Gerhard 2003), in particular its star formation history, the timescales for the formation of the halo, thick and thin disks and, in some cases, could lead to a revision of the nucleosynthesis theory in stars, mainly in type II supernova explosions. For instance, concerning    the   well   known    overabundances  of the named $\\alpha$-elements (O, Mg, Si, Ca etc) with respect to Fe in metal-poor stars,  there   are  claims  indicating  that  the  level   of  this overabundance  at a given  metallicity is  different depending  on the stellar  population  studied:  halo,  thick-disk  or  thin-disk  stars, but also for a particular $\\alpha$-element. This poses a difficult challenge to our current understanding of nucleosynthesis in stars. Furthermore, abundance  studies  in  extreme metal-poor  stars ([Fe/H]$<-2.5$) have revealed   very   peculiar  abundance ratios. For instance, there are a significant number of stars with [C,N,Mg,Si/Fe]$>1$), with a considerable  degree of dispersion (Norris, Ryan \\&  Beers 2001;  Depagne et  al. 2000;  Truran et al.  2002; Aoki  et  al. 2002,  among  many others). This might  well provide information about  the level of  mixing of the interstellar  medium at early  epochs as  well as  on nucleosynthesis  processes in  zero metallicity stars.  A common characteristic of many stellar abundance analyses is that they are usually based on the use of a small number of atomic lines for a specific element, in many cases also in the same ionisation state. Of course, for some elements there are very few spectral lines available in dwarf stars, in some cases just one (e.g. the Li I at 6708 {\\AA}) but in others, the use of several abundance indicators does not always lead to agreement between them. In fact, in  this era of precision in abundance analyses, the longstanding problem of the [O/Fe] vs. [Fe/H] relationship in metal-poor stars still remains to  be solved. The {\\it oxygen  problem} is that  several potential  indicators of  the oxygen abundance  (ultraviolet OH lines,  the forbidden  [OI] 6300  {\\AA} and 6363  {\\AA} lines, the  high excitation lines of 7770 {\\AA} O I  triplet, and the near infrared vibration-rotation lines from  the OH  molecule's  ground  state), may  give different results (see the recent discussions by Nissen et al. 2002, and Fulbright \\& Johnson 2003, for details). This fact provokes some doubts regarding our current understanding of line formation in the atmosphere of metal-poor stars and, in consequence, regarding the implications to be drawn from the abundance analyses for the important nucleosynthetic and galactic issues mentioned above.  In this paper, we address the difficulty using an example: magnesium. It is widely  accepted that the [Mg/Fe] ratio in metal-poor  stars increases  for decreasing  metallicity, reaching  a  plateau-like value $\\sim  0.3-0.4$  dex  at  [Fe/H]$\\leq -1$. Magnesium is  an $\\alpha$-element which is produced in  the interior of  massive stars during the hydrostatic burning phases of the stellar evolution. Nevertheless, recent studies have revealed that the [Mg/Fe] ratio shows greater scattering at a given metallicity than other $\\alpha$-elements like, Si, Ca or Ti (we exclude here O); moreover, there is a sudden increase in the [Mg/Fe] ratio at [Fe/H]$\\sim -0.6$ and/or a clear distinction of the [Mg/Fe] ratio depending on which population the star belongs to (see e.g. Prochaska et al. 2000; Qiu et al. 2003 and references therein).  Similarly to oxygen, there are several atomic and molecular Mg lines available to derive the abundance of  this element in dwarfs, including two  Mg II  lines  at 7787  and  7796 {\\AA} not commonly used in the literature and which  have a similar excitation energy ($\\chi=9.99$ eV) to those of the O I triplet ($\\chi=9.14$ eV). From the  point of view  of classical 1D  analysis, this  implies that these  Mg II lines  should form  at essentially  (however, see below)  the same depth in the atmosphere as the oxygen triplet lines.  Furthermore, since neutral oxygen and  single ionised magnesium  are  both  majority species in metal-poor dwarf stars, they are expected to be affected in a similar  way by stellar  inhomogeneities (granulation, Asplund  2003). Furthermore,  given the high  excitation  energy of  these  lines,  they  are expected  to  be as sensitive to the temperature structure of the atmosphere as the oxygen triplet. Departures  from LTE in these  Mg II lines  are also expected, for  the  same  reason. In brief, we wonder whether analysis of these Mg II would give abundances in agreement with those obtained with the more commonly used Mg I lines or, similarly to oxygen, there is also a {\\it  magnesium  problem}.  The  paper is  organised as  follows: Section  2 describes  the observations  and the sample  of stars.  Section  3 presents the abundance analysis  of the selected Mg  I and Mg  II lines considering also  the corrections for N-LTE effects.  In Section  4 the  results are compared  with   previous  studies  and  finally   our  conclusions  are summarised in Section 5. ",
        "conclusions": "We have analysed several Mg I and two Mg II lines in a sample of 19 mildly metal-poor dwarf stars in a search for differences between these two Mg abundance indicators. We found a good agreement between both in the LTE and N-LTE cases. Similar to the features found in other studies, the [Mg/Fe] ratios derived here show significant scatter at [Fe/H]$\\sim -0.6$ which is currently ascribed to the existence of thick and thin disk stars at about this metallicity, stars which are chemically discrete from each other. Using the oxygen abundances from the literature, we also studied the [O/Mg] vs. [Fe/H] relationship. The [O/Mg] ratios derived are constant over the range of metallicity studied but significantly larger than the nucleosynthetic predictions obtained from SNII models. There is no easy explanation to this problem. It could be both spectroscopic and theoretical but, in any case, the determination of any abundance ratio in metal-poor stars urgently requires a detailed modelling (3D) of their atmospheres as well as a simultaneous N-LTE study. Only thus could we safely draw conclusions concerning the chemical evolution of the Galaxy."
    },
    "0402/astro-ph0402642_arXiv.txt": {
        "abstract": "The main properties of the observed extrasolar planets are reviewed with respect to their relevance to the formation scenario of planetary systems. Results of numerical computations of embedded planets in viscously evolving disks are presented. Emphasis is given to the accretion and migration process. New calculations on inviscid disks are shown.  The second part of the talk concentrates on resonant planetary systems. Among the observed extrasolar systems there are 3 confirmed cases, Gl 876, HD 82943 and 55 Cnc, where the planets orbit their central star in a low order mean motion resonance. Results of numerical simulations modeling the formation and evolution of such systems are presented. ",
        "introduction": "Since the first discovery of extrasolar planets around main sequence stars in 1995, their number has risen to over 100 planets today. For an always up-to-date list see e.g. {\\tt {http://www.obspm.fr/encycl/encycl.html}}, maintained by J.~Schneider. Apart from two discoveries by the transit method of OGLE-objects \\citep{2003Natur.421..507K, 2003A&A...402..791D}, all others have been found by radial velocity measurements. Among all systems, there are at least 12 planetary systems with two or more planets; a summary of their properties has recently been given by \\citet{2002marcy-systems}. Longer observations to study trends in velocity curves may still increase this number. Of all stars investigated, planets have been found so far around about 7 \\%.  The orbital properties of these planetary systems display features quite different from those of our own solar system. Some statistical properties are summarized in \\citet{2002RvMA...15..133M}, \\citet{2003A&A...398..363S} and additional data may be obtained from {\\tt {http://www.obspm.fr/encycl/encycl.html}}. As an example we display in Fig.\\ref{fig:eccen} the eccentricity versus mass distribution of all plants detected so far. While the planets close to the star (hot Jupiters) tend to have circular orbits due to tidal circularization the planets further out have a rather broad eccentricity distribution which is essentially indistinguishable from that of binary stars. Additionally, the cumulative eccentricity-period distribution of stellar binaries and planetary systems are also nearly identical for periods larger than about 10 days \\citep{2002RvMA...15..133M}. This may imply a common origin of the eccentricities, caused for example by dynamical gravitational interaction.  \\begin{figure} \\centering \\includegraphics[width=0.95\\linewidth]{e-rx.eps} \\caption{Eccentricity vs. distance for the extrasolar planets. For comparison the solar system planets and the two OGLE objects are also included. \\label{fig:eccen}} \\end{figure} A summary of the main properties of the exo-planets is outlined in the following. The reasons why those characteristics present problems to the standard planet formation scenarios, is given just below the individual points, see also \\citet{2002Natur.419..355L}. \\begin{list}{I}{~} \\item[{\\color{black} $\\bullet$}] Large Masses (0.2-17 M$_{\\mbox{Jup}}$) \\\\ Tidal interaction between the planet and the protoplanetary disk results for planets larger than about 1 Jupiter mass ($M_{Jup}$) in the creation of a density gap in the disk \\citep{1980MNRAS.191...37L, 1993prpl.conf..749L}. In turn, this gap will reduce the accretion onto the planet and limit the mass to about 1 $M_{Jup}$, just as in our solar system \\citep{1993ARA&A..31..129L}. \\item[{\\color{black} $\\bullet$}] Small Distance (0.04 - 3.3 AU)\\\\ The high disk temperature in the vicinity of the star makes it difficult to allow for condensation of solid material. Thus, the cores of the protoplanets are expected to have been formed beyond the so called snow line at a distance of about 3 AU and larger from the star \\citep{2000ApJ...528..995S}. \\item[{\\color{black} $\\bullet$}] {Large Eccentricities} (0.0 - 0.91)  \\\\ It is usually assumed that planets form in a protoplanetary disk, which is essentially in a Keplerian rotation around the star. Hence, directly after formation, planets are expected to orbit on circular orbits. \\item[{\\color{black} $\\bullet$}] {Resonant Orbits} (2:1, 3:1)\\\\ The condensation process in the disk and the subsequent oligarchic \\citep{1998Icar..131..171K, 2003Icar..161..431T} growth of planetesimals towards larger planets can occur theoretically at arbitrary orbits. For stability reasons, only the planets cannot be on orbits separated by less than about 10 Hill radii. \\item[{\\color{black} $\\bullet$}] {Free Floaters} (in Jupiter-mass range) \\\\ Recently, objects in the planetary mass range (a few $M_{Jup}$) but not orbiting stars have been found \\citep{2000MNRAS.314..858L, 2001ApJ...556..830B}. Although not directly part of a planetary system today, the origin of these objects may be related to planet formation, followed by ejection processes. It is believed that bodies of a few $M_{Jup}$ are to small to have formed by direct gravitational collapse. \\end{list} All the raised features may have their origin in an evolution process of the young planets still embedded with the protoplanetary disk from which they formed originally. ",
        "conclusions": "Through high resolution numerical computations of planets still embedded in protoplanetary disks it is possible to calculate the accretion migration rate of planets for the whole mass range from a few earth masses to over 1 Jupiter mass. The inferred timescales can be significantly shorter than the typical lifetimes of the disk.  Thus, we arrive at the problem of stopping the migration close to the stars, as well as for Jupiter in our own solar systems. In the case of the close-in planets it is believed that tidal interaction with the central star possibly through magnetic fields may be responsible for braking the planet. In the case of the solar system, early dissipation of the solar nebula must have occurred to prevent the planets from moving further in.  Even though a lot of progress has been achieved during the last years, several problems will have to be addressed in the future. Among those are inclined orbits which may lead to disk warping (see $\\beta$ Pic disk), the inclusion of radiative transport to study observational effects, planets in binaries to investigate the constraints on the formation. The study of turbulent (MHD) and radiative disks has just begun and possibly will lead in the future to a change of some aspects of the migration and evolution scenario of young planets."
    },
    "0402/astro-ph0402532_arXiv.txt": {
        "abstract": "We present results of numerical simulations of sequences of binary-single scattering events of black holes in dense stellar environments.  The simulations cover a wide range of mass ratios from equal mass objects to 1000:10:10$~\\msun$ and compare purely Newtonian simulations to simulations in which Newtonian encounters are interspersed with gravitational wave emission from the binary. In both cases, the sequence is terminated when the binary's merger time due to gravitational radiation is less than the arrival time of the next interloper.  We find that black hole binaries typically merge with a very high eccentricity ($0.93 \\le e \\le 0.95$ pure Newtonian; $0.85 \\le e \\le 0.90$ with gravitational wave emission) and that adding gravitational wave emission decreases the time to harden a binary until merger by $\\sim30$ to $40\\%$.  We discuss the implications of this work for the formation of intermediate-mass black holes and gravitational wave detection. ",
        "introduction": "Recent observations suggest that large black holes may reside in the centers of some stellar clusters.  X-ray observations in the last few years have shown unresolved sources in galaxies offset from their nuclei and with fluxes that, if isotropic, correspond to luminosities of $L \\approx 10^{39}$ to $10^{41}~{\\mathrm {erg}}\\; {\\mathrm s}^{-1}$ \\citep*[e.g.,][]{fsm97, cm99, metal01, fzm01}.  Many of these sources are associated with stellar clusters (\\citealt{fsm97}; \\citealt*{alm01}).  The strong variability observed in these sources suggests that they are black holes, and if the observed fluxes are neither strongly beamed nor super-Eddington, the implied masses are as high as $M\\ga 10^{3}\\;{M}_{\\odot}$.  The fact the sources are non-nuclear implies masses $M\\la 10^{6}\\;{M}_{\\odot}$ since a larger mass would have rapidly sunk to the center of the host galaxy due to dynamical friction \\citep[$< 10^{9}$~yr for a dispersion velocity of $100\\kms$ and a separation from the galaxy nucleus of $10^{2}~\\mathrm{pc}$ as in the case of M82;][]{ketal01}.  In addition, optical observations of the globular clusters M15 and G1 show velocity profiles consistent with central black holes with masses of $2.5\\times 10^{3}~\\msun$ and $2.0\\times 10^{4}~\\msun$, respectively \\citep*{getal00, getal02, vetal02, grh02}, although \\citet{betal03} demonstrate with their N-body simulations that the observations of G1 can be explained without a large black hole.  Such intermediate-mass black holes (IMBHs) would be in a different mass category, and thus likely indicative of a different formation scenario, from either $3$ - $20~M_{\\odot}$ stellar-mass black holes, which are thought to be the result of core-collapse supernovae, or $10^{6}$ - $10^{10}~M_{\\odot}$ supermassive black holes, which are found in the centers of many galaxies.  Several models have been proposed to account for the origin of IMBHs. \\citet{mr01} and \\citet{setal02} suggest that they are the remnants of massive ($M \\ga 200~\\msun$) Population~III stars.  The low metallicity of these stars precludes cooling through metal line emission and enables them to reach masses much larger than ordinary main sequence stars.  These large stars avoid significant mass loss due to stellar winds or pulsations, and the star may collapse to form a black hole with almost the same mass as the progenitor star.  \\citet{pzm02} and \\citet*{gfr04} show with numerical simulations that the core of a young stellar cluster may collapse rapidly such that direct collisions of stars will lead to runaway growth of a single object with as much as $10^{-3}$ of the original cluster mass over the course of a few million years.  \\citet{mh02} propose that over a Hubble time stellar-mass black holes in dense globular clusters may grow by mergers to the inferred IMBH masses.  In their model, a black hole with mass greater than $50~\\msun$ will interact with other massive objects to form binaries that will merge due to gravitational radiation.  The merger process may proceed more quickly in the presence of encounters with a third black hole or another black hole binary \\citep{mh02b} if the encounters shrink the binary's orbit, as is known to happen with hard (tight) binaries \\citep{h75}.  Wherever and however IMBHs formed, the best candidates are found in stellar clusters where three-body encounters are important.  An IMBH in a cluster, whether formed there or later swallowed by the cluster, will find its way to the center.  As all of the heaviest objects in a cluster sink to the center in a process known as mass segregation, the IMBH will interact primarily with other massive objects and binaries \\citep{sh93, fetal02}.  A single IMBH will tend to acquire companions through exchanges with binaries because the most massive pair of objects in a three-body encounter preferentially end up in the binary \\citep*[e.g.,][]{hhm96}.  The IMBH binary will encounter other objects in the dense center of its host cluster, harden further, and ultimately merge.  These events are important sources of gravitational waves.  The Advanced LIGO (Laser Interferometer Gravitational Wave Observatory) detector is expected to be capable of detecting mergers of IMBHs with $M \\la 100~\\msun$ \\citep*{b00}, and LISA (Laser Interferometer Space Antenna) is expected to detect the earlier inspiral phase of an IMBH merger \\citep*{d00}.  In order to predict the gravitational wave signature of the inspiral, the expected separations and eccentricities of the binaries must be known.  Because three-body encounters alter the orbital parameters, simulations are needed to predict their distributions as well as the source population and event rates.  The three-body problem has been studied extensively, but with every new generation of computing power, our understanding of the problem advances with a wider range of numerical simulations and a changing perspective on this rich but conceptually simple problem.  Previous studies of the three-body problem have tended to focus on the case of equal or nearly equal masses \\citep[e.g.,][]{h75, hb83} though other mass ratios have been studied \\citep[e.g.,][]{fh82, sp93, hhm96}.  The nearly equal mass case does not apply to the case of an IMBH in the core of a stellar cluster.  In addition the vast majority of previous work has studied the effect of a single encounter on a binary.  To determine the ultimate fate of an IMBH, simulations of sequences of encounters are needed. Furthermore, to our knowledge no previous work has considered the effects of orbital decay due to gravitational radiation between encounters, which we expect to be important for very tight binaries.  In this paper we present numerical simulations of sequences of high-mass ratio binary-single encounters.  We describe the code used to simulate the encounters in \\S~\\ref{nummeth}.  Next, we present results of the simulations of sequences of encounters on a range of mass ratios with Newtonian gravity (\\S~\\ref{newt}) and with gravitational radiation between encounters (\\S~\\ref{grseq}) and show that including gravitational radiation decreases the duration of the sequence by $\\sim 30$ to $40\\%$.  In \\S~\\ref{imbhformation} and \\S~\\ref{gwdetection} we discuss the implications of these results for IMBH formation and gravitational wave detection. ",
        "conclusions": "\\label{conclusions} We present results of numerical simulations of sequences of binary-single black hole scattering events in a dense stellar environment.  We simulate three-body encounters until the binary will merge due to gravitational radiation before the next encounter. In half of our simulations, we include the effect of gravitational radiation between encounters.   \\emph{1. Sequences of high mass ratio encounters.}  Our simulations cover a range of mass ratios including those corresponding to IMBHs in stellar clusters.  Because the binaries simulated are tightly bound, the encounters steadily shrink the binary's semimajor axis until it merges.  The eccentricity, however, jumps chaotically between high and low values over the course of a sequence.  Merger usually occurs at high eccentricity since gravitational radiation is much stronger then.    \\emph{2. Gravitational wave emission between encounters.}  The inclusion of gravitational radiation between encounters affects the simulations in several ways.  The extra source of shrinking caused by gravitational wave emission has the effect of shortening the sequence in terms of both the number of encounters and the total time, and the circularization from gravitational waves has the effect of decreasing the final eccentricity of the binary before it merges.  \\emph{3. IMBH formation.}  Our simulations directly test the IMBH formation model of \\citet{mh02}.  We find that there is sufficient time to build up to $1000~\\msun$ when starting from $50~\\msun$, but our simulations also show that if there are a thousand $10~\\msun$ black holes in the globular cluster, the seed black hole would only be able to grow to $240~\\msun$ before exhausting half of the black holes in the cluster.  In addition, the probability of the binary's remaining in the cluster during a growth from $50$ to $240~\\msun$ is small.  In order to avoid ejection from the cluster with a reasonable probability, either the black hole must have a larger mass at the onset of dynamical encounters, the cluster's escape velocity must be larger, or the black hole must grow by some additional mechanisms such as by Kozai-resonance-induced mergers, two-body captures, or from smaller interlopers.  \\emph{4. Gravitational wave detection.}  The mergers of binary black hole systems are strong sources of detectable gravitational waves.  We find that the merging binary will typically start with very high eccentricity.  By the time the binary is detectable by the Advanced LIGO detector, it will have completely circularized, but when detectable by LISA, it may have moderate eccentricity ($0.1 \\la e \\la 0.2$) such that it will display general relativistic effects such as pericenter precession and still possibly be detectable with circular templates.  We find a high rate of mergers in the first few hundred million years of a globular cluster.  This suggests that recently formed, nearby super star clusters are promising sources for gravitational radiation from IMBH coalescence.  Further work in this study will be to include a distribution of interloper masses instead of a single population of $10~\\msun$ black holes.  A mass distribution of black holes is a more realistic model of a cluster core and could change the outcomes of the sequences. Exchanges will be more important since encounters with the more prevalent smaller black holes may do most of the hardening until a more massive black hole exchanges into the binary."
    },
    "0402/astro-ph0402656_arXiv.txt": {
        "abstract": "A survey of the northern hemisphere for astrophysical point sources with continuous emission of high-energy cosmic rays is presented. Around 47 Mio extensive air showers with primary energies above $\\approx$ 300~TeV measured by the KASCADE detector field are selected for this analysis. Besides the sky survey, a search for signal excess in the region of the galactic plane and of selected point source candidates has been performed. There is no evidence for any significant excess. This is valid for an analysis of all recorded showers as well as for a data set enhanced in gamma-ray induced showers. An upper flux limit of around $3\\times 10^{-10}$ m$^{-2}$s$^{-1}$ for a steady point source that transits the zenith is obtained. Additionally, the distribution of the arrival directions of extensive air showers with energies above 80~PeV were studied by an autocorrelation analysis. ",
        "introduction": "The origin of high-energy cosmic rays is still unknown. Their sources are obscured due to the deflection of the charged particles in galactic magnetic fields. Only neutral particles, i.e.~neutrons or high-energy photons, can point back to their sources. But besides restrictions from the known acceleration processes, several limitations for the detection of neutral particles arise from propagation effects and from the large background of charged cosmic rays.  Neutrons can reach the Earth if their energy and hence their decay length is comparable with the distance of the source. A decay length of 1~kpc corresponds to a neutron energy of about $10^{17}$~eV, the distance to the galactic center (8.5~kpc) to an energy of roughly $10^{18}$~eV. There are models proposing neutron emission from the center of the galaxy \\citep{Hayashida:1999}. They could explain the excess of events from this direction (which is unfortunately not visible from the \\mbox{KASCADE} site) measured by the AGASA experiment \\citep{Hayashida:1999} at around $10^{18}$~eV. In the present work the possibility of neutrons originating from more nearby sources ($\\approx 1$~kpc) is analysed with the data set of KASCADE.  TeV $\\gamma$-ray sources are potential sites for the acceleration of high-energy cosmic rays. Several sources have been established in the last decade. The energy spectrum of some of them extend up to 80~TeV (e.g. \\cite{Tanimori:1998b}). Limitations for the detection of PeV $\\gamma$-ray sources are: very low fluxes due to the steeply falling energy spectra of the sources, a possible cutoff in their spectra, and the attenuation of $\\gamma$-rays on their way to the observer as a result of the interaction with low energy photons from the cosmic microwave background radiation \\citep{Coppi:1997}. A review of the status of high-energy $\\gamma$-ray astronomy can be found in \\citep{Ong:1998,Hoffmann:1999}.  All TeV $\\gamma$-ray sources have been observed by imaging atmospheric Cherenkov telescopes. % While these instruments are very sensitive to $\\gamma$-ray showers (low energy threshold, good rejection of hadron-induced extensive air showers), the major drawback for the search of new sources is their small field of view and their limitation to dark observation periods. Large field arrays measuring the secondary particle distribution of extensive air showers on ground have large fields of view and duty cycles independent of day and night periods. Unfortunately, they suffer from the huge background of hadron-induced extensive air showers. Only the Tibet air shower array and the Milagro experiment reported the detection of a TeV $\\gamma$-ray source (Crab-Nebula) \\citep{Amenomori:1999,Atkins:2003}.  Several sky surveys for steady $\\gamma$-ray sources have been performed by air shower experiments at energies from 1~TeV to 1~PeV (CYGNUS \\citep{Alexandreas:1991b}, CASA \\citep{McKay:1993}, Milagrito \\citep{Wang:2001}, Tibet \\citep{Amenomori:2002}, HEGRA \\citep{Aharonian:2002}). None of the experiments have found point sources. In this article a survey of the northern sky for primary energies above 300~TeV in the declination range 15\\Grad \\ $ < \\delta <$ 80\\Grad \\ with data recorded by the KASCADE detector field is presented. Compared to the experiments mentioned above, KASCADE measures in a different energy region or observes a different part of the sky. An analysis of the arrival-direction distribution of charged cosmic rays on large angular scales covering the effects of their diffusion in the galactic magnetic field is presented elsewhere \\citep{Antoni:2003c}. ",
        "conclusions": "A search for small scale anisotropies in a data set of about 47 Mio extensive air showers measured with the KASCADE experiment has been presented. The arrival direction of all events in this data set as well as a subset of muon-poor events, i.e.~extensive air showers which are more similar to $\\gamma$-ray induced showers, are analysed. No evidence for any point-like source has been found, which is in accordance with theoretical expectations. This is valid for the sky survey as well as for a detailed analysis of a region around the galactic plane and known TeV $\\gamma$-ray sources. Upper limits for the detection of point-like sources are determined to be around $10^{-10}$ m$^{-2}$ s$^{-1}$. Also no clustering of the arrival direction for showers with primary energies above 80 PeV is visible."
    },
    "0402/astro-ph0402183_arXiv.txt": {
        "abstract": "The {\\it COS~B} high energy $\\gamma$-ray source 2CG~075+00, also known as GeV~J2020+3658 or 3EG~J2021+3716, has avoided identification with a low energy counterpart for over twenty years. We present a likely identification with the discovery and subsequent timing of a young and energetic pulsar, PSR J2021+3651, with the Wideband Arecibo Pulsar Processor at the Arecibo Observatory.  PSR~J2021+3651 has a rotation period $P \\sim 104$~ms and $\\dot P \\sim 9.6\\times10^{-14}$, implying a characteristic age $\\tau_c \\sim 17$ kyr and a spin-down luminosity $\\dot E \\sim 3.4\\times 10^{36}$ erg~s$^{-1}$.  The pulsar is also coincident with the {\\it ASCA} source AX J2021.1+3651. The implied luminosity of the associated X-ray source suggests the X-ray emission is dominated by a pulsar wind nebula unresolved by {\\it ASCA}. The pulsar's unexpectedly high dispersion measure (DM~$ \\sim 371$~pc cm$^{-3}$) and the $d \\ge 10$~kpc DM distance pose a new question: is PSR J2021+3651 an extremely efficient $\\gamma$-ray pulsar at the edge of the Galaxy?  This is a question for {\\it AGILE} and {\\it GLAST} to answer. ",
        "introduction": " ",
        "conclusions": "The DM of PSR J2021+3651 is by far the highest known in the Galactic longitude range $55^{\\circ} < l < 80^{\\circ}$, which is mainly an inter-spiral arm direction. Using the recent \\citet{cl02} update to the \\citet{tc93} dispersion measure model gives a distance of $\\sim 12.4$ kpc, at the outer edge of the last spiral arm used in the model. It is possible that there are further contributions from clouds in the Cygnus region not included in the model, where there is known to be excess gas at $d\\sim 1.5$~kpc (J. Cordes, private communication, 2002).  However, there are no obvious HII regions within the Arecibo beam seen in either Very Large Array (VLA) 20-cm radio or Midcourse Space Experiment (MSX) 8.3~$\\mu$m images (available from the NASA/IPAC Infrared Science Archive).  The likelihood that PSR J2021+3651 is unrelated to the X-ray source is extremely small given the rarity of such young, energetic, pulsars and the consistency of the timing position with the $\\sim 1'$ positional uncertainty of the {\\it ASCA} source. However, the high DM is somewhat surprising given the X-ray absorption quoted by RRK (n$_{\\rm H}$=(5.0$\\pm0.25)\\times 10^{21}$~cm$^{-2}$), where the errors represent the 90\\% confidence region. The total Galactic HI column density in this direction as estimated from the FTOOL {\\tt nh}, which uses the HI map of \\citet{dl90}, is $1.2\\times 10^{22}$~cm$^{-2}$.  This should be a good approximation if the source is truly at the far edge of the outer spiral arm. Noting that the {\\it ASCA} image shows faint, softer emission in the region (Figure~\\ref{fig:xray}), and given the likely possibility of either associated thermal X-ray flux from a supernova remnant or a nearby massive star, we fit the {\\it ASCA} spectrum of RRK, adding a thermal component to the absorbed power-law model. Accounting for $\\sim4$\\% of the photon flux with a MEKAL thermal plasma model (see \\citet{log95} and references therein) of temperature $kT\\sim 0.1$~keV in {\\tt XSPEC} \\citep{arn96} statistically improves the fit ({\\it F}-test chance probability of 2.5\\%). The best-fit absorption for this three component model is n$_{\\rm H}$=7.6$\\times 10^{21}$~cm$^{-2}$ with a 90\\% confidence region of (4.1 -- 12.3)$\\times 10^{21}$~cm$^{-2}$, consistent with the total Galactic column density. The best-fit photon index is $\\Gamma=1.86$, still consistent with the 1.47 -- 2.01 range in RRK derived from the simple absorbed power-law model.  Hence the X-ray absorption does not force us to adopt a smaller distance than is suggested by the DM and the value of n$_{\\rm H}$ may be contaminated by thermal emission. Upcoming {\\it Chandra} observations of the source may well clarify this.  For a distance $d_{10}=d/10$~kpc, the inferred isotropic X-ray luminosity $L_{\\rm X}=4.8\\times 10^{34} d_{10}^2$ erg s$^{-1}$ (2 -- 10~keV). The X-ray efficiency $\\eta_{\\rm X}=L_{\\rm X}/\\dot E$ is 0.01$d^2_{10}$. Compared to the total pulsar plus nebula X-ray luminosity of other spin-powered pulsars, this is somewhat high, but within the observed scatter \\citep{pccm02,che00}. Upcoming Chandra observations will determine what fraction of the X-ray flux comes from a compact nebula, and will allow us to search for X-ray pulsations.  2CG~075+00 = 3EG J2021+3716 = GeV J2020+3658 is the 10th brightest Galactic source above 1 GeV. It has a hard (photon index $\\Gamma=1.86\\pm 0.10$) spectrum with no sign of a break out to 10 GeV, and has low variability, similar to the known $\\gamma$-ray pulsars. PSR J2021+3651's position is outside the 99\\% contour of the 3EG catalog position, but consistent with the 95\\% contour of both the \\citet{lm97} GeV catalog position and the RRK position. The only other potential counterpart within any of the error contours is WR 141. Although such binary systems have the potential for producing $\\gamma$-rays from colliding winds, no convincing association with a known high-energy ($>$ 100~MeV) $\\gamma$-ray source has yet been made for a Wolf-Rayet system \\citep{rbt99}. The pulsar's positional coincidence coupled with the high inferred spin-down luminosity strongly suggests this pulsar emits $\\gamma$-rays. Unfortunately, confirming this by folding archival {\\it EGRET} data is problematic due to the likelihood of significant past timing noise and glitches, which make the back-extrapolation of the rotational ephemeris uncertain.  Assuming the distances given by the \\citet{cl02} model, the known $\\gamma$-ray pulsars have $\\gamma$-ray efficiencies $\\eta_\\gamma=L_{\\gamma}/\\dot E$ mostly between 0.0001 and 0.03 (assuming 1 sr beaming). Assuming the $\\gamma$-rays are 100\\% pulsed, the inferred $\\gamma$-ray efficiency for PSR J2021+3651, $\\eta_\\gamma =0.15 d_{10}^2$ in the 100~MeV to 10~GeV range, would be by far the most efficient $\\gamma$-ray pulsar.  Note however that PSR B1055$-$52 \\citep{tbb+99} has a similar efficiency if one uses the DM distance given by \\citet{tc93}, instead of the newer \\cite{cl02} model. PSR B1706$-$44, which has nearly identical spin parameters, has a $\\gamma$-ray efficiency of $\\sim 0.01$, while Vela with a similar spin-down energy has a low efficiency of $\\sim 0.0004$. Even if the pulsar is located within the Perseus arm at a distance of 5~kpc, it would still be very efficient. If it is the $\\gamma$-ray source, it  would further strain the relationship of efficiency to spin-down energy $\\eta_{\\gamma}\\propto L_{sd}^{-1/2}$ expected from theory \\citep{zh00}. While there is currently no observational evidence for a closer distance, increased DM from an intervening source in this relatively crowded direction would not be surprising. We note that the DM derived distances for several other young pulsars recently discovered within {\\it EGRET} error boxes, if the true counterparts, also tend to have high inferred $\\gamma$-ray efficiencies \\citet{rob02}.  We plan to continue timing observations of this source, which will refine the position, provide a contemporaneous timing solution for future high-energy pulse searches, and allow us to monitor for glitches.  {\\it AGILE} and {\\it GLAST} should be able to unambiguously determine if the pulsar is the counterpart of the $\\gamma$-ray source."
    },
    "0402/astro-ph0402460_arXiv.txt": {
        "abstract": "In this paper we present the $V$ and $I$ luminosity functions and color-magnitude diagrams derived from wide-field ($23\\arcmin\\times23\\arcmin$) $BVI$ photometry of the intermediate metallicity ([Fe/H]$\\sim-1.3$) Galactic globular cluster M12. Using observed values (and ranges of values) for the cluster metallicity, reddening, distance modulus, and age we compare these data to recent $\\alpha$-enhanced stellar evolution models for low mass metal-poor stars. We describe several methods of making comparisons between theoretical and observed luminosity functions in order to isolate the evolutionary timescale information the luminosity functions contain. We find no significant evidence of excesses of stars on the red giant branch, although the morphology of the subgiant branch in the observed luminosity function does not match theoretical predictions in a satisfactory way. Current uncertainties in $T_{\\mbox{eff}}$-color transformations (and possibly also in other physics inputs to the models) make more detailed conclusions about the subgiant branch morphology impossible. Given the recent constraints on cluster ages from the WMAP experiment (Spergel et al. 2003), we find that good fitting models that do not include He diffusion (both color-magnitude diagrams and luminosity functions) are too old (by $\\sim1-2$ Gyr) to adequately represent the cluster luminosity function. The inclusion of helium diffusion in the models provides an age reduction (compared to non-diffusive models) that is consistent with the age of the universe being $13.7\\pm0.2$ Gyr (Bennett et al. 2003). ",
        "introduction": "One of the observational tools available for the study of low-mass ($\\sim0.5-1~M_{\\odot}$), metal-poor stars is the luminosity function (LF) of Galactic globular clusters (GGCs).  A GGC usually presents a large, chemically homogeneous, and coeval stellar population --- samples of a kind that cannot easily be extracted from halo field stars. In particular, the LF counts of evolved stars (from the main-sequence turnoff to the tip of the red giant branch) are directly related to the rate of nuclear fuel consumption \\citep{re88}.  Thus, a high-precision LF indirectly reflects interior stellar physics, and thus can complement observations of stellar surface conditions.  Previous LF studies of GGCs have given hints that non-standard physics might be required in the models.  First, in the case of several metal-poor clusters an excess of stars on the sub-giant branch (SGB) has been noted.  In a study that combined the LFs of M68, NGC 6397, and M92, \\citet{st91} found an excess located just brighter than the main-sequence turnoff (MSTO).  A similar excess was also observed in M30 by \\citet{bo94} and confirmed by \\citet{be96}, \\citet{gu98}, and \\citet{sa99}.  Such an excess could be the result of enhanced energy transport in the cores of main-sequence (MS) stars nearing hydrogen exhaustion (e.g. Faulkner \\& Swenson 1993).  To date excesses have only been seen in extremely metal-poor clusters and at low statistical significance; the LFs of M5 \\citep{sand96} and M3 \\citep{rood99} did not reveal such excesses. Second, studies by \\citet{st91}, \\citet{bv92}, \\citet{bo94}, \\citet{be96}, and \\citet{sa99} have shown that there may exist an over-abundance of red giant branch (RGB) stars relative to the number of MS stars in some clusters. This may be indicative of physical processes such as core rotation \\citep{va98} and/or deep mixing \\citep{la00}.  The observed LFs of other GGCs, however, have been argued to agree with theory \\citep{de97,zo00,rood99}. While the LF of M3 derived by \\citet{rood99} is one of the largest to date, the study by \\citet{zo00} has a smaller overall statistical significance due to the smaller overall star samples.  In general, only thorough studies of large star populations in GGCs will help confirm or deny the reality of these kinds of excesses.  A physical effect that is becoming part of standard stellar models is the gravitational settling of heavy elements \\citep{ri02}. Solar models show \\citep{pr94,ri96,ba97,gu97} that helioseismic data and the inferred radius of the convection zone \\citep{cgt91} can only be matched by theory if the surface abundance of helium has decreased with time.  Stellar evolution models for low mass, metal-poor stars that include the effects of He diffusion have been previously constructed \\citep{pr91,ch92,st97,ch01} to investigate the observational effects on GGC LFs and color-magnitude diagrams (CMDs). One important result of these studies was that by adding He to the core of the stars, and consequently displacing H, the duration of the MS lifetime is shortened.  This results in a lower MSTO luminosity for a given age and has strong implications for the derivation of GGC ages \\citep{ch92,va02}.  Using these diffusive models GGC ages may be reduced as much as $10-15$\\% ($\\sim1-2$ Gyr) compared to models that do not incorporate diffusion \\citep{st97,va02}, although the reduction could be as little as $4-7$\\% ($\\sim0.5-1$ Gyr) depending on the presence of complete or partial ionization in the models \\citep{ri02,gr03}. Given the recent analysis of the WMAP data from the cosmic background  radiation, we now have a tight cosmological upper-limit on the age of the universe of $13.7\\pm0.2$ Gyr \\citep{be03,sp03}.  As a consequence, theoretical stellar evolution models that do not include diffusion (at least for He) may produce globular cluster isochrones that are too old.  In this paper we present $BVI$ photometric data derived from wide-field CCD photometry of the GGC M12 (NGC 6218, C 1644-018).  This bright, large, intermediate metallicity ([Fe/H]$\\sim1.3$) cluster should provide an interesting comparison to other well-studied intermediate clusters such as M3 and M5.  M12 is also an extreme ``second parameter'' cluster [$(B-R)/(B+V+R)=0.92$; \\citet{le94}]. Ground-based photometric data on M12 have most recently been presented by von Braun et al. (2002; hereafter VB02), Rosenberg et al. (2000; hereafter R00), and Brocato et al. (1996; hereafter B96). Space-based data on M12 was published by \\citet{pi02} as part of a \\textit{HST} GGC snapshot survey. Sato, Richer, \\& Fahlman (1989; hereafter S89) presented a deep CMD and LF for the inner regions of M12, but to date no LF of the evolved stellar populations of M12 has been determined.   In this study we focus on comparing our data (in the form of LFs and CMDs) to three sets of stellar evolution models.  In $\\S2$ we discuss the observations, data reduction and photometry, photometric calibration, and comparisons to existing photometry. In $\\S3$, we present the photometry in the form of CMDs and derived fiducial lines. $\\S4$ discusses the cluster reddening, distance modulus, age, and metallicity which are the necessary input parameters to the comparison of the theoretical LF with the observed.  The computation of the observed LF and incompleteness corrections are discussed in $\\S5$.  In $\\S6$ we compare the data to theoretical CMDs and LFs.  Our conclusions are presented in $\\S7$. ",
        "conclusions": "In this paper we have presented the $VI$ luminosity functions (LFs) and $BVI$ color-magnitude diagrams (CMDs) of the Galactic globular cluster (GGC) M12 from wide-field CCD photometry. Given constraints on the cluster age, metallicity, distance modulus, and reddening we compare our data to three sets of theoretical stellar evolution models for metal-poor, $\\alpha$-enhanced, low mass stars.  We find that neither the \\citet{bv01} nor the \\citet{ki02} models are able to adequately describe the SGB region of the M12 LF.  While we find no statistically significant excesses of stars, the observed SGB LF has a noticeably different slope than predicted.  We find the theoretical description of the SGB region of the cluster LF to be sensitive to the selection of color-$T_{\\mbox{eff}}$ transformations, and to a lesser degree to age and metallicity.  On the other hand, we find agreement between the observed and predicted numbers of MS-to-RGB stars; M12 does not appear to have an excess of RGB stars compared to theory.  In the context of the Langer et al. (2000) claim that extremely blue (``second parameter'') clusters are explained by deep mixing (during the RGB phase) and resulting envelope helium enrichment, the M12 LF should have shown this RGB excess.  In contrast to this, we find the LF to be similar to that of M3 [$(B-R)/(B+V+R)=0.08$; Lee et al.(1994)], another cluster of nearly identical metallicity that does not shown an excess of RGB stars. We find in our analysis that regardless of the differences in input physics in these two models, the adoption of the ZW metallicity scale is incompatible with observations. Assuming a metallicity for M12 on the ZW scale, adequate theoretical description of both the CMD and LF data would require either (1) a model older than recent estimates of the age of the universe (see below), or (2) a distance modulus that would be $2-3\\sigma$ larger than we determined from subdwarf fitting.   Analysis of the WMAP experimental data \\citep{be03,sp03} has now placed new restrictions to the age of the universe, and hence the possible ages of GGCs.  Taking into account a GGC formation timescale of $\\sim1$ Gyr this implies that the possible ages of GGCs can be no larger than $\\sim13$ Gyr.  We find that the models of \\citet{bv01} require the use of a 16 Gyr model in order to account for the observed properties of the M12 LF given the metallicity of M12 on the ZW scale.  A 14 Gyr model would provide a similar fit but require the use of a distance modulus that is $\\sim0.15$ greater (just above our $1\\sigma$ error) than the value we derived from subdwarf fitting. While comparisons between observations and the models of \\citet{bv01} have been shown to prefer the ZW metallicity scale \\citep{bv01,va02}, use of the CG scale with these models still implies an age for M12 of 13-14 Gyr. In contrast, the theoretical models of \\citet{ki02} and \\citet{va02} imply a cluster age of $12-13$ Gyr given the metallicity of M12 on the CG scale, in better agreement with the WMAP upper limit.  We attribute this to the use of He diffusion and gravitational settling (only VandenBerg et al. 2002), although uncertainties due to color-$T_{\\mbox{eff}}$ transformations are still of comparable importance. Previous work (see $\\S1$) has already shown the input of diffusion tends to reduce the cluster ages by $0.5-2$ Gyr, and therefore provides the simplest explanation of our observations.  Clearly for gravitational settling to become a necessary part of stellar evolution models, confirmation of this age reduction and consistency with LF observations will be needed for other clusters.  As a consequence, the reduction of systematic errors in the distance modulus determination (and therefore cluster metallicity and reddening) will be crucial to this analysis \\citep{gr03}. Further observations of stellar surface conditions, such as the $^7$Li Spite plateau \\citep{sp82,ry99} and metal abundance variations in the GGC evolved populations, will also help to place crucial constraints on gravitational settling as well \\citep{va02,ch01}."
    },
    "0402/nucl-th0402012_arXiv.txt": {
        "abstract": "\\noindent Coulomb dissociation is an especially simple and important reaction mechanism. Since the perturbation due to the electric field of the (target) nucleus is exactly known, firm conclusions can be drawn from such measurements. Electromagnetic matrixelements and astrophysical $S$-factors for radiative capture processes can be extracted from experiments. The dissociation of neutron halo nuclei is studied in a zero range model using analytical methods. Of special interest for nuclear structure physics is the appearence of low lying electric dipole strength in neutron rich nuclei. We use effective range methods to study it. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0402/astro-ph0402306_arXiv.txt": {
        "abstract": "We have undertaken a long-term project, Planets in Stellar Clusters Extensive Search (PISCES), to search for transiting planets in open clusters. NGC~2158 is one of the targets we have chosen -- an intermediate age, populous, rather metal poor cluster. In this paper we present the results of a search for variable stars in the data from the first season of monitoring at the FLWO 1.2 m telescope. This is the first variability search ever conducted in this cluster. We present a catalog of 57 variable stars, most with low amplitude variability. Among the variables is a cataclysmic variable (CV) which underwent a 2.5 mag outburst. If it is a member of NGC~2158, this would be the fourth CV known in an open cluster. We have also found five $\\delta$ Scuti stars, three of which we have two or more detectable modes of pulsation. Of the 57 variables discovered, 28 have $R$-band amplitudes of 5\\% or below. Six of those vary at or below the 2\\% level, including one with 0.08\\% variability. ",
        "introduction": "}  We have undertaken a long-term project, Planets in Stellar Clusters Extensive Search (PISCES), to search for transiting planets in open clusters. In Mochejska et al.\\ (2002, hereafter Paper~I) we demonstrated that we are capable of attaining the requisite photometric precision, based on one season of data for NGC~6791.  In this paper we present the results of a search for variable stars in our second target, NGC~2158, based on the data from the first observing season.  NGC~2158 is a rich open cluster of intermediate age (2-3 Gyr). It is rather metal poor ([Fe/H]=-0.46) and is located at a distance of 3600 pc (Carraro et al.\\ 2002, hereafter Ca02; Christian, Heasley and Janes 1985). Despite being one of the most populous galactic open clusters, it has never been investigated for variability. The only published CCD photometry for this cluster is that of Christian et al.\\ (1985), Piersimoni et al.\\ (1993) and Ca02. An extensive catalog of photographic photometry and proper motions has been published by Kharchenko, Andruk \\& Schilbach (1997, hereafter KAS97).  The paper is arranged as follows: \\S 2 describes the observations, \\S 3 summarizes the data reduction procedure and \\S 4 contains the variable star catalog. Concluding remarks are found in \\S 5. ",
        "conclusions": "}  In this paper we present $VR$ photometry for 57 newly discovered variable stars in the intermediate age, populous, rather metal poor cluster NGC~2158. We have identified 34 eclipsing binaries, 19 other periodic variables and 4 long period or non-periodic variables. Among these variables we have discovered five $\\delta$ Scuti stars. For three of them we have identified two or more modes of pulsation.  We have also found a cataclysmic variable which underwent a 2.5 mag outburst.  If it is a member of NGC~2158 (which at this point does not seem very likely) this would be the fourth CV known in an open cluster.  Of the 57 variables discovered, 28 have $R$-band amplitudes of 5\\% or below. Six of those vary at or below the 2\\% level, including one with 0.08\\% variability. This, together with the results from Paper~I, shows that we are capable of attaining requisite photometric precision to detect variability at the level expected for planetary transits."
    },
    "0402/astro-ph0402240_arXiv.txt": {
        "abstract": "We present optical and near-infrared follow-up observations of the X-Ray Flash (XRF) of July 23 2003. Our observations in the R-band cover the temporal range from 4.2 h to 64 days after the high energy event. We also present the results of multicolor imaging extending to the K-band on three epochs. The light curve of the R-band afterglow the first week after the burst is similar to the light curve for long duration Gamma-Ray Bursts (GRBs), i.e.,\\ a broken power-law with a late time slope of $\\alpha\\approx2.0$ ($F_{\\nu} \\propto t^{-\\alpha}$). Furthermore, the spectral energy distribution (SED) has a power-law ($F_{\\nu} \\propto \\nu ^{-\\beta}$) shape with slope ${\\beta}\\approx1.0$. However, the decay slope at $t<1$ day is shallow, consistent with zero. This is in qualitative agreement with the prediction that XRFs are off-axis classical GRBs. After the first week there is a strong bump in the light curve, which peaks at around 16 days. The SED after the peak becomes significantly redder. We discuss the possible interpretations of this bump, and conclude that an underlying supernova is the most likely explanation since no other model appears consistent with the evolution of the SED. Finally, we present deep spectroscopy of the burst both in the afterglow and in the bump phase. A firm upper limit of $z=2.3$ is placed on the redshift of XRF~030723 from the lack of Ly$\\alpha$ forest lines in the spectrum of the afterglow. The lack of significant absorption and emission lines in either of the two spectra excludes a spectroscopic redshift determination. ",
        "introduction": "X-Ray Flashes (XRFs) are transient sources of X-ray photons which are distributed iso\\-tro\\-pi\\-cal\\-ly on the sky. Their existence as a class with a different nature than X-Ray Bursters in the Galaxy was first put forward by Heise et al.\\ (2001) based on data from the BeppoSAX satellite. They introduced an operational definition for XRFs, namely a fast transient source with duration less than 1000~s in the Wide Field Camera (covering 2--25~keV) that did not trigger the Gamma Ray Burst Monitor (covering 40--700~keV). BeppoSAX detected 20 such events in the 6 years it was operational (Heise, in 't Zand, \\& Kippen 2003). A HETE-2 burst is classified as an XRF if its X-ray fluence ($S$) exceeds its gamma-ray fluence, i.e., $\\log [ S($2--30~keV)/$S$(30--400~keV)$] >0$ (Lamb, Donaghy, \\& Graziani 2003). Classical Gamma-Ray Bursts (GRBs) are defined as having $\\log [ S($2--30~keV)/$S$(30--400~keV)$] < -0.5$, while bursts belonging to the intermediate class are classified as X-ray rich GRBs, following Castro-Tirado et al.\\ (1994).  XRFs can be interpreted as the same phenomenon as classical GRBs with the difference being that the high energy spectra for XRFs are softer than for GRBs (Heise et al.\\ 2001; Barraud et al.\\ 2003). The high energy spectra ($\\nu F_{\\nu}$) of GRBs are well described by the so called Band function (Band et al.\\ 1993), which is composed of two smoothly connected power-laws. The energy at which the two power-laws connect is referred to as $E_\\mathrm{peak}$ and is where most of the energy usually is emitted. For classical GRBs $E_\\mathrm{peak}$ is typically a few 100~keV \\cite{preece00}. The high energy spectra of XRFs are also well fitted by the Band function, but with values of $E_\\mathrm{peak}$ below 100~keV and in some cases even below 10~keV \\citep{kippen02,barraud03}. For classical GRBs there is a correlation between spectral hardness, defined as the fluence in the X-ray band divided by the fluence in the gamma-ray band, and the fluence in the gamma-ray band in the sense that hard bursts have a higher total fluence \\cite{nemiroff94}. XRFs follow this correlation \\cite{barraud03}. These properties of GRBs and XRFs would be naturally explained if XRFs were GRBs at very large redshifts. However, the lack of excessive time dilation compared to GRBs argues against this interpretation \\cite{heise01}. Moreover, Amati et al.\\ (2002) find in a sample of 12 BeppoSAX GRBs with known redshifts that the total isotropic equivalent energy radiated in the 1--10000~keV range is positively correlated with $E_\\mathrm{peak}$. This also argues against very large redshifts for XRFs. Instead the current view is that XRFs are either {\\it i)} the result of classical GRBs seen off-axis (Yamazaki et al.\\ 2002, 2003; Dado et al.\\ 2003; Rhoads 2003), {\\it ii)} so-called dirty fireballs, which are relativistic jets with a larger baryon load and hence (assuming external shocks) lower $\\Gamma$-factor than those of classical GRBs (Dermer, Chiang, \\& B{\\\"o}ttcher 1999; Heise et al.\\ 2001), or {\\it iii)} fireballs with large $\\Gamma$-factors and/or low baryon loading that in the case of internal shocks lead to the emission of less energetic photons (Zhang \\& M{\\'e}sz{\\'a}ros 2002; Barraud et al.\\ 2003).  Prior to XRF~030723, four XRFs had been localized to arcmin accuracy. The first was the BeppoSAX burst XRF~011030 (e.g., Heise et al.~2001) for which a likely afterglow was detected at radio (Taylor et al.~2001) and X-ray wavelengths (Harrison et al.~2001). The first XRF localized by HETE-2 (XRF~011130) did not lead to an unambiguous afterglow detection despite extensive efforts at optical, X-ray and radio wavelengths. XRF~020427, localized by BeppoSAX, exhibited an X-ray afterglow, localized to arcsec accuracy by the Chandra X-ray Observatory (Fox 2002). The first HETE-2 XRF for which an afterglow was identified was XRF~020903 (Ricker et al.~2002). A candidate optical and radio afterglow was detected and a bright galaxy coincident with the transient was found at $z = 0.251$ (Soderberg et al.~2003). Finally, Watson et al.\\ (2004) suggest that the Integral burst GRB~031203 was actually an XRF based on the modeling of a dust-echo in our galaxy observed by XMM (Vaughan et al.\\ 2004). For GRB~031203 a radio afterglow (Frail 2003), but no optical afterglow, has been detected. The likely redshift of this burst is only 0.105 based on the spectrum of a galaxy coincident with the position of the radio afterglow (Prochaska et al.~2004).  HST observations have been obtained of the fields of XRF 011030, 020903, and 020427. In all cases a candidate host galaxy has been identified. XRFs 011030 and 020427 exhibited blue $\\mathrm{R}\\sim 24$ host galaxies, typical of GRB host galaxies (Bloom et al.~2003). The galaxies are probably not at a very high redshift ($z\\leq3.5$; Bloom et al.~2003). The host galaxy of XRF~020903 exhibits a complex morphology (Levan et al.\\ 2002; Soderberg et al.~2003). A single emission line at 8485 {\\AA} has been detected from the host galaxy of XRF~020427 (van Dokkum \\& Bloom 2003).  XRF~030723 was detected by the FREGATE, WXM, and SXC instruments on-board the HETE-2 satellite on July 23.26965 2003 UTC (HETE trigger H2777). The event was localized with the SXC to a 2 arcmin radius error circle at high galactic latitude (b = 50$^\\mathrm{o}$) in the constellation Pisces Aus \\cite{prigozhin03}. The burst duration (T$_{90}$) was 25 s. The total fluence in the 7--30~keV band was $\\sim$2$\\times$10$^{-7}$ ergs cm$^{-2}$ and in the 30--400~keV band the fluence was less than 7$\\times$10$^{-8}$ ergs cm$^{-2}$. It was hence clearly an XRF according to the HETE-II definition. Observations of the X-ray afterglow to XRF~030723 has been reported by Butler et al.\\ (2003). In the radio band only an upper limit of 180 $\\mu$Jy at 8.46 GHz (July 26.42 UT) has been reported (Soderberg, Berger, \\& Frail 2003). In this paper we present a comprehensive optical and near-infrared (near-IR) study of its afterglow - the first such study for any XRF. ",
        "conclusions": "We have presented the first detailed multi-color study of the afterglow of an XRF. The properties of the early light curve are in agreement with the predictions for off-axis classical GRBs both in the fireball and cannonball models. After the first week there is a strong bump in the light curve, peaking at around 16 days after the XRF. The SED after the peak of the bump becomes significantly redder and can no longer be fit by a power-law. We have considered the possible interpretations of this bump, e.g., an underlying supernova, a second jet, a refreshed shock or an encounter with a density enhancement. We conclude that a SN is the most likely explanation since no other model can qualitatively explain the evolution of the SED. We also present deep spectroscopy of the burst both in the afterglow and in the bump phase. Unfortunately, no redshift determination has resulted from the spectra as there are no significant absorption or emission lines detected. A firm upper limit of $z\\leq2.3$ on the redshift is established from the lack of any Ly$\\alpha$ forest lines in the afterglow spectrum. If the bump in the light curve is indeed due to a supernova the most likely redshift is in the range $\\sim$0.3--1.0. The lack of significant absorption lines in the afterglow spectrum and the lack of any indication of extinction in the afterglow SED is indicative that the progenitor was in an environment with lower density than typically found for GRBs. The host galaxy has not been unambiguously detected. From the upper limit of $\\mathrm{R}\\gtrsim27.0$ and the preferred redshift range of $z\\approx0.3$--$1.0$ we conclude that the host is a dwarf galaxy."
    },
    "0402/astro-ph0402289_arXiv.txt": {
        "abstract": "Long-term (up to 10000d) monitoring has been undertaken for 41 Seyferts in the near infrared (1.25 -- 3.45$\\mu$m). All but two showed variability, with amplitudes at $K$ in the range $<$ 0.1 to $>$1.1 mag. The timescale for detectable change is from about one week to a few years.  Where contemporary observations of variability in x-rays, UV or visible light exist, it is found that the near-infrared varies in a similar way, though in some cases the shorter-wavelength IR bands are diluted by underlying galaxy radiation.  A simple cross-correlation study indicates that there is evidence for delays of up to several hundred days between the variations seen at the shortest wavelengths ($U$ or $J$) and the longest ($L$) in many galaxies. In particular, the data for Fairall 9 now extend to twice the interval covered in earlier publications and the delay between its UV and IR outputs is seen to persist.  An analysis of the fluxes shows that, for any given galaxy, the colours of the variable component of its nucleus are usually independent of the level of activity. The state of activity of the galaxy can be parameterized.  Taken over the whole sample, the colours of the variable components fall within moderately narrow ranges. In particular, the $H-K$ colour is appropriate to a black body of temperature 1600K. The $H-K$ excess for a heavily reddened nucleus can be determined and used to find $E_{B-V}$, which can be compared to the values found from the visible region broad line ratios.  Using flux-flux diagrams, the flux within the aperture from the underlying galaxies can often be determined without the need for model surface brightness profiles. In many galaxies it is apparent that there must be an additional constant contribution from warm dust. ",
        "introduction": "This programme was started with the aim of making a reliable study of the variability of Seyfert galaxies in the near infrared ($JHKL$ bands; 1.25, 1.65, 2.2 and 3.45$\\mu$m). Previous work had indicated that some Seyfert galaxies do vary, but the data remained somewhat sparse and controversial. At the commencement of the programme it was believed, for example, that Seyfert 1 galaxies varied but Seyfert 2 galaxies did not. A 450-day programme by Penston et al (1974) appeared to confirm this viewpoint. Lebofsky \\& Rieke (1980) reported on variations observed in several Seyferts, of which that in 3C\\,120 was the most spectacular, but the numbers of observations were very limited.  Some results from this programme have already been published. One of the most unexpected was the discovery by Clavel, Wamsteker and Glass (1989) that there was a delay of ~400 days between the UV variations of Fairall 9 and the response of its infrared output. This was satisfactorily interpreted according to a dust reverberation model (Barvainis, 1987). NGC\\,1566 showed a delay of 2 $\\pm$ 1 month between the UV and the IR (Baribaud et al, 1992). Similarly, a possible delay of $\\sim85$d was seen in NGC\\,3783 (Glass, 1992). The infrared output of NGC\\,1068 was shown to have increased by a factor of two between the early 1970s and 1995 by Glass (1995). The galaxy NGC\\,2992 showed an energetic flare-like outburst with a decay time of about 900d (Glass, 1997a) and NGC\\,7469 showed a dramatic but relatively short-lived decline in its nuclear energy output in October 1989 (Glass 1998).  More recently, Salvati et al (1993) observed an outburst in NGC\\,4051 and Nelson (1996) reported one in Mkn\\,744, with a 32 $\\pm$ 7 day delay between the visible and IR light curves. Oknyanskij \\& Horne (2001) have re-examined the published data and concluded that the observed delays are consistent with dust sublimation radii $R \\propto \\sqrt{L_{\\rm UV}}$. Oknyanskij (2002) includes two additional objects. The angular diameter of the dust ring or shell as viewed fron the earth is of order tenths of milliarcsec.  Throughout this paper it is assumed that the excess infrared radiation from a Seyfert when compared to a less active galaxy comes from dust that is heated by ultraviolet radiation originating near a central massive black hole. It is usually assumed that this dust is distributed in a torus; e.g., Pier \\& Krolik (1992). The maximum temperature that dust can reach without sublimating, in the neighbourhood of 1600K, together with the ultraviolet luminosity of the nucleus, sets the inner radius of its distribution and the corresponding light travel time gives the delay in the response of a dust grain to a change in the ultraviolet flux from the nucleus.  During part of the present programme, $UBVRI$ photometry was obtained by Winkler et al (1992), Winkler (1997) and other observers. ",
        "conclusions": "Thirty-nine of the forty-one Seyfert galaxies in this sample have shown variability during the observing programme; some of them by over one mag. These variations reflect changes in the ultraviolet output of the central engine in the same way as the visible-light continuum and spectral line observations do; i.e., after a certain time-lag.  Lags are common and most conspicuous at $L$, but difficult to determine accurately. Several of the lower-luminosity galaxies in the present sample have been shown to vary and to exhibit probable lags on time scales of tens of days. These will repay study with more frequent sampling. There is some suggestion that a decline in the $L$ flux following a drop in the UV from the central engine may show a longer time constant than its counterpart at $K$, perhaps because of increased radial extent of the dust zone which contributes most at this wavelength.  For the most part, there are strong linear correlations between the $J,H,K$ and $L$ fluxes. The colours of the variable components are therefore easy to determine, provided that the variations have sufficient amplitude and the photometry is sufficiently accurate. There is a high degree of uniformity of colours among the variable components which allows an estimate of the nuclear reddening to be made in heavily reddened cases. It also clear that the near-infrared emission mechanism is the same in all Seyfert galaxies.  In many cases, the fluxes seen through the measuring aperture can be separated into variable and underlying galaxy components without recourse to surface-brightness modelling. In a number of galaxies, a further non-variable component, probably arising from cool dust, is clearly present at $L$ and may be modelled.  \\begin{table} \\caption{Sample data table, for III Zw 2. See web version for remainder.} \\begin{tabular}{llllll} MJD  &  $J$  &  $H$  &  $K$  &  $L$  & $L_{\\rm err}$ \\\\ 3717 & 13.18 & 12.25 & 11.11 &  9.82 & .12 \\\\ 6695 & 13.61 & 12.55 & 11.41 &  9.86 & .14 \\\\ 6989 & 13.71 & 12.67 & 11.54 & 10.01 & .09 \\\\ 7436 & 13.60 & 12.70 & 11.63 &  9.95 & .11 \\\\ 7437 & 13.55 & 12.72 & 11.61 & 10.28 & .10 \\\\ 7342 & 13.55 & 12.68 & 11.65 &       &     \\\\ 7842 & 13.54 & 12.74 & 11.67 & 10.32 & .16 \\\\ 8132 & 13.67 & 12.78 & 11.68 & 10.19 & .05 \\\\ 8172 & 13.59 & 12.69 & 11.62 &  9.91 & .06 \\\\ 8227 & 13.56 & 12.61 & 11.57 & 10.00 & .13 \\\\ 8529 & 13.79 & 12.86 & 11.67 & 10.43 & .22 \\\\ 8574 & 13.64 & 12.76 & 11.60 & 10.07 & .17 \\\\ 8576 & 13.75 & 12.75 & 11.67 & 10.10 & .18 \\\\ 8874 & 13.75 & 12.84 & 11.67 & 10.13 & .05 \\\\ 9966 & 13.84 & 12.89 & 11.79 & 10.41 & .12 \\\\ 10655& 13.61 & 12.66 & 11.56 & 10.31 & .17 \\\\ 10742& 13.64 & 12.61 & 11.69 & 10.00 & .16 \\\\ 11117& 13.60 & 12.71 & 11.60 & 10.43 & .16 \\\\ \\end{tabular} \\end{table}"
    },
    "0402/astro-ph0402076_arXiv.txt": {
        "abstract": "Preliminary results are reported on the spectral and timing properties of the spectacular 2001 superburst of 4U\\,1636$-$536 as seen by the RXTE/PCA. The (broad-band) power-spectral and hardness properties during the superburst are compared to those just before and after the superburst. Not all of the superburst emission can be fitted by pure black-body radiation. We also gathered BeppoSAX/WFC and RXTE/ASM data, as well as other RXTE/PCA data, obtained days to months before and after the superburst to investigate the normal X-ray burst behavior around the time of the superburst. The first normal X-ray burst after the 2001 superburst was detected $\\simeq$23~days later. During inspection of all the RXTE/ASM data we found a third superburst. This superburst took place on June 26, 1999, which is $\\simeq$2.9~yrs after the 1996 superburst and $\\simeq$1.75~yrs before the 2001 superburst. The above findings are the strongest constraints observed so far on the duration of the cessation of normal X-ray bursts after a superburst and the superburst recurrence times. ",
        "introduction": "Superbursts are a recently discovered new mode of unstable thermo-nuclear burning on the surface of a neutron star. These neutron stars reside in so-called low-mass X-ray binaries (LMXBs) in which it continuously receives fresh material from a donor star. A large fraction of these LMXBs are known to exhibit short (seconds to minutes) so-called Type~I X-ray bursts which are due to unstable thermo-nuclear burning of H and/or He (for reviews, see Lewin et al.\\ \\cite{Letal1993}; Strohmayer \\&\\ Bildsten \\cite{SB2003}). These `normal' X-ray bursts appear as rapid ($\\sim$1--10\\,s) increases in the X-ray flux, followed by an exponential-like decline. They recur with a frequency (typically hours to days) which is (partly) set by the supply rate of the fresh fuel. The (net) burst X-ray spectra are well described by black-body emission from a compact object with $\\sim$10\\,km radius and color temperature of $\\sim$1--2\\,keV. The color temperature increases during the burst rise and decreases during the decay, reflecting the heating and subsequent cooling of the neutron star surface. Typical integrated burst energies are in the 10$^{39}$ to 10$^{40}$\\,erg range.  The durations and output energies of superbursts are roughly a factor of a thousand more than that of the normal X-ray bursts. They display, however, the usual relatively fast rise and exponential decay, as well as the hardening during the rise and subsequent softening during the decay, as seen in normal X-ray bursts. Also, the (net) superburst X-ray spectra are described in the same way as the normal X-ray bursts. For an overview of the superburst properties we refer to Kuulkers \\cite{K2003}, and references therein.  The current view is that the superbursts are due to the energy release of unstable burning of C, which is the left over from the stable and unstable burning of H and He (Cumming \\&\\ Bildsten \\cite{CB2001}; Strohmayer \\&\\ Brown \\cite{SB2002}). Additional energy may be released through photo-desintegration triggered by the superburst (Schatz et al.\\ \\cite{Setal2003}).  The first superburst was discovered with the BeppoSAX/WFCs to come from the X-ray burster 4U\\,1735$-$444 (Cornelisse et al.\\ \\cite{Cetal2000}). Another superburst was independently found with the RXTE/PCA, originating from the X-ray burster 4U\\,1820$-$303 (Strohmayer \\&\\ Brown \\cite{SB2002}; see also Ballantyne \\&\\ Strohmayer \\cite{BS2004}). Thereafter, six more events have been seen to occur in five other X-ray bursters with the BeppoSAX/WFCs (Cornelisse et al.\\ \\cite{Cetal2002}; Kuulkers et al.\\ \\cite{Ketal2002}, see also in 't Zand et al.\\ \\cite{Zetal2004}), RXTE/PCA (Strohmayer \\&\\ Markwardt \\cite{SM2002}) and RXTE/ASM (Wijnands \\cite{W2001}; Kuulkers \\cite{K2002}). The fact that only eight such events have been found so far indicates that they must be rare. Two superbursts from 4U\\,1636$-$536 were seen to occur $\\simeq$4.7~yrs apart (Wijnands \\cite{W2001}); observational estimates of the recurrence times are on the order of a year (see, e.g., Kuulkers \\cite{K2002}; in 't Zand et al.\\ \\cite{Zetal2003}).  4U\\,1636$-$536 is a regular X-ray burster and a so-called `atoll' source (e.g., van der Klis et al.\\ 1990). We here present preliminary results focussed on the X-ray behavior of the source during and around the superburst which occurred on 2001 Feb 22.  \\subsection{The 2001 superburst as seen by the RXTE/PCA}  A first account of the RXTE/PCA data during the 2001 superburst from 4U\\,1636$-$536 is given by Strohmayer \\&\\ Markwardt \\cite{SM2002}. Note that this is event was also seen by the RXTE/ASM (Wijnands \\cite{W2001}; see also below). Strohmayer \\&\\ Markwardt  \\cite{SM2002} found a highly coherent 582\\,Hz pulsation during a small part ($\\sim$800\\,s) of the superburst.  \\begin{figure} \\includegraphics[bbllx=33pt,bblly=70pt,bburx=568pt,bbury=550pt,angle=-90,width=0.48\\textwidth]{kuulkerse_f1.ps} \\caption{RXTE/PCA 2--16\\,keV light curve ({\\it Top}) and hardness (ratio of the count rates in the 6.5--16\\,keV and 2--6.5\\,keV energy bands) curve ({\\it Bottom}) of the 2001 superburst of 4U\\,1636$-$536. Note that slew data are not included, and that the count rates are not corrected for background and dead time. Standard 2 data are used from PCU's 0, 2 and 3; the time resolution is 16\\,s.} \\end{figure}  The light curve (see also Strohmayer \\&\\ Markwardt \\cite{SM2002}), as well as the hardness curve, of the superburst are shown in Fig.~1. Clearly, they resemble the curves usually seen for normal X-ray bursts, except for the duration of a few hours. A high time resolution (0.125\\,s) light curve of the superburst is shown by Kuulkers \\cite{K2003}, including data from the slew just before the third orbit of data on 4U\\,1536$-$536. This light curve showed an increase in count rate of a few 1000\\,c\\,s$^{-1}$ about 125\\,s before the precursor X-ray burst. It was pointed out to us by Tod Strohmayer that this increase is actually due to the turn-on of PCU 3; this was not taken into account when doing the collimator response correction. The superburst, therefore, did {\\it not} start $\\sim$125\\,s before the precursor X-ray burst, but between $\\sim$145\\,s and $\\sim$46\\,min before the precursor.  \\begin{figure} \\includegraphics[bbllx=78pt,bblly=23pt,bburx=587pt,bbury=702pt,angle=-90,width=0.48\\textwidth]{kuulkerse_f2.ps} \\caption{Color-color diagram of 4U\\,1636$-$536 using Standard 2 data between Feb 28, 1996 and April 10, 2002. The data on Feb 22, 2002 are indicated in red. Soft and hard color are defined as the ratio of the count rates in the 3.5--6.4\\,keV and 2.0--3.5\\,keV bands, and, respectively, the 9.7--16\\,keV and 6.4--9.7\\,keV bands. The count rates are normalized to the Crab (see Di Salvo et al.\\ \\cite{DSetal2002}; van Straaten et al.\\ \\cite{Setal2003}) and not corrected for dead time.} \\end{figure}  A color-color diagram (CD) of all public data on 4U\\,1636$-$536 between Feb 28, 1996 and April 10, 2002 is shown in Fig.~2 (this is an extension of the data set used by Di Salvo et al.\\ \\cite{DSetal2002}). The superburst (shown in red) started at the bottom part of the CD, i.e., in the so-called `banana' branch. During the rise to superburst peak the source hardened, while during the decay it softened, with the hard colors even dropping below pre-superburst values. At the end of the observation the source resumed its pre-superburst colors again.  We investigated the broad-band timing behavior using the high-time resolution modes available during parts of the observation (see Strohmayer \\&\\ Markwardt \\cite{SM2002}). No strong features are seen before and after the superburst, except for some very-low frequency noise. This is as expected from previous analyses of the timing behavior at the corresponding portion of the CD (see Di Salvo et al.\\ \\cite{DSetal2002}). The high-timing data during the superburst, however, show a hint of a quasi-periodic oscillation near 1150\\,Hz. A more careful analysis is in progress.  Using the Standard 2 mode data we extracted X-ray spectra at 16\\,s time resolution throughout the 2001 superburst. As is common practice for normal X-ray bursts, we subtracted the average pre-superburst X-ray spectrum from the superburst spectra. The residual spectra were modeled by pure black-body emission (keeping the interstellar absorption, $N_{\\rm H}$, fixed at 3$\\times$10$^{21}$~atoms\\,cm$^{-2}$, see e.g., Christian \\&\\ Swank \\cite{CS1997}). As the superburst flux increased the color temperature increased from $\\simeq$1\\,keV up to 2.35\\,keV, while during the decay the color temperature decreased again. The inferred black-body radius was more or less constant ($\\simeq$6\\,km at 6\\,kpc). The fits are, however, not perfect throughout the whole burst, with the strongest deviations seen when the superburst flux had dropped to about half its peak value. The deviations are very similar to those seen during the superburst from 4U\\,1820$-$303, and may show the influence of the superburst emission on the accretion disk surrounding the neutron star (Strohmayer \\&\\ Brown \\cite{SB2002}; Ballantyne \\&\\ Strohmayer \\cite{BS2004}). The total integrated superburst flux is about 0.65$\\times$10$^{42}$\\,erg (at 6\\,kpc).  \\subsection{BeppoSAX/WFC and RXTE/ASM measurements around the 2001 superburst}  \\begin{figure} \\includegraphics[bbllx=52pt,bblly=50pt,bburx=553pt,bbury=552pt,angle=-90,width=0.48\\textwidth]{kuulkerse_f3.ps} \\caption{{\\it Top}: RXTE/ASM (1.5--12\\,keV) light curve of 4U\\,1636$-$536 from Jan 6, 1996 to Sep 4, 2003. The data shown are 3-day averages and normalized to the Crab count rate. With dotted lines we indicate the times of the 3 superbursts seen from 4U\\,1636$-$536 (see text). {\\it Bottom}: The BeppoSAX/WFC (2--28\\,keV) light curve at 15-min time resolution during the interval Jan 21 to Apr 9, 2001 is shown in black. Data which were less than 0.5~day apart were connected to guide the eye. They were normalized to the Crab count rate. In gray we show the RXTE/ASM 1-day average data points, as well as in red the RXTE/PCA 15-min average light curves, in the same period; within an RXTE/PCA observation the data are connected. In the top part of the bottom panel we show with vertical bars the times of occurrence of normal X-ray bursts.} \\end{figure}  In Fig.~3 (bottom) we show in black the WFC light curves acquired around the 2001 superburst of 4U\\,1636$-$536, as well as in gray the RXTE/ASM 1-day average and in red the RXTE/PCA light curves. The BeppoSAX/WFC observations did not cover the superburst; the RXTE/ASM, however, did cover it (Wijnands \\cite{W2001}; see also below). In the top of the bottom panel figure we annotate with bars the times of normal X-ray bursts as seen by the BeppoSAX/WFC (black) and the RXTE/PCA (red). Clearly, the source showed normal X-ray bursts before the superburst, with an average rate of about 6 per day. Observations after the 2001 superburst are sparse; the first BeppoSAX/WFC and RXTE/PCA observations were $\\sim$23~days, respectively, $\\sim$42~days after the superburst. During both observations a normal X-ray bursts was seen. We note that the peak count rates during the normal X-ray bursts are a factor of $\\sim$1.5--2 higher than that reached during the precursor burst during the superburst. This is consistent with an early ignition of a H/He layer by the superburst as it starts (Cumming \\cite{C2003}).  \\subsection{RXTE/ASM superbursts} \\label{asm}  In the top panel of Fig.~3 we show the long-term RXTE/ASM 3-day average light curve of 4U\\,1636$-$536. The source varies on time scales of days to months to years. It may be interesting to note that this light curve resembles that of KS\\,1731$-$260 during the first 4 year of the RXTE mission. Apart from the 2001 event, a superburst was also seen by the RXTE/ASM on June 19, 1996 (Wijnands \\cite{W2001}). In Fig.~4 we show the latter event on the left, whereas the 2001 event is shown on the right. During the inspection of the individual RXTE/ASM dwell light curves (see, e.g., Wijnands \\cite{W2001}), we found, however, another candidate superburst. This candidate, which occurred on June 26, 1999, is shown in the middle panels of Fig.~4. Although the peak of the light curve (top) consists of 1 data point only, the data points after that are on average higher than data before and after. To guide the eye, we have connected the data points; the light curve of the candidate is consistent with a fast rise plus slower exponential decay light curve. In the bottom panel we show the hardness curve in the same period; the hardness values during the peak and during the decay, as well as their evolution, are very similar to those of the superbursts in 1996 and 2001. Based on the similarities we regard the 1999 event as a real superburst; if true, then the recurrence times between the 1996 and 1999 events and the 1999 and 2001 are $\\simeq$2.9~yrs and $\\simeq$1.75~yrs, respectively.  \\begin{figure} \\includegraphics[bbllx=33pt,bblly=70pt,bburx=573pt,bbury=770pt,angle=-90,width=0.48\\textwidth]{kuulkerse_f4.ps} \\caption{RXTE/ASM 1.5--12\\,keV light curves ({\\it top}) and hardness curves ({\\it bottom}) around the superbursts seen from 4U\\,1636$-$536 in 1996 (Wijnands \\cite{W2001}; {\\it left}), 1999 (see text; {\\it middle}) and 2001 (Wijnands \\cite{W2001}; Strohmayer \\&\\ Markwardt \\cite{SM2002}; {\\it right}). For the light curves we show the individual 90-s average RXTE/ASM dwell data points. Hardness is defined as the ratio of the count rates in the 5--12\\,keV and 1.5--5\\,keV bands; values are averages of nearby dwells. Data points are connected in light gray to guide the eye.} \\end{figure} ",
        "conclusions": "In this proceedings paper we have shown some preliminary results on the analysis of RXTE/PCA, RXTE/ASM and BeppoSAX/WFC data obtained during and around the 2001 superburst seen from the X-ray burster 4U\\,1636$-$536. We found that during the first observation after this superburst, i.e., 23 days later, a normal X-ray burst occurred. Previously, we found that normal X-ray bursts seem to cease to occur for up to a month after the superbursts of Ser\\,X-1 (Cornelisse et al.\\ \\cite{Cetal2002}) and KS\\,1731$-$260 (Kuulkers et al.\\ \\cite{Ketal2002}). Our observed upper limit on the duration of the cessation of normal X-ray bursts after the 2001 superburst puts strong constraints on the modeling of the thickness of the layer where the superburst originates as well as where the energy is deposited (Cumming \\&\\ Macbeth \\cite{CM2004}).  We also investigated the entire RXTE/ASM data set on 4U\\,1636$-$536 and found evidence for a third superburst, which occurred in 1999 (the first superburst occurred in 1996). The shortest observed superburst recurrence time is now $\\simeq$1.75~yrs, which is more or less in agreement with previous estimates (roughly one superburst per source per year; Kuulkers \\cite{K2002}; in 't Zand et al.\\ \\cite{Zetal2003}). Superburst recurrence times may constrain the original composition of the material which was deposited on the neutron star (Cumming \\cite{C2003}; see also Cumming \\cite{C2003b}, and the discussion in Kuulkers \\cite{K2003}).  \\begin{theacknowledgments} We thank Tod Strohmayer for pointing out to us about the turn-on of the various PCU's during the slew towards the source when the superburst occurred. \\end{theacknowledgments}"
    },
    "0402/gr-qc0402031_arXiv.txt": {
        "abstract": "The recent type Ia supernova data suggest that the Universe is accelerating now and decelerated in recent past. This may provide the evidence that the standard Friedmann equation needs to be modified. We analyze in detail a new model in the context of modified Friedmann equation using the supernova data published by the High-$z$ Supernova Search Team and the Supernova Cosmology Project. The new model explains recent acceleration and past deceleration. Furthermore, the acceleration of the expansion of the Universe is almost zero in the future. ",
        "introduction": "The type Ia supernova (SN Ia) observations support that the expansion of the Universe is speeding up rather than slowing down \\cite{sp97,tonry,riess04}. A spatially flat universe is favored by the measurements of the anisotropy of the cosmic microwave background (CMB) \\cite{pdb00,bennett03}. The observation of type Ia supernova SN 1997ff at $z\\sim 1.7$ also provides the evidence that the Universe is in the acceleration phase and was in the deceleration phase in recent past \\cite{agr,mstagr}. The transition from the deceleration phase to the acceleration phase happens around the redshift $z_{\\rm T}\\sim 0.4$ \\cite{mstagr,daly}. The simplest way of explaining the acceleration phenomenon is through the cosmological constant with the equation of state parameter $\\omega_{\\rm \\Lambda}=-1$. One easily generalizes the cosmological constant model to dynamical cosmological constant models, such as the dark energy model with negative constant equation of state parameter $-1\\le \\omega_{\\rm Q}<-1/3$ and the holographic dark energy models \\cite{holo,kaloper}. In general, a scalar field $Q$ called quintessence field that slowly evolves down its potential $V(Q)$ takes the role of a dynamical cosmological constant \\cite{quint,johri,aasss,uavs}. The energy density of the quintessence field must remain very small compared with radiation and matter at early epoches and evolves in a way that it started to dominate the Universe around the redshift $0.4$. If we remove the null energy condition restriction $\\omega_{\\rm Q}\\ge -1$ to allow supernegative $\\omega_{\\rm Q}<-1$, then the phantom energy models arise \\cite{phantom}. More exotic equation of state is also possible, such as the Chaplygin gas model with the equation of state $p=-A/\\rho$ and the generalized Chaplygin gas models with the equation of state $p=-A/\\rho^\\alpha$ \\cite{chaply}. There are fields other than the quintessence field, such as tachyon field, as dark energy \\cite{tachyon}.  Although there are many dark energy models proposed in the literature, the nature of dark energy is still unknown. Therefore it is also possible that the observations show a sign of the breakdown of the standard cosmology. In brane cosmology, our observable universe is a 3-brane which is located at the $Z_2$ symmetry fixed plane of a five dimensional Ads (Anti-de Sitter) spacetime with the extra dimension to be an orbifold $S_1/Z_2$. Bin\\'etruy, Deffayet and Langlois showed that the Friedmann equation takes the unconventional form $H^2\\sim \\rho +\\rho^2$ in brane cosmology \\cite{brane}. In addition, Chung and Freese argued that almost any relationship between $H$ and $\\rho$ was possible with the help of extra dimension \\cite{chung}. Along this line of reasoning, the usual Friedmann equation $H^2=8\\pi G\\rho/3$ is modified to a general form $H^2=g(\\rho)$ \\cite{freese02,dvali,sen03,zhu03,wang,gong03a,gong03}. In these models, no exotic matter form is needed and the Universe is composed of the ordinary matter and radiation only. The function $g(x)$ satisfies the following conditions: (1) $g(x_0)=1$; (2) $g(x)\\approx x$ when $z\\gg 1$; (3) $g(x_0)>3x_0g'(x_0)/2$, where $x=\\Omega_{\\rm m0}(1+z)^3+\\Omega_{\\rm r0}(1+z)^4$, $x_0=\\Omega_{\\rm m0}+\\Omega_{\\rm r0}\\approx \\Omega_{\\rm m0}$ and $\\Omega_{\\rm r0}=8.40\\times 10^{-5}$ \\cite{bennett03}. In Cardassian model, $g(x)=x+Bx^n$. In generalized Cardassian model, $g(x)=x[1+Bx^{\\alpha(n-1)}]^{1/\\alpha}$ \\cite{freese02}. For Chaplygin gas like model, $g(x)=x+(A'+B'x^\\beta)^{1/\\beta}$ \\cite{gong03a}. In Ref. \\refcite{dvali}, Friedmann equation is modified to be $H^2+H^\\alpha\\sim \\rho$ based on the idea of embedding a three brane into a five dimensional flat spacetime. The modified model can be mapped into dark energy model. For example, $g(x)=x$ for standard cosmology; $g(x)=x+\\Lambda$ for $\\Lambda$CDM model; $g(x)=x+Bx^{1+\\omega_{\\rm Q}}$ for dark energy models with constant equation of state parameter $\\omega_{\\rm Q}$. In this paper, we discuss the new model proposed in Ref. \\refcite{gong03} in detail. We first derive the model with the help of extra dimension, then we use different SN Ia data sets to fit the model. ",
        "conclusions": ""
    },
    "0402/astro-ph0402649_arXiv.txt": {
        "abstract": "A survey of 329 nearby galaxies (redshift z $<$ 0.045) has been conducted to study the distribution of mass and light within spiral galaxies over a range of environments.  The 18 observed clusters and groups span a range of richness, density, and X-ray temperature, and are supplemented by a set of 30 isolated field galaxies.  Optical spectroscopy taken with the 200-inch Hale Telescope provides separately resolved \\Halpha and \\fnii major axis rotation curves for the complete set of galaxies, which are analyzed to yield velocity widths and profile shapes, extents and gradients.  \\HI line profiles provide an independent velocity width measurement and a measure of \\HI gas mass and distribution.  \\Ib images are used to deconvolve profiles into disk and bulge components, to determine global luminosities and ellipticities, and to check morphological classification.  These data are combined to form a unified data set ideal for the study of the effects of environment upon galaxy evolution. ",
        "introduction": "\\label{ch:intro}  This study concerns the effects of the cluster environment on the distribution of mass and light in spiral galaxies.  These data are part of a large survey of nearby clusters (redshift $z \\le 0.045$) designed to study the dependence of mass--to--light distributions upon cluster densities by comparing core galaxies to periphery and low density region members, to evaluate the global properties of rotation curves, and to examine fundamental galaxy properties in varied environments.  There are several other large samples of optical rotation curves in the literature at present, including Courteau's 1992 sample of 350 which focuses primarily upon field galaxies, and Mathewson \\& Ford's 1996 sample of 2447 in the southern hemisphere.  This data set contains the rotation curves of 329 northern hemisphere cluster and field galaxies, with \\HI observations for most to provide additional, independent constraints upon mass distributions.  The velocity profiles of spiral galaxies were originally believed to fall in a Keplerian decline at large radii, following an inverse $r$ law near the edge of the optical disk where the light (and presumably mass) dropped off.  When the first measured rotation curves did not show such a trend they were assumed to probe only the inner region, rather than that predicted to exist at large radii for a pure exponential disk model.  As hardware improvements extended the sensitivity and depth of observations, however, more and more of the outer regions of galaxy disks were observed.  It became clear that both optical and 21~cm line observations showed velocity profiles which rose in the inner regions but often remained flat at their outer extents, implying the presence of a considerable amount of dark matter within the optical radius (\\cf Bosma 1978, review by Faber \\& Gallagher 1979).  In the early 1980's Rubin and collaborators began a pioneering study of spiral galaxies to attempt to constrain mass distributions by analyzing optical velocity profiles and light distributions (Rubin, Burstein, Ford, \\& Thonnard 1985, and references therein).  They deduced that velocity profiles correlated well with basic galaxy properties as luminosity (brighter galaxies having larger velocity widths which rose more slowly) and morphological type (earlier type galaxies having larger velocity widths). This initial sample of field galaxies was later extended to include a set of cluster spirals (Burstein, Rubin, Ford, \\& Whitmore 1986; Rubin, Ford, \\& Whitmore 1988; Whitmore, Forbes, \\& Rubin 1988), to evaluate the importance of the cluster environment upon the mass distributions.  They found a population of cluster galaxies with rotation curves which appeared to decline at large radii, and argued that these galaxies had deficient halos which had either been stripped through an interaction with the cluster medium (\\eg ram pressure sweeping) or with close neighbors (tidal stripping), or had never been allowed to form.  Later studies (Guhathakurta \\etal 1988; Forbes \\& Whitmore 1989; Distefano \\etal 1990; Amram \\etal 1993; Adami \\etal 1999) found conflicting results regarding these trends, and thus part of the motivation for this work has been to re-examine these claims within a larger sample of galaxies.  The strong effect of the cluster environment upon the \\HI gas envelope of spiral galaxies has been well established (\\cf Haynes \\& Giovanelli 1986, Magri \\etal 1988, Haynes 1989, Solanes \\etal 2001).  Galaxies interacting with the hot intracluster medium of moderate to high X-ray luminosity clusters are observed to be strongly \\HI deficient and to have lost \\HI gas preferentially at outer radii (Haynes 1989), and the pattern of \\HI deficiency correlates strongly with clustercentric radius.  The stripped galaxies appear to have an undisturbed spiral morphology, though biased towards early types, implying that the mechanisms in play are either too weak to directly affect the stellar material or that the effect can only be recognized on a longer timescale (\\eg decreased young star formation from a depleted \\HI gas reservoir).  This paper is the first of three in a series.  The present paper (\\pone) details the observations which were made in the optical and in the radio, and defines the data analysis procedures and the modeling technique used to determine mass-to-light ratios\\ for the galaxies within the sample.  The second paper (Vogt \\etal 2004a; \\ptwo) investigates the evidence for galaxies which are currently infalling into the cores of clusters on a first pass, or show evidence of a previous passage through the core, and the role of gas stripping mechanisms in a morphological transformation of the field spiral population into cluster S0s.  The third paper (Vogt \\etal 2004b; \\pthree) explores differences in fundamental galaxy properties (size, mass, and luminosity) as a function of environment, whether spiral disks within rich cluster cores may have coalesced from their halos at an early epoch, and examines the relationship between \\HI gas stripping by the hot intracluster medium and the consequential suppression of young star formation across the disks. ",
        "conclusions": "This study concerns the effects of the cluster environment upon the distribution of mass and light in spiral galaxies.  We have observed a set of 329 nearby (redshift $z < 0.045$) galaxies, (1) to study the dependence of mass--to--light distributions upon cluster environment, (2) to evaluate the global properties of optical rotation curves (spatially extended velocity profiles), and (3) to determine the robustness of the Tully--Fisher relation within differing environments.  The 16 Abell clusters cover a wide range of cluster richness, density, and luminosity, and are supplemented by the less rich groups Cancer and \\NGC 507, and 30 isolated field galaxies.  We have examined the distribution of mass and light within the target galaxies through a study of optical emission lines, the total \\HI gas, and optical photometry. Optical spectroscopy were taken with the 200-inch Hale Telescope and provide separately resolved \\Halpha and \\fnii major axis rotation curves for the complete set of galaxies, which are analyzed to yield velocity widths and velocity profile shapes and gradients.  \\HI line profiles provide an independent velocity width measurement and a measure of \\HI gas deficiency. \\Ib images are used to deconvolve profiles into disk and bulge components and to determine global luminosities and inclination.  This paper is the first in a series of three linked papers concerning this data set.  The present paper (\\pone) covers the observational data collected; the remaining two (Vogt \\etal 2003; \\ptwo\\ and \\pthree) contain the results of the analysis of cluster properties and environmental effects."
    },
    "0402/astro-ph0402155_arXiv.txt": {
        "abstract": "\\label{abs} This work analyses the  properties of groups of galaxies in the surroundings of clusters. On the basis of a very large public Virgo Consortium  Simulation, we identified systems of galaxies in a wide range of masses. Systems with masses greater than $ M_{cut}= 4 \\times 10^{14}~M_{\\odot} ~h^{-1} $ are considered \"host\", whereas smaller systems are taken as groups. Our results show that groups properties are affected by the proximity of massive hosts. Physical properties such as velocity dispersion, internal energy ($E$) and virial radius, show an increment, whereas the mean density decreases as the host-group distance is smaller. By analysing groups with different properties, we find that the low mass and the weakly bounded ($E > 0$) subsamples, are strongly affected by the presence of the host; on the other hand, massive groups and groups with $E < 0$ do not show dependence on the host-group distance. Using a sample of groups identified in the final version of the 2dF Galaxy Redshift Survey, we find a very similar velocity dispersion behaviour in the observational data compared to results in the simulation. We also study the dependence of the groups velocity dispersion on the host masses in both, observations and simulation; finding that the larger is the host mass the higher is the effect on its vicinity. ",
        "introduction": "\\label{int} In hierarchical models for galaxy formation, structures grow due to the accretion of smaller systems which formed at earlier times. According to this scheme, these systems should not only be in the cluster itself, but also in its vicinity, harbouring an important number of groups or small clusters of galaxies. In this sense, clusters of galaxies represent high density enhancements in the distribution of galaxies.  On the basis of this scenario, it is natural to ask whether physical properties of structures depend on its proximity to high density regions, which, as previously mentioned, may be well represented by rich clusters of galaxies; or, if on the contrary, they evolve independently on the presence of denser regions.  Several works have focussed on the evolution of substructures inside massive systems in simulations or rich galaxy clusters (Tormen et al. 1998, Ghigna et al. 2000, Taylor \\& Babul 2003, De Lucia et al. 2003); nevertheless, not much attention has been paid to structures in the outer regions of these systems. Hence, it is not well known if some of the behaviours of substructures inside the clusters, could also be applicable to groups in their periphery.  Einasto et al. (2003) have investigated the properties of loose groups in the Las Campanas Redshift Survey (LCLGs), in the vicinity of rich clusters of galaxies, including clusters from the Abell and APM catalogues, X-ray clusters and also a sample of the richest groups from the LCLG catalogue itself. By using an additional sample, consisting of all those LCLGs which do not neighbour a rich cluster, they have been able to compare the properties obtained for dense-environment LCLGs with those of typical LCLGs. They found that in most cases, the observed richness of groups near rich clusters is larger than the corresponding value for groups in the comparison sample. The same effect is shown when using Abell counts as a measure of groups richness. The harmonic radius and the velocity dispersion are also somewhat larger in the neighbourhood of rich clusters than in typical loose groups. They also found a strong mass segregation; indicating that loose groups in the vicinity of clusters have masses that are larger than mean masses of groups drawn from the comparison sample. The same trend is observed in the luminosities of LCLGs, showing that groups around rich clusters seem to be more luminous. At last, they found that these properties of loose groups do not show any dependence on clusters richness.  Lemson \\& Kauffmann (1999), also address the issue of properties of dark matter haloes as a function of local density; concluding that mass is the only halo property which correlates significantly with local environment. In fact, high mass haloes are underepresented in low density regions, whereas in high density regions the situation is the other way round.  In this work we analyse the properties of groups in the vicinity of rich clusters; taking into account their masses, internal energies, velocity dispersions and virial radii. To this end, we use data extracted from a high resolution numerical simulation and from the final version of the 2dF Galaxy Redshift Survey (2dFGRS) (Colless et al. 2003). The outline of this paper is as follows. In section \\ref{simulacion} we describe the simulation that we used and give a brief description of the group identification algorithm; in section \\ref{2dF} we give a description of the 2dFGRS and the identification of groups in this catalogue; section \\ref{resultados} contains the results obtained for both, the simulation (subsection \\ref{resulsim}) and the observational data (subsection \\ref{resul2dF}); a summary and a discussion of our results are presented in section \\ref{conclusiones}. ",
        "conclusions": "\\label{conclusiones}  Given the large volume and high resolution of the  publicly available VLS, we could identify a large number of systems of galaxies in a wide range of masses. Using a mass threshold we defined several samples of host systems and groups.  Our results show that the median velocity dispersion of groups of galaxies increases as the distance to the host clusters is smaller. Taking subsamples of groups, considering high and low masses, we find that the  previous behaviour is not present in the high mass subsample whereas low mass groups account for the observed effect. This low mass subsample was used to study the behaviour of groups according to their internal energies. We found that the internal energy of groups increases towards smaller values of host-group distances, showing up that the gravitational equilibrium of groups is influenced  by the presence of massive systems. By splitting this sample in two groups subsamples with positive and negative energies, we found that only those with positive energies increase the median velocity dispersion towards smaller host distances. In order to unmask the source of these effects, we studied the dependence on the distance of complementary physical properties. By considering smaller host-groups distances, the virial radius increases and the opposite trend is observed for the mean density.  A possible scenario describing the above results, may be placed within the framework of the tidal forces exerted by the  massive clusters over neighbouring systems. In the proximity of a cluster, its potential has a strength which is similar to the one exerted by the group itself; indicating that tidal forces are likely to affect groups properties.   Finally, analysing different  subsamples of host clusters we found that, the more massive is the host cluster the greater is the effect on the group velocity dispersion, which is also consistent with the previous model.  Using the final version of the 2dF redshift survey, we repeated the velocity dispersion analysis performed to the VLS groups. In this case we take a subsample of VLS groups, which reproduces the velocity dispersion distribution of groups identified in the 2dFGRS. The trend we find for the observational data is in excellent agreement with the result we obtain for the simulation subsample. Moreover, 2dF groups velocity dispersion show a dependence on the host masses very similar to that found in the simulation. The good agreement  between the results coming from the numerical simulation and the observational data suggests that no astrophysical mechanisms, other than the  gravitational forces, are needed in order to explain the dynamical properties of groups of galaxies."
    },
    "0402/astro-ph0402225_arXiv.txt": {
        "abstract": "We present \\sauron\\ integral-field observations of the S0 galaxy  NGC\\,7332. Existing broad-band ground-based and HST photometry reveals a double disk structure and a boxy bulge interpreted as a bar  viewed close to edge-on. The \\sauron\\ two-dimensional stellar kinematic maps  confirm the existence of the bar and inner disk but also uncover the  presence of a cold counter-rotating stellar component within the central  $250$~pc. The \\hbeta\\ and \\oiii\\ emission line maps show that the  ionised gas has a complex morphology and kinematics, including  both a component counter-rotating with respect to the stars and a fainter co-rotating one. Analysis of the absorption line-strength maps show that NGC\\,7332 is young everywhere. The presence of a large-scale bar can explain most of those properties, but the fact that we see a  significant amount of unsettled gas, together with a few peculiar  features in the maps, suggest that NGC\\,7332 is still evolving. Interactions  as well as bar-driven processes must thus have played an important  role in the formation and evolution of NGC\\,7332, and presumably of S0  galaxies in general. ",
        "introduction": "\\label{Sec:Introduction}  The nature of S0 galaxies has been a matter of controversy since their  appearance in the first morphological classification schemes \\citep{hubble36,devac59,vdbergh60a,vdbergh60b} and many authors refer to  them as normal spirals stripped of their gas  \\citep[e.g.,][]{sandage78,dressler80}. Studies of S0 galaxies in clusters  suggest that interactions between galaxies can lead to the loss of  interstellar matter in the disk and convert spirals into S0s. Other  theories favour gas loss through ram-pressure stripping from the intergalactic medium \\citep{gunn72} or via galactic winds \\citep{faber76b},  or assert that S0 galaxies are primordial galaxies which have entirely  consumed their gas due to a high star formation rate \\citep{larson80}.  When studying the formation of S0 galaxies, one is often drawn to examine the order of events during galaxy assembly. Early formation models favour the formation of the bulge before the disk, which is supported by several photometric studies \\citep[e.g.,][]{caldwell83,terndrup94,pb96,peletier99} and by the fact that S0s behave like ellipticals in the widely known scaling relations (i.e.\\ color-magnitude, Mg$_2$-$\\sigma$ or fundamental plane).  Numerical simulations raised the possibility of an alternative scenario \\citep[e.g.,][]{pfenniger93}, in which bulges are formed via secular evolution processes after the disk \\citep[e.g.,][]{sw93,courteau96,carollo99,bureau99b,merrifield99}.  Although such studies naturally focus on the two major components, i.e.\\ the large-scale disk and the bulge, galaxies often show more structures,  e.g.\\ central disks, stellar kinematically decoupled components (KDCs)  and multiple gaseous components. These are commonly observed in galaxies  all along the Hubble sequence and are often used to emphasize the role of  mergers during the formation of early-type galaxies  \\citep[][]{toomre77,wr78,schweizer86,fi88}.  In the last $15$ years, evidence for the presence of KDCs in early-type  galaxies (e.g.,\\ NGC\\,4406, \\citealt{bender88}; NGC\\,4365, \\citealt{wbm88})  has usually been revealed through long-slit measurements. This can introduce  a priori assumptions about the geometry of the central structures, and  several works have tried to overcome this limitation by mapping the galaxies  with multiple slits at different positions angles  \\citep{fisher94,fisher97,ss99,michele02}. However, the arrival of a new  generation of integral-field units (IFUs) providing contiguous field-of-views  (FOVs), such as \\oasis\\ (e.g., NGC\\,4621, \\citealt{wernli02}) or \\sauron\\  \\citep{bacon01,tim02}, offers a much more efficient way to detect KDCs.  Early \\sauron\\ results \\citep{bacon01,tim02} reveal a variety of  structures much richer than usually recognized in early-type galaxies.  It is therefore important to use the information contained in these  features to elucidate the key processes at work during galaxy  formation. The galaxies concerned can become the ideal benchmarks  against which to test the predictions of galaxy formation and evolution  theories. NGC\\,7332, the galaxy studied in this paper, may be such a keystone.  NGC\\,7332, a boxy edge-on S0 galaxy\\footnote{The galaxy is classified  S0p in the RC3 \\citep{RC3} due to its boxy bulge \\citep{sandage61}.},  was studied extensively in the past. The galaxy is mainly known for a  bright counter-rotating and a faint co-rotating \\oiii\\ gas component with  respect to the stars \\citep{bertola92,fisher94}. Those gas structures were  confirmed by \\citet{plana96}, who mapped the \\halpha\\ emission via  Fabry-Perot observations. NGC\\,7332 exhibits a rather regular broad-band  morphology \\citep{pb97} and its colours are somewhat bluer than those of  elliptical galaxies of the same luminosity. Spectral analysis of the central  regions reveals a luminosity-weighted age of about $6$~Gyr \\citep{vazari99}.  In this paper, we use \\sauron\\ observations to characterize new and known features and improve our understanding of the morphology and dynamics of NGC\\,7332. We focus the analysis on the distribution and kinematics of  the stellar and gaseous components, but we also discuss the corresponding  \\sauron\\ line strengths maps. We present strong evidence for the presence  of a central KDC and complex two-dimensional (2D) gas kinematics.  The main characteristics of NGC\\,7332 are provided in Table~\\ref{Tab:Params}. Section~\\ref{Sec:Obsdata} summarises the observations discussed in this paper and their reduction. A photometric analysis and comparison with the literature is presented in Section~\\ref{Sec:Phot}. In Sections~\\ref{Sec:Skinem} and \\ref{Sec:Gkinem}, we present and discuss the kinematic maps of stars and gas, and describe the methods used to disentangle their respective spectral contributions. The line-strength analysis is presented in Section~\\ref{Sec:Popu}. We finally discuss possible formation and evolution scenarios for NGC\\,7332 in Section~\\ref{Sec:Discussion} and summarise our conclusions in Section~\\ref{Sec:Conclusions}. ",
        "conclusions": "\\label{Sec:Conclusions}  We presented in this paper new \\sauron\\ integral-field spectroscopy observations of the edge-on S0 galaxy NGC\\,7332, along with a discussion of existing ground and space-based imaging.  The photometric analysis reveals a boxy bulge, a central disk and  evidence for the presence of a bar (Section~\\ref{Sec:Phot}; but see  also \\citealt{ldp00}). The \\sauron\\ observations provide unprecedented  coverage of the stellar and gas kinematics. In particular, the stellar  kinematics displays a rather smooth velocity field with rotation along  the major-axis and a weak dependence of rotation on galactic height. We  also discovered a stellar kinematically decoupled component (KDC) misaligned  with respect to the galaxy's kinematic major-axis, which may be related  to a dip in the centre of the velocity dispersion map ($\\rm{r}\\lesssim2\\farcs5$). We  finally provided kinematic evidence for the presence of a larger central  stellar disk ($\\rm{r}\\lesssim7\\arcsec$).  NGC\\,7332 presents significant ionised gas emission in the spectral range studied, in particular in the \\hbeta\\ ($4861$~\\AA) and \\oiii\\ ($4959$ and $5007$~\\AA) lines. We used a sophisticated technique to separate the  absorption and emission lines, yielding both pure absorption spectra from  which to derive the stellar kinematics and pure emission spectra for the gas  kinematics (and distribution). The emission maps reveal a very complex gas  morphology and kinematics. This is found especially in \\oiii\\, which is  mainly counter-rotating with respect to the stars. The gas maps  (Figure~\\ref{Fig:Gaskin}) show significant amounts of gas at high relative  velocities outside of the equatorial plane, whereas the central parts display  rather ordered motions. We also found traces of a faint co-rotating ionised gas  component.  The analysis of the absorption line-strengths in NGC\\,7332 show that its stellar populations are generally young ($5\\pm2$~Gyr), not only in the disk but also in the bulge, in agreement with previous studies \\citep{bp94,vazdekis96,terlevich02}. The metallicity indices (i.e.\\ Fe5015,  Mg$b$, Fe5270) show an increase in the centre, contrasting with the  rather homogeneous \\hbeta\\ index. There is also weak evidence of a  decrease at $r\\approx10$\\arcsec, between the central and the outer disks.  The existence of a large-scale bar in NGC\\,7332 can simultaneously explain  most of the features found in the galaxy. The boxy morphology can be  explained by the thickness of the bar after buckling, which also leads  to homogeneous stellar populations across the bulge and disk. The stellar  kinematics can similarly be explained by the particular orbital structure  of barred disks \\citep[see][]{ba04}. The KDC, on the other hand, was most  likely formed through late gas infall, an accretion event perhaps related  to the complex ionised gas distribution currently seen.  As emphasized by the unique data set presented in this paper, the S0 galaxy NGC\\,7332 possesses a number of peculiar properties helping to unravel  its formation and evolution processes. There is strong observational evidence  that interactions played a role in the shaping of NGC\\,7332, while the bar  has and is certainly significantly affecting its evolution. Given the amount  of gas yet to settle and the relatively young age of the stellar populations  in NGC\\,7332, it probably represents an earlier evolutionary stage of more  regular edge-on S0 galaxies (i.e.\\ NGC\\,4570, \\citealt{vdbemsellem98};  NGC\\,3115, \\citealt{eric03}; NGC\\,3377, \\citealt{bacon01}). The morphology  and dynamics of present-day S0 galaxies can thus probably be understood  through the combined effects of both interactions and bar-driven processes."
    },
    "0402/astro-ph0402539_arXiv.txt": {
        "abstract": "The vertical structure of black hole accretion disks in which radiation dominates the total pressure is investigated using a three-dimensional radiation-MHD calculation.  The domain is a small patch of disk centered 100 Schwarzschild radii from a black hole of $10^8$ M$_\\odot$, and the stratified shearing-box approximation is used.  Magneto-rotational instability converts gravitational energy to turbulent magnetic and kinetic energy.  The gas is heated by magnetic dissipation and by radiation damping of the turbulence, and cooled by diffusion and advection of radiation through the vertical boundaries. The resulting structure differs in several fundamental ways from the standard Shakura-Sunyaev picture.  The disk consists of three layers. At the midplane, the density is large, and the magnetic pressure and total accretion stress are less than the gas pressure.  In lower-density surface layers that are optically thick, the magnetic pressure and stress are greater than the gas pressure but less than the radiation pressure.  Horizontal density variations in the surface layers exceed an order of magnitude.  Magnetic fields in the regions of greatest stress are buoyant, and dissipate as they rise, so the heating rate declines more slowly with height than the stress.  Much of the dissipation occurs at low column depth, and the interior is cooler and less radiation-dominated than in the Shakura-Sunyaev model with the same surface mass density and flux.  The mean structure is convectively stable. ",
        "introduction": "Black holes in X-ray binary star systems and active galactic nuclei accrete material through surrounding disks of gas.  When the accretion rate is greater than about 1\\% of the Eddington value, the radiation pressure in the central regions of the disk may exceed the gas pressure \\citep{ss73}.  The disk structure and the radiation emerging are governed by removal of angular momentum from the gas, and conversion of released gravitational energy to photons.  In the standard model, the angular momentum transfer and dissipation processes are unspecified.  Both are assumed to occur at rates proportional to the height-integrated total pressure, and the heating is balanced by diffusion of photons to the disk surfaces \\citep{ss73}. The model is unstable to perturbations in mass accretion rate \\citep{le74} and heating rate \\citep{ss76}, and to vertical convective displacements \\citep{bkb77}.  Any one of these instabilities might substantially modify the disk structure.  However, if the stresses are magnetic in nature, buoyancy of the fields may lead to stresses scaling with the gas pressure alone, rather than the total pressure \\citep{sc89}.  In this case the viscous and thermal instabilities are absent.  The instabilities may also be eliminated if part of the accretion energy is dissipated outside the disk, in a hot corona \\citep{sz94}.  Observed properties that are not understood using the standard Shakura-Sunyaev model include the X-ray emission \\citep{em78} and UV-optical spectra \\citep{b04} of active galactic nuclei, and the steady disk luminosities of some stellar-mass black hole X-ray binary systems accreting at 1\\% to 50\\% of the Eddington rate \\citep{gd04}.  The most likely physical mechanism of angular momentum transfer in radiation-dom\\-in\\-at\\-ed disks is magnetic stresses.  The magneto-rotational instability or MRI \\citep{bh91} leads to turbulence in which the energy of differential orbital motion is converted to magnetic fields and gas motions \\citep{hgb96}.  In isothermal MHD simulations in which heating and cooling are assumed to balance, magnetic fields generated in the turbulence are partly expelled to form magnetized coronae above and below the disk \\citep{ms00}.  Disk structure may be affected also by the locations and rates of heating and cooling.  Candidate physical heating mechanisms include resistive dissipation of magnetic fields, microscopic viscous dissipation of gas motions, and radiative damping of compressible turbulence \\citep{ak98}.  Cooling processes may include radiation diffusion, convection \\citep{ak01}, and photon bubble instability \\citep{g98}.  Here I examine the effects of heating and cooling on the vertical structure of radiation-dominated disks, using a radiation-MHD simulation. ",
        "conclusions": "A patch of radiation-dominated accretion disk lying 100 Schwarzschild radii from a $10^8$ M$_\\odot$ black hole is simulated including physical processes of energy release, dissipation, and cooling.  The patch has surface density $10^6$ g~cm$^{-2}$, and its net vertical magnetic flux is zero.  The structure that develops has density greatest at the midplane, stresses greatest in optically-thick surface layers, and dissipation found throughout but intermittent in time and space.  The stresses result from magnetic forces.  The fields in the surface layers are buoyant.  Heating occurs mainly by numerical dissipation of the fields, and physical radiation damping of the turbulence, while cooling occurs by diffusion and advection of radiation through the boundaries.  The vertical structure may prove to vary with mass of the central body, location in the disk, surface density, and magnetic flux.  The time-averaged structure in the simulation is hydrodynamically convectively stable.  Thermal instability may be prevented by outflow through the boundaries.  Viscous stability is not tested, as there is no net radial flow in the shearing box.  Photon bubble instability may be present.  The fastest linear modes \\citep{bs03} in the time-averaged structure are found in the surface layers, have wavelengths shorter than the vertical grid spacing, and grow at $0.9\\Omega_0$.  Resolved modes, with wavelengths greater than six grid zones, are expected to grow more slowly than the MRI.  No clear signs of these modes are found.  The likely direct effect of photon bubbles is to further reduce the cooling time."
    },
    "0402/astro-ph0402013_arXiv.txt": {
        "abstract": "The presence of a kHz frequency in LMXBs has been expected from scaling laws, by analogy with the QPO phenomenon in HMXB X-ray pulsars. Interpretation of the two kHz frequencies, observed in accreting neutron stars,  in terms of non-linear resonance in strong-field gravity led to the prediction of twin  QPOs in black hole systems, in a definite frequency ratio (such as 2/3). The imprint of a subharmonic of the 401 Hz rotation rate in the frequencies of the QPOs detected in the accreting millisecond pulsar is at once a signature of non-linear resonance and of coupling between accretion disk modes and the neutron star spin. ",
        "introduction": "This is a story of the origin of high-frequency quasi-periodic oscillations (HF QPOs) observed in the X-ray flux of low-mass X-ray binaries (LMXBs). We believe that they really are oscillations, resonant oscillations of a body of fluid nearly in equilibrium---the accretion disk. Accretion disk are so poorly understood that their detailed numerical and analytic modeling has failed to reveal phenomena approximating the behavior of real HF QPOs. Even so, they must exhibit behavior generally observed in any sufficiently complicated system, e.g., resonances.  Physics has principles and techniques allowing a meaningful discussion of the behavior of systems whose detailed structure is not known. The application of these principles has allowed us to successfully predict \\cite{kmw,ka00} millisecond variability in LMXBs, at a time when at most $\\sim 50\\,$Hz oscillations had been seen, and the appearance of pairs of hHz QPOs with rational (as in 3:2) frequency ratios in black holes, at a time when only single QPOs had been reported. We see no reason to abandon these principles now, when the field of QPO modeling  has been thrown into turmoil with the discovery \\cite{rudy} of a subharmonic frequency difference in the accreting millisecond pulsar. This is where they take us. High-frequency QPOs are an accretion disk phenomenon reflecting fundamental properties of gravity. A spinning neutron star may directly perturb the accretion disk and influence the frequencies of its oscillation via a non-linear coupling. ",
        "conclusions": ""
    },
    "0402/astro-ph0402296_arXiv.txt": {
        "abstract": "{% CMB anisotropy data could put powerful constraints on theories of the evolution of our Universe.  Using the observations of the large number of CMB experiments, many studies have put constraints on cosmological parameters assuming different frameworks. Assuming for example inflationary paradigm, one can compute the confidence intervals on the different components of the energy densities, or the age of the Universe, inferred by the current set of CMB observations. The aim of this note is to present some of the available methods to derive the cosmological parameters with their confidence intervals from the CMB data, as well as some practical issues to investigate large number of parameters.  }{% Les observations des anisotropies du fond diffus cosmologique (FDC) peuvent placer de fortes contraintes sur les th\\'eories d'\\'evolution de notre Univers. L'utilisation de telles donn\\'ees a permis de contraindre differents param\\`etres de diff\\'erents cadres th\\'eoriques: l'age de l'Univers, son contenu baryonique, etc. Le but de cette contribution est de pr\\'esenter differentes m\\'ethodes possibles pour extraire les param\\`etres cosmologiques et leurs intervalles de confiance des donn\\'ees du FDC. Des questions pratiques sur l'utilisation de grands nombres de param\\`etres sont aussi abord\\'ees.} ",
        "introduction": "The extraction of information from cosmic microwave background (CMB) anisotropies is a classic problem of model testing and parameter estimation, the goals being to constrain the parameters of an assumed model and to decide if the {\\em best--fit model} (parameter values) is indeed a good description of the data. Maximum likelihood is often used as the method of parameter estimations.  Within the context of the class of models to be examined, the probability distribution of the data is maximized as a function of the model parameters, given the actual, observed data set\\footnote{In recent Bayesian analyses, quoting the mean of the product of the likelihood and prior functions as best model, is preferred}. Once found, the best model must then be judged on its ability to account for the data, which requires the construction of a {\\em statistic} quantifying the {\\em goodness--of--fit} (GoF). Finally, if the model is retained as a good fit, one defines {\\em confidence} intervals on the parameter estimation.  The exact meaning of these confidence intervals depends heavily on the method used to construct them, but the desire is always the same -- one wishes to quantify the `ability' of other parameters to explain the data (or not) as well as the best fit values. Given the quality of the current data, and the aim of the analysis -- precise determination of the cosmological parameters -- much attention should be put on the robustness and accuracy (unbiased techniques) of the methods used.  I review the different ways of estimating the likelihood function of the parameters focusing on the use of the angular power spectrum ($C_\\ell$'s). Then, some methods to compute the goodness of fit and the confidence intervals will be discussed. Finally, some practical issues for such computations will be addressed. In this review, I take the temperature fluctuations as the observed quantity. The same approaches could be applied for the polarisation signal of the CMB. ",
        "conclusions": "In order to derive the cosmological parameters in a given framework from the temperature fluctuation of the CMB, many steps are needed. When the observed power spectrum is derived, one could use different techniques to estimate successively the (approximated) likelihood value of the family of models (parameters) investigated, the best model and its goodness of fit, and finally the confidence intervals on each parameter. Each of these steps may be highly cpu and memory consuming. With better and better observations, sensitivity and sky coverage, brute force maximum likelihood methods become impossible. Many approximations and techniques have then been developed during the last years, allowing to analyze more and more data with increasing speed. When the appropriate method is used, this leads to an unbiased estimate of the cosmological parameters. These developments have demonstrated that efficient methods could be developed to take full advantage of data at the Planck accuracy and allow to determine parameters of cosmological relevance to a remarkably high accuracy.  This is opening the golden road of precision cosmology."
    },
    "0402/astro-ph0402543_arXiv.txt": {
        "abstract": "We present the results of a 3.3-year project to monitor the flux density of Sagittarius~A* at 2.0, 1.3, and 0.7~cm with the Very Large Array.  Between 2000.5 and 2003.0, 119 epochs of data were taken with a mean separation between epochs of eight days.  After 2003.0, observations were made roughly once per month for a total of nine additional epochs.  Details of the data calibration process are discussed, including corrections for opacity and elevation effects as well as changes in the flux density scales between epochs.  The fully calibrated light curves for Sgr~A* at all three wavelengths are presented.  Typical errors in the flux density are 6.1\\%, 6.2\\%, and 9.2\\% at 2.0, 1.3, and 0.7~cm, respectively.  There is preliminary evidence for a bimodal distribution of flux densities, which may indicate the existence of two distinct states of accretion onto the supermassive black hole.  At 1.3 and 0.7~cm, there is a tail in the distribution towards high flux densities.  Significant variability is detected at all three wavelengths, with the largest amplitude variations occurring at 0.7~cm.  The rms deviation of the flux density of Sgr~A* is 0.13, 0.16, and 0.21~Jy at 2.0, 1.3, and 0.7~cm, respectively.  During much of this monitoring campaign, Sgr~A* appeared to be relatively quiescent compared to results from previous campaigns.  At no point during the monitoring campaign did the flux density of Sgr A* more than double its mean value.  The mean spectral index of Sgr~A* is $\\alpha=0.20\\pm0.01$ (where $S_\\nu\\propto\\nu^\\alpha$), with a standard deviation of 0.14.  The spectral index appears to depend linearly on the observed flux density at 0.7~cm with a steeper index observed during outbursts.  This correlation is consistent with the expectation for outbursts that are self-absorbed at wavelengths of 0.7~cm or longer and inconsistent with the effects of simple models for interstellar scintillation.  Much of the variability of Sgr~A*, including possible time lags between flux density changes at the different wavelengths, appears to occur on time scales less than the time resolution of our observations (8 days). Future observations should focus on the evolution of the flux density on these time scales. ",
        "introduction": "}  Observations of stellar proper motions in the central $1''$ (0.04~pc) of the Milky Way suggest that a $4\\times10^6$~M$_\\odot$ black hole is located at the dynamical center of the Galaxy \\citep{sch02,ghe03b}. In the radio, the bright ($\\sim1$~Jy), compact source called Sagittarius A* (Sgr~A*) appears to be closely associated with the supermassive black hole \\citep{men97}.  For over two decades, the radio flux density of Sgr~A* has been known to vary \\citep{bro82}, but the cause of this variability remains unclear.  The radio variability tends to increase towards shorter wavelengths, and significant fluctuations on roughly weekly time scales are observed at wavelengths of 2~cm and shorter \\citep{zha01}.  Based on this observed increase in fractional variability towards shorter wavelengths, the variability of Sgr~A* at centimeter wavelengths has been suggested to be intrinsic to the source \\citep{zha92}. At millimeter and sub-millimeter wavelengths, the flux density of Sgr~A* is even more variable \\citep{wri93, tsu99,zha03}.  Changes of up to a factor of four in the flux density have been observed at 1.3~mm using the Sub-Millimeter Array (SMA) \\citep{zha03}.  Short-term variability of Sgr~A* has been detected in X-rays.  In 2002, a multi-wavelength campaign was undertaken to simultaneously monitor Sgr~A* at centimeter, millimeter, infra-red, and X-ray wavelengths.  Based on 500 kilo-seconds of Chandra observations, \\citet{bag03} finds that flares of a factor of $\\sim10$ occur roughly once per day.  (In this paper, we call events on hour-to-day time scales ``flares'', while events on time scales of weeks are referred to as ``outbursts''.)  The strongest X-ray flare during this observation occurred on 29 May 2002 and showed a factor of 20 increase in the flux of Sgr~A*.  To date, three additional large X-ray flares have been observed between June 2000 and October 2003.  On 27 October 2000, the flux of Sgr~A* detected by Chandra increased by a factor of 45 during a 10~ksec flare \\citep{bag01}.  A factor of 20 increase in the flux of Sgr~A* at 2--10~keV was detected by XMM-Newton during the beginning of a flare on 4 September 2001, but the observations unfortunately did not include the entire event \\citep{goldwurm03}. More recently, a factor of 160 flare lasting 2.7~ks was detected with XMM-Newton on 3 October 2002 \\citep{porquet03}.  Recently, Sgr~A* has been detected for the first time at infrared wavelengths.  In the infrared, Sgr~A* appears to show both long-term and short-term variability.  The short-term variability occurs on time scales of $\\sim1$~hr, similar to the time scale of X-ray events \\citep{gen03}.  The day to week time scale of the long-term infrared variability is more similar to the time scale of radio outbursts \\citep{ghe03b}.  In order to understand the detailed nature of the variability of Sgr~A* in the radio, \\citet{zha01} (hereafter ZBG01) combined Very Large Array\\footnote{The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.}  (VLA) data from frequent monitoring of Sgr~A* between 1990.1 and 1993.5 \\citep{zha92} with twenty years of additional archival data from 1977 to 1999.  At 1.3~cm, the flux density of Sgr~A* typically varied by 30\\%, and occasional outbursts of 100\\% were observed.  A comparison of the light curves at 20, 6.0, 3.6, 2.0, and 1.3~cm showed that the largest amplitude variations in flux density occurred at the shortest wavelengths (ZBG01).  Using a Maximum Entropy Method \\citep{pre89} as well as a classic periodogram augmented with CLEAN to search for periodicities, ZBG01 found a periodicity of $106\\pm10$ days ($1.1(\\pm0.1)\\times10^{-7}$~Hz) at 3.6, 2.0, and 1.3~cm.  The light curves were consistent with no phase offset between the three wavelengths.  The mean profile of the 106 day cycle at 1.3~cm had a broad peak, roughly 25\\p5 days wide at full width at half maximum (FWHM), indicating that the variability of Sgr~A* is most likely quasi-periodic in nature.  Analysis of data presented in ZBG01 was hindered, however, by highly irregular sampling intervals.  Observations were only made on a regular basis (with sampling intervals from 1 to 28 days) between 1990.1 and 1991.5, a small fraction of the entire study.  Between 1991.5 and 1993.5, the sampling was less frequent with a maximum sampling interval of 120 days.  Gaps in the archival data were as large as 1200 days.  Additional evidence for periodic variability on the order of 100 days has been detected at 2 and 3~mm using 46 epochs of data taken between 1996 and 2000 with the Nobeyama Millimeter Array (NMA) \\citep{tsu02}. When the data are folded with a period of 106 days, the resulting cycle divides roughly equally into a ``high'' and ``low'' activity state.  During the high activity state, the flux density of Sgr~A* can vary by as much as a factor of three, while Sgr~A* is relatively quiescent in the low activity state.  Evidence for a period of 57 days was suggested by \\citet{fal99} based on 11~cm data from the Green Bank Interferometer (GBI).  This result suggested that the 106-day periodicity could be a harmonic of a higher frequency periodicity.  Due to sparse sampling over much of the 20~yr baseline, the data presented by ZBG01 were not sensitive to periods shorter than $\\sim50$ days.  A new monitoring campaign in which the flux density of Sgr~A* was regularly monitored on weekly time scales was necessary to search for these short periodicities and detect the shape of outbursts and phase offsets between the different wavelengths.  From June 2000 to October 2003, over 170 hours of VLA observing time were used to monitor the flux density of Sgr~A* at 2.0, 1.3, and 0.7~cm.  The results of this 3.3-year project are presented in this paper.  Section \\ref{cal} explains in detail the data calibration, including corrections for opacity effects and errors in absolute flux calibration.  The fully calibrated flux densities for Sgr~A* at 2.0, 1.3, and 0.7 cm are presented in \\S\\ref{res} and are also listed in their entirety in Table \\ref{table}.  The characteristics of the light curve, including the spectral index and time delay between wavelengths, are also discussed.  This paper deals primarily with data calibration and the general characteristics of the light curve and spectrum.  Results from the analysis of these data for periodic or quasi-periodic signals will be presented in subsequent papers.  Although \\S \\ref{res} briefly compares our light curves to those from monitoring campaigns at other wavelengths, detailed comparisons are presented elsewhere. \\citet{zha03} compares the radio light curve to preliminary data from the Sub-Millimeter Array.  A separate paper also discusses evidence for correlated events in the radio and X-ray and possible implications for accretion models for Sgr~A* \\citep{zha04}.  Finally, six Very Long Baseline Array (VLBA) observations were made as part of this monitoring campaign.  The results from these observations are presented in \\citet{bower04}. ",
        "conclusions": "In this paper, we have presented results from a 3.3-year campaign to monitor the flux density of Sgr~A* at centimeter wavelengths using the VLA.  The largest amplitude variations are observed at 0.7~cm, consistent with variability increasing towards shorter wavelengths (ZBG01).  Overall, however, Sgr~A* appears to be more quiescent than during previous monitoring campaigns.  The spectral index of Sgr~A* appears to be strongly correlated with the 0.7~cm flux density.  This result strongly favors an emission mechanism in which outbursts are intrinsic to Sgr~A* and are not the result of interstellar scintillation.  Regular monitoring of Sgr~A* at the VLA will continue through at least June 2004 with monthly observations at 2.0, 1.3, and 0.7~cm.  These additional data will be useful in detecting periods longer than 100 days.  Much of the activity of Sgr~A* appears to take place on time scales $\\simgt1$~hr and less than the time resolution of our data (8 days). However, the monitoring campaigns in the radio and sub-millimeter have only minimally probed these time scales.  To date, hourly time scales have only been systematically probed in X-rays and the infrared.  In both cases, there is significant variability on time scales of hours. Monitoring of the centimeter flux density on these short time scales will be necessary to determine the duration and shape of outbursts as well as detect any time lag between wavelengths.  Additional coordinated multi-wavelength campaigns specifically aimed at probing these time scales should also be implemented.  Only with simultaneous coverage and fine sampling will we be able to ascertain the relationship between radio outbursts and X-ray flares and constrain the emission mechanism for Sgr~A*."
    },
    "0402/astro-ph0402069_arXiv.txt": {
        "abstract": "\\noindent Spatially homogeneous and isotropic cosmological models, with a perfect fluid matter source and non-vanishing cosmological constant, are studied. The equations governing linear perturbations of the space-time and the variation of energy density are given. The complete solution of the problem is obtained for $C^{\\infty}$ perturbations, using a comoving time. The Sachs-Wolfe fluctuations of the temperature of the cosmic background radiation are obtained for the relatively growing density perturbations. It is found that the observable celestial microwave fluctuation pattern underwent a reversal approximately two billion years ago. What is observed today is a negative image of the last scattering surface with an attenuation of the fluctuations, due to the presence of the cosmological constant. \\vskip 12pt \\noindent Keywords: {\\it cosmic microwave background -- large scale structure of the universe} ",
        "introduction": "Recent observations by the {\\it High-z Supernova Search Team} \\citep{Tonry} corroborate the data on an accelerating expansion of the Universe. Three independent lines of evidence [those from the {\\it Wilkinson Microwave Anisotropy Probe} ({\\it WMAP}) measurements \\citep{WMAP} of the cosmic microwave background  radiation {\\it (CMBR)}, the {\\it Sloan Digital Sky Survey} ({\\it SDSS}) \\citep{Tegmark} and the observations of type-Ia supernova spectra] converge on the value $\\Omega_{\\Lambda}=0.70\\pm 0.04$ of dark energy. All data are consistent, within a 20\\% error bar, with a time-independent dark energy distribution, as is the case with the cosmological constant $\\Lambda$ and flat space ($k=0$). These developments attracted attention to cosmological models in the presence of a $\\Lambda$ term in Einstein's gravitational equations. Many excellent review papers are available \\citep  {Carroll,Peebles,Carroll2,Silk} on the implications of a cosmological constant in models of the Universe. While the cosmological constant has little effect on the large-scale structure and  dynamics of the Universe, the microwave fluctuations will be significantly different in $\\Lambda \\neq0$ models \\citep{Silk}. The scalar perturbations of $\\Lambda \\neq0$ models with multicomponent fluids (i.e., cold or hot dark matter, photons and massless neutrino) have been computed by \\citet{Bond}, \\citet{Fukugita}, \\citet{Holtzman}, \\citet{Vittorio}, \\citet{HS} and \\citet{Stompor} in the linear approximation. These numerical results have been compared by \\citet{Stompor} and found to be in reasonable agreement. \\citet{HS} claim that for $\\Lambda\\neq0$, one must use a numerical approach. In our work, however, we present an analytic treatment of the collisionless dust which represents the late-time evolution.  A complete solution describing the effect of the cosmological constant on perturbation dynamics is lacking in the literature, this including the rotational and gravitational wave perturbations. The temperature fluctuations of the CMBR have been computed in numerical schemes \\citep{Bond,Fukugita,Holtzman,Vittorio,Stompor} and \\citet{Multamaki} investigate the integrated contributions. Known as the Sachs-Wolfe effect, this is the variation in the redshift of a photon travelling freely in the universe rippled with perturbations. Initial work, both on the zero-pressure model of the current state of the Universe, and on the radiation-dominated era, was done by \\citet{SW}. The pressure-free universe in the presence of a cosmological constant has been considered by \\citet{Heath} and \\citet{Lahav}. They used the redshift parameter to investigate the behaviour of the density contrast. A more recent work by \\citet{Vale} uses the formalism of \\citet{Padmanabhan} to compute linear perturbations of a dust-filled Friedmann-Robertson-Walker (FRW) universe in the presence of a cosmological constant. They find that the presence of $\\Lambda$ inhibits the growth of the fluctuations in time. [A different approach to the perturbation problem of a different (de Sitter) model, using the gauge-independent formulation, has been pursued by \\citet{Barrow}, for the investigation of stability under perturbations.]  Neither of these earlier works reaches a complete perturbative picture. Although the Sachs-Wolfe effect is computed by numerical methods, there is no analytic treatment so far available for $\\Lambda\\neq0$. Our aim here, therefore, is to consider, in an analytic framework, the Sachs-Wolfe effect and its contributions to the fluctuations of the CMBR in the presence of a cosmological constant. We follow the treatment of the perturbations due to \\citet{White}, thereby relaxing the momentum conditions of the original Sachs-Wolfe work. Both works yield all $C^\\infty$ perturbations although this was only shown by \\citet{White}. Correspondingly, in the present paper we also obtain all $C^\\infty$ perturbations. In Sec. 2, we enlist the perturbed field equations using the conformal time $\\eta$. While retaining the form of the metric functions, we change to the comoving time $t$ in the field equations so obtained. The reason for this is that the unperturbed radius $a$ has a simple analytic dependence on $t$.  The general solution of the linear field equations in a dust-filled universe is obtained in Sec. 3. We find that all first-order fields, except the wave solutions, can be expressed in terms of only two complex incomplete elliptic integrals $E$ and $F$. This is achieved by using the mirror symmetries of the elliptic integrals.  In Sec. 4, we obtain the Sachs-Wolfe effect and its integral contribution which vanishes in the $\\Lambda\\rightarrow 0$ limit. An unexpected feature of the transfer function is that it changes sign and reverses the pattern of the temperature fluctuation on the celestial sphere at late times. In Sec. 5, we investigate the physical interpretation of our solution. Taking the current experimental values of the cosmological parameters, we obtain the comoving time elapsed between the emission at the surface of last scattering and reception of the photon. We find that the Sachs-Wolfe contribution to the CMB fluctuations in the presence of the cosmological constant is damped by a factor of $3$ today, because of the proximity, on the cosmological time scale, of the moment of reversal of the microwave temperature fluctuations and the reception. This feature of the SW effect has remained concealed in earlier works, because of the numerical methods involved. ",
        "conclusions": "The results of our work are relevant in two major contexts: first, we find that the observed CMBR power spectrum is attenuated by a significant overall factor due to the presence of the cosmological constant. Our finding is consistent with the results of the  numerical studies in \\cite{HS}. This phenomenon may help understanding structure formation by taking into account a higher level of the initial matter fluctuations. The second observation of interest that we make is the natural complexification of the perturbed fields in the normal picture of Legendre. This complex description may prove useful in a treatment of quantum fluctuations, a phenomenon that occurs also in the Fourier expansion of the wave solutions."
    },
    "0402/astro-ph0402633_arXiv.txt": {
        "abstract": "We have obtained {\\it HST} and ground-based observations of a sample of 20 O-type stars in the LMC and SMC, including six of the hottest massive stars known (subtypes O2-3) in the R136 cluster.  In general, these data include (a) the {\\it HST} UV spectra in order to measure the terminal velocities of the stellar winds, (b) high signal-to-noise, blue-optical data where the primary temperature- and gravity-sensitive photospheric lines are found, and (c) nebular-free H$\\alpha$ profiles, which provide the mass-loss rates.  We find that the older (FOS) {\\it HST} data of the R136 stars (which were obtained without the benefits of sky measurements) suffered from significant nebular emission, which would increase the derived mass-loss rates by factors of $\\sim 3$, all other factors being equal.  We also find several stars in the SMC for which the N~III $\\lambda\\lambda 4634,42$ and He~II$\\lambda 4686$ emission ``f\" characteristics do not appear to follow the same pattern as in Galactic stars. Since He~II emission is due to the stellar wind (which will be weaker in SMC for stars of the same luminosity), while N~III emission is a complex NLTE effect affected mostly by temperature, it would not be surprising to find that these features do not correlate with each other or with luminosity in SMC stars in the same was as they do in Galactic stars, but theory does not provide a clean answer, and analysis of more stars (both SMC and Galactic) are needed to resolve this issue.  The line-blanketed non-LTE atmosphere code FASTWIND was then used to determine the physical parameters of this sample of stars. We find good agreement between the synthetic line profiles for the hydrogen, He~I, and He~II lines in the majority of the stars we analyzed; the three exceptions show evidence of being incipiently resolved spectroscopic binaries or otherwise spectral composites.  One such system is apparently an O3~V+O3~V eclipsing binary, and a follow-up radial velocity study is planned to obtain Keplerian masses.  Although we did not use them to constrain the fits, good agreement is also found for the He~I $\\lambda 3187$ and He~II $\\lambda 3203$ lines in the near-UV, which we plan to exploit in future studies.  Our effective temperatures are compared to those recently obtained by Repolust, Puls \\& Herrero for a sample of Galactic stars using the same techniques.  We find that the Magellanic Cloud sample is 3,000-4,000$^\\circ$K hotter than their Galactic counterparts for the early through mid-O's.  These higher temperatures are the consequence of a decreased importance of wind emission, wind blanketing, and metal-line blanketing at lower metallicities. ",
        "introduction": "\\label{Sec-hist}   Translating the observed characteristics of O and early B stars into physical properties has historically proven to be be very challenging.   Because of their high effective temperatures, most of their flux is in the far UV, where even spacecraft cannot observe (due to interstellar extinction), as the peak of the flux distribution $F_\\lambda$ will be below the Lyman limit for stars hotter than about 32,000$^\\circ$K.  This means that the bolometric corrections are quite significant ($-1$ to $-4$~mags), with a steep dependence on the effective temperature.  Thus getting the effective temperature right is the key to determining the other physical properties of a hot, massive star.   Since we are always observing the fluxes of these stars far on the tail of the Rayleigh-Jeans distribution, the observed colors of these stars show little sensitivity to effective temperature. For instance, the Kurucz (1992) model atmospheres predict that a dwarf ($\\log g=4.0$) with $T_{\\rm eff}$=50,000$^\\circ$K will have $(U-B)_o=-1.151$ and a $(B-V)_o=-0.326$, while a dwarf with $T_{\\rm eff}$=40,000$^\\circ$K will have a $(U-B)_o=-1.124$ and $(B-V)_o=-0.311$ (Massey 1998a). The difference in these colors is well within the error of most photometry, and yet the bolometric correction (BC) is $-4.5$~mag for the former, and $-3.8$~mag for the latter, using the conversion of Vacca, Garmany, \\& Shull (1996). This uncertainty in the BC would result in an uncertainty of 0.15~dex in the log of the mass, using the mass-luminosity relation implied by the Schaller et al.\\ (1992) $Z=0.02$ evolutionary tracks ($L\\sim M^2$; see discussion in Massey 1998a); i.e., 50$\\cal M_\\odot$ vs.~70$\\cal M_\\odot$. Reddening complicates the interpretation of the photometry further, of course, and although reddening-free indices can be constructed assuming an average reddening law, even an ``optimal\" reddening-free index using space-craft accessible UV photometry (such as {\\it F170W}) remains too degenerate to be useful (Massey 1998a).   In principle, spectroscopy should allow us to resolve this degeneracy nicely, as the relative strengths of He~I and He~II (O stars) and Si~IV and Si~III (early B supergiants) are quite sensitive to the effective temperature (and only secondarily to the $\\log g$), as confirmed by non-LTE calculations (Auer \\& Mihalas 1972; Kudritzki, Simon, \\& Hamann 1983; Lennon et al.\\ 1991).  Indeed, a difference of one spectral subtype (O5~V to O6~V, say) roughly corresponds to a difference in effective temperature of 0.02~dex in $\\log T_{\\rm eff}$, and a corresponding difference of only 0.13~mag in the bolometric correction, and hence 0.025~dex in the mass of the star ($50\\cal M_\\odot$ vs.~53$\\cal M_\\odot$) one would deduce from the mass-luminosity relationship.  However, the absolute calibration of the relationship between effective temperature and line strengths (or, equivalently, spectral subtypes) requires reliable stellar atmosphere models. Unfortunately, the physics of these stellar atmospheres is quite complicated. The strong lines are all formed under non-LTE conditions, as first shown by Auer \\& Mihalas (1972), and in addition stellar winds provide a significant source of heating for the photosphere through the backscattering of radiation (Hummer 1982, Abbott \\& Hummer 1985). Each decade has seen an improvement in our understanding of the physics of these stellar atmospheres, along with the development the numerical techniques to include these effects in a model atmosphere. We identify four stages in the evolution of these models and hence in the effective temperature scale of O-type stars:  \\begin{enumerate} \\item {\\it The introduction of non-LTE.} The non-LTE models of Auer \\& Mihalas (1972) and Kudritzki (1975, 1976) were relatively simple by today's standards, in being plane-parallel and including no metals, but in fact required the development of a number of innovative techniques in order to include the effects of non-LTE in line formation. (See the review by Kudritzki \\& Hummer 1990).  The Auer \\& Mihalas (1972) models were the first to correctly reproduce the observed line strengths of the helium and hydrogen spectral lines in O-type stars, as was shown by Conti \\& Alschuler (1971).  This led to the first modern effective temperature scale for O stars, that of Conti (1973). This work was extended to the newly defined O3 spectral class (Walborn 1971b) by Kudritzki (1980) and Simon et al.\\ (1983).   \\item {\\it The introduction of mass-loss.} Abbott \\& Hummer (1985) showed that the presence of stellar winds had a significant effect on the He~I/He~II line ratios and hence on the effective temperatures deduced for O-stars.  Although the photospheric lines  are formed in a nearly static part of the atmosphere, the scattering of radiation by the stellar winds back into the photosphere results in substantial heating of the surface layers, an effect known as ``wind blanketing\".   This effect is quite significant, as a high (but realistic) mass-loss rate would result in a 42,000$^\\circ$K model matching an O3~V star, while the same model without mass loss would match that of an O5.5~V star. In other words, the greater the mass-loss rate, the cooler the effective temperature is for a given spectral subtype.  Conti (1988) revised the effective temperature scale to somewhat lower values, presumably to take this effect into account.  Of special note is the revised effective temperature scale of Vacca et al.\\ (1996), which has generally become the standard against which other work is judged.  This scale was based on modeling drawn from the contemporaneous literature, and as such represented the best that the field had to offer at the time. Although it included {\\it some} work that was based on wind-blanketed models, the vast majority of the data was not (e.g., Herrero et al.\\ 1992), and so the derived temperature scale was significantly higher than Conti (1988).  \\item {\\it The inclusion of hydrodynamics and metals.} The next decade saw substantial improvements in the modeling, with the introduction of spherical extension and a more sophisticated treatment of the stellar wind as well as non-LTE treatment of the metal lines. Rather than just using the mass-loss rates, the hydrodynamics of the stellar winds in both the sub- and supersonic regions was included. These improvements were pioneered by the stellar atmospheres group in Munich (e.g., the {\\it unified model atmospheres} concept; see Gabler et al.\\ 1989), which used them for the ``quantitative spectroscopy of hot stars\" (Kudritzki et al.\\ 1989, Kudritzki \\& Hummer 1990, Kudritzki 1991). (A similar but independent approach was taken by Schaerer \\& Schmutz 1994, who made the first attempt to include the opacity of metal lines.)  Sellmaier et al.\\ (1993) demonstrated that in addition to the effect of wind-blanketing, the stellar winds produced emission that partially filled in the He~I lines, which has a strong effect on the effective temperatures as a function of the He~I/He~II line ratios. Another seminal work from this period was Puls et al.\\ (1996), who analyzed a large sample of Galactic and Magellanic Cloud O stars using UV and optical data. The UV spectra were used to determine the terminal velocities of the stellar winds using the strong resonance doublets N~V $\\lambda\\lambda 1239,43$, Si~IV $\\lambda\\lambda 1394,1403$, and C~IV $\\lambda\\lambda 1548,51$. The mass-loss rates were then determined primarily by observations at H$\\alpha$, combined with an assumption (and occasional adjustment) of $\\beta$, a parameter that characterizes the steepness of the velocity law in the stellar wind (Sec.~\\ref{Sec-term}). The classical MK optical/blue region (3800\\AA\\ to 4900\\AA) still provides the primary diagnostics of the surface gravity (from the hydrogen Balmer-line profiles) and effective temperatures (from the He~I and He~II lines).  \\item {\\it The full inclusion of line blanketing.}  A significant improvement in hot star model atmospheres has been the inclusion of full line blanketing.  The CMFGEN code, described by Hillier \\& Miller (1998) and Hillier et al.\\ (2003) is one such example.  Originally developed and used for fitting emission-line features in the expanding atmospheres of Wolf-Rayet stars, CMFGEN has been only recently used for the analysis of absorption lines in O-type stars (i.e., Martins, Schaerer \\& Hillier 2002, Crowther et al.\\ 2002, Hillier et al.\\ 2003, Bouret et al.\\ 2003). Although only a few stars have been fit, these studies suggest that the effective temperature scale of Vacca et al.\\ (1996) is too high\\footnote{This would be expected in any event given the lack of wind-blanketed model fits in the literature at the time of the Vacca et al.\\ (1996) study.}.  (See also discussion in Martins et al.\\ 2002). Similarly, WM-basic, a code developed by Pauldrach, Hoffmann, \\& Lennon (2001), includes full line-blanketing, and its use by Bianchi \\& Garcia (2002) also suggests that a lowering of the effective temperature scale is in order. (WM-basic lacks the Stark broadening and the co-moving frame treatment needed to compute useful synthetic spectra of the quasi-photospheric hydrogen and helium lines, but is very useful to fitting the metal lines found in the UV spectrum.) FASTWIND (``Fast Analysis of STellar atmospheres with WINDS\"), first described by Santolaya-Rey, Puls, \\& Herrero (1997), has now been modified to include an approximate---but highly realistic---treatment of line-blocking, with a similar effect on the effective temperatures (Herrero, Puls, \\& Najarro 2002, Repolust, Puls, \\& Herrero 2004).\\footnote{The use of an {\\it approximate} treatment of line-blanketing and blocking is necessitated by the need for reasonable computational times for a model.  We note for comparison that a single run of CMFGEN, with its more rigorous treatment, requires 9 hours on a 1.3-GHz Pentium IV processor, according to Smith, Norris, \\& Crowther (2002).   By contrast, a single run of FASTWIND requires less than 7 minutes on the slower 750-Mhz SparcIII machine we used in this study. In general, something like 10 to 30 models are needed to fit a star. Comparisons of the flux distribution of FASTWIND with those of WM-basic (Repolust et al.\\ 2004) and CMFGEN (Herrero et al.\\ 2002) have so far shown very good agreement.}  \\end{enumerate}  In this series of papers, we will use new observational data with the latest generation of model atmosphere code to derive a new effective temperature scale for O and early B-type stars, exploring for the first time the effect that metallicity plays on the effective temperature scale (and other derived physical parameters) of these stars.  This is crucial for many astrophysical applications, such as deriving the initial mass function from H-R diagrams (Massey 2003), or in modeling expanding shells, super-bubbles, and normal H~II regions, where having an accurate census of the ionizing flux and amount of mechanical energy being supplied to the region by stars is crucial (Oey \\& Kennicutt 1997).   The series will eventually encompass hot stars in the SMC, LMC, Milky Way, and the Andromeda Galaxy, which span a factor of  6.8 in metallicities, at least as measured by the oxygen abundances (see Massey 2003 and references therein).  Here we begin by studying a sample of hot stars in the Magellanic Clouds.  Analysis of  a second sample of Magellanic Cloud stars is currently underway, and will be published shortly. ",
        "conclusions": "We have attempted to model the spectra of 20 O-type stars in the Magellanic Clouds, succeeding in obtaining adequate or good fits in 17 cases. We suspect that the other three stars are binaries, and expect that a few more in our sample are in reality composite spectra, given the statistical expectations that a third of (Galactic) O-type stars are spectroscopic binaries.  The physical parameters of the stars which were successfully modeled are well determined, and should permit us to refine the effective temperature scale.  In describing the spectral features in our sample, we were struck by the fact that many of our SMC stars appear to be more luminous than their ``f\" characteristics (N~III $\\lambda 4634, 42$ and He~II $\\lambda 4686$) would indicate.  Although one cannot rule out a binary explanation, we did not see such discrepancies amoung our LMC sample, suggesting that the effects of lower metallicity may be responsible. N~III emission is a complex NLTE effect dependent primarily on temperature, while He~II $\\lambda 4686$ is formed in the stellar wind, and its formation will be affected by the temperature, wind density, and amount of EUV flux, all of which will be affected by the metallicity.  Thus one would not expect {\\it a priori} that these ``f\" characteristics would exhibit the {\\it same} behavior at low metallicity as they do in Galactic objects. We have identified here everal examples where the He~II emission is weaker than one might expect given the star's $M_V$, while the referee Nolan Walborn has kindly reminded us of two additional examples in the SMC, Sk~80 and AV~83, where the He~II $\\lambda 4686$ emission is possibly {\\it stronger} than in many Galactic counterparts.  More modeling of both Galactic and SMC stars are needed to understand this issue.   In the second paper of this series we will roughly double the sample of Magellanic Cloud stars that have been analyzed in this manner.  However, it is tempting to compare the effective temperatures we have derived here with those recently found for Galactic stars by Repolust et al.\\ (2004).  We show this comparison in Fig.~\\ref{fig:results}, where we have separated the effective temperatures by luminosity class.  For comparison, we also show the effective temperature scales determined by Conti (1988) (dashed line) and Vacca et al.\\ (1996).  The data for the giants are too sparse as yet to draw any conclusions. However, for both the dwarfs and the supergiants we see that the Magellanic Cloud stars have {\\it significantly} higher temperatures than their Galactic counterparts in the range from the earliest types through mid-O. By the late O-types, there is less of a difference. The data are also too sparse yet to draw firm conclusions about differences in the effective temperature scales for O stars for the SMC vs.\\ LMC. In the sample we have analyzed so-far we have mainly very early O-type stars in the LMC, and somewhat later types in the SMC.  We will be able to address this more fully in the second paper in this series, where we complete our Magellanic Cloud sample.   Nevertheless, this result is quite intriguing.  As described in Sec.~\\ref{Sec-hist}, we would expect the lower mass-loss rates (due to the lower metallicities) in the Clouds to result in a higher effective temperature in comparison to a star of the same spectral subtype in the Milky Way: first, there will be less wind emission affecting the He~I.  In addition, the smaller metallicity will lead to reduced {\\it wind-blanketing} and {\\it line-blanketing}.  Together, these combine to result in effective temperatures which are {\\it $\\sim3000-4000^\\circ$K} (10\\%) greater at O5~V for the Magellanic Cloud stars in our sample.  We note with some irony that despite the substantial improvements in the stellar models over the years, the older effective temperature scales of Conti (1973) and Conti (1988), based primarily on the original non-LTE models of Auer \\& Mihalas (1972), have held up remarkably well for the Galactic stars.  Work over the next few years will result in an improved scale that will take metallicity into account, but for now the Conti (1988) scale is to be preferred over more recent editions (e.g., Vacca et al.\\ 1996).  Three recent studies have analyzed a limited sample of Magellanic Cloud O stars using CMFGEN. Crowther et al.\\ (2002) studied four extreme O supergiants (luminosity class ``Iaf+\") in the Clouds, and derived temperatures that are considerably lower than those shown in Fig.~\\ref{fig:results}a for our sample. Hillier et al.\\ (2003) analyzed the SMC stars extreme supergiant Az~83 (O7 Iaf+) and Az 69 (OC7.5 III((f))). The supergiant is also cooler than what we find, but this may be due to the fact that the more extreme supergiants have lower effective temperatures. The giant star agrees well with the Galactic giants. Bouret et al.\\ (2003) have examined the spectra of five O dwarfs and one O-type giant in the NGC~346 cluster in the SMC. Here the results are more mixed. The temperature of the star they classify as an O2~III is in accord with what we find (Fig.~25b), as are some of the results for the dwarfs.  Others of their dwarfs have effective temperatures lower than what Repolust et al.\\ (2004) found for Galactic stars of the same spectral type. Are the differences due to the different models being employed? We withhold any judgement until a similar sized sample of stars have been analyzed with both codes.  In the second paper in this series we will complete our Magellanic Cloud sample.  At that time, we will also examine the effect that the new data have on the wind-momentum luminosity relation, and give a comparison between the masses derived from the present atmosphere fits and those of stellar evolutionary models."
    },
    "0402/astro-ph0402319_arXiv.txt": {
        "abstract": "We discuss the observational properties of a remarkably faint triply-imaged galaxy revealed in a deep $z'$-band Advanced Camera for Surveys observation of the lensing cluster Abell 2218 ($z=$0.175). A well-constrained mass model for the cluster, which incorporates the outcome of recent Keck spectroscopic campaigns, suggests that the triple system arises via a high redshift ($z>6$) source viewed at high magnification ($\\simeq\\times$25). Optical and infrared photometry from Hubble Space Telescope and the Keck Observatory confirms the lensing hypothesis and suggests a significant discontinuity occurs in the spectral energy distribution within the wavelength interval 9250--9850\\AA. If this break is associated with Gunn-Peterson absorption from neutral hydrogen, a redshift of 6.6$\\,<z<$\\,7.1 is inferred. Deep Keck spectroscopy conducted using both optical and infrared spectrographs fails to reveal any prominent emission lines in this region. However, an infrared stellar continuum is detected whose decline below 9800\\AA\\ suggests a spectroscopic redshift towards the upper end of the range constrained photometrically, i.e. $z\\simeq$7. Regardless of the precise redshift, the source is remarkably compact ($\\lapprox 1\\,h_{70}^{-1}$kpc) and faint ($z_{F850LP}=$28.0) yet is undergoing vigorous star formation at a rate $\\simeq$2.6~$M_{\\odot}$yr$^{-1}$. An intriguing property is the steep slope of the ultraviolet continuum implied by the photometry which may suggest that the source is representative of an early population of galaxies responsible for cosmic reionization. Independent verification of these results is highly desirable but our attempts highlight the difficulty of studying such sources with present facilities and the challenges faced in pushing back the frontiers of the observable universe beyond $z\\sim$6.5. ",
        "introduction": "Reionization was a landmark event which imprinted a signature over the scale of the entire universe. After decades of lower limits on the redshift at which it occurred, recent observations of QSOs discovered by the Sloan Digital Sky Survey (Becker et al.\\ 2001; Djorgovski et al.\\ 2001, Fan et al.\\ 2002) suggest that reionization was just finishing at $z\\sim6-6.5$. The discovery of $z\\simeq6.5$ galaxies with strong Lyman~$\\alpha$ emission (Hu et al.\\ 2002, Kodaira et al.\\ 2003) is illustrative of possible sources which may be responsible. Analysis of recent temperature and polarization fluctuation data from the \\textit{WMAP} satellite suggests an optical depth for Thompson scattering of $\\tau=0.17\\pm$0.04, implying that reionization began at higher redshift, perhaps as early as $z\\sim 15-20$ (Kogut et al.\\ 2003, Spergel et al.\\ 2003).  The discovery of star-forming galaxies at $z\\simeq6.5$ is an important step toward understanding the nature of the sources responsible for the end of cosmic reionization. However, to explore the earlier stages implied by the \\textit{WMAP} results, it is necessary to push the search for star forming systems to higher redshifts. With current facilities this is technically very challenging. It is perhaps salutary to note that the current redshift frontier ($z=6.5$, corresponding to an observed Lyman~$\\alpha$ wavelength of 9200\\AA ) is coincident with the wavelength at which optical CCD detectors fall significantly in their quantum efficiency. Ground-based infrared spectroscopy, necessary for exploring sources at higher redshift, is especially difficult at faint limits.  Color-based searches for $z>$6 sources with the Advanced Camera for Surveys ({\\em ACS}) on-board the Hubble Space Telescope ({\\em HST}) is now recognized as a valuable way of locating $z>6$ sources. Promising results have been obtained by utilizing the long wavelength F850LP ($z'$-band) filter in conjunction with deep infrared imaging on {\\em HST} or with large ground-based telescopes (Bouwens et al.\\ 2003, Yan et al.\\ 2003, Stanway et al.\\ 2003, Dickinson et al.\\ 2003). In view of significant contamination of red `drop-out' samples by cool stars (e.g. Stanway et al.\\ 2003), the primary challenge lies in spectroscopically verifying faint candidates (e.g. Dickinson et al.\\ 2000). Neither optical nor infrared spectrographs on the current generation of ground-based telescopes may have the sensitivity to give convincing results unless strong emission lines are present. Although most of the distant sources found beyond $z\\simeq$5 have been identified via strong Lyman~$\\alpha$ emission (Stern \\& Spinrad 1999, Spinrad 2003), some sources should have weak or no Lyman~$\\alpha$ emission (e.g. Spinrad et al.\\ 1998). Indeed, a significant fraction ($\\sim 75$\\%) of the most intensely star-forming galaxies located by color selection techniques at $z\\simeq3$, reveal Lyman $\\alpha$ only in absorption (Shapley et al.\\ 2003).  Gravitational magnification by foreground clusters of galaxies, whose mass distributions are tightly constrained by arcs and multiple images of known redshift, has already provided new information on the abundance of high redshift objects (Kneib et al 1996, Santos et al.\\ 2003). Particularly high magnifications ($\\simeq\\times$10-50) are expected in the {\\it critical regions} which can be located precisely in well-understood clusters for sources occupying specific redshift ranges (Ellis et al.\\ 2001). Although the volumes probed in this way are far smaller than those addressed in panoramic narrow band surveys (Hu et al.\\ 1998, Malhotra et al.\\ 2001, Hu et al.\\ 2003) or the color-based surveys cited earlier, if the surface density of background sources is sufficient, lensing may provide the necessary boost for securing the first glimpse of young cosmic sources beyond $z\\simeq$6.5 (Santos et al.\\ 2003).  In the course of studying the detailed rest-frame properties of the image pair of a lensed $z=5.576$ galaxy in the cluster Abell~2218 (Ellis et al 2001), we have discovered a new faint pair of images in a deep $z'$-band {\\em ACS} observation of this cluster. Based on the geometrical configuration of this pair and its photometric properties, we concluded that the images most likely arise via strong magnification of a distant $z>$6 source. The implied high redshift of this source led us to explore its properties in more detail.  A plan of the paper follows. We present the photometric observations in Section 2. Section 3 discusses redshift constraints determined independently from the lensing model of Abell 2218 and the spectral energy distribution based on broad-band photometry. Section 4 summarizes our attempts to detect Lyman~$\\alpha$ emission spectroscopically and discusses the implications of a continuum discontinuity seen in the infrared spectrum. We discuss the source properties and implications further in Section 5. Throughout we assume a cosmological model with $\\Omega_{\\rm M}=0.3$, $\\Omega_\\Lambda=0.7$ and H$_0$=70 km~s$^{-1}$~Mpc$^{-1}$. ",
        "conclusions": ""
    },
    "0402/astro-ph0402405_arXiv.txt": {
        "abstract": "We measure the anisotropy of dark matter flows on small scales ($\\le 500$ kpc) in the near  environment of haloes using a large set  of simulations.  We rely on two different  approaches to quantify the anisotropy  of the cosmic infall: we  measure  the flows  at  the haloes'  virial  radius  while describing  the infalling matter via fluxes through  a spherical shell; we measure the spatial and kinematical  distributions of  satellites and substructures  around haloes detected  by the  subclump finder  ADAPTAHOP first  described in  the appendix \\ref{s:ADAPT}.   The   two  methods  are   found  to  be  in   agreement  both qualitatively and  quantitatively via one  and two points statistics.   \\\\ The peripheral and  advected momentum is correlated  with the spin  of the embeded halo at a  level of $30\\%$ and $50\\%$.  The  infall takes place preferentially in the  plane perpendicular to the  direction defined by the  halo's spin.  We computed  the excess  of equatorial  accretion both  through rings  and  via a harmonic expansion  of the  infall.  \\\\ The  level of anisotropy  of infalling matter  is  found  to be  $\\sim15  \\%$.   The  substructures have  their  spin orthogonal to  their velocity vector  in the halo's  rest frame at a  level of about $5\\%$,  suggestive of  an image of  a flow along  filamentary structures which  provides   an  explanation  for  the  measured   anisotropy.   Using  a `synthetic'  stacked halo,  it is  shown  that the  satellites' positions  and orientations relative to the direction of  the halo's spin are not random even in projection. The average ellipticity  of stacked haloes is $10\\%$, while the alignment excess  in projection reaches $2\\%$.  All  measured correlations are fitted by a simple 3 parameters  model.  \\\\ We conclude that a halo does not see its environment  as an isotropic perturbation, investigate  how the anisotropy is  propagated   inwards  using  perturbation  theory,   and  discuss  briefly implications for weak lensing, warps and the thickness of galactic disks. ",
        "introduction": "Isotropy  is one of  the fundamental  assumptions in  modern cosmology  and is widely verified on  very large scales, both in large  galaxies' surveys and in numerical  simulations.   However   on  scales  of  a  few   Mpc,  the  matter distribution  is structured  in  clusters  and filaments.   The  issue of  the anisotropy down to galactic and cluster scales has long been studied, as it is related to the search for large scale structuration in the near-environment of galaxies.   For   example,  both  observational   study  (e.g.   \\citet{West}, \\citet{Plionis},      \\citet{Kitz})      and     numerical      investigations (e.g.  \\citet{Falten}) showed  that galaxies  tend  to be  aligned with  their neighbours  and support the  vision of  anisotropic mergers  along filamentary structures.  On smaller  scales, simulations of rich clusters  showed that the shape  and  velocity  ellipsoids  of  haloes  tend  to  be  aligned  with  the distribution   of   infalling  satellites   which   is  strongly   anisotropic (\\citet{Tormen}).  However the point is still moot and recent publications did not confirm such an anisotropy  using resimulated haloes; they proposed 20$\\%$ as  a  maximum  for  the  anisotropy  level  of  the  satellites  distribution (\\citet{Vitvitska}).  When considering preferential directions  within the large scale cosmic web, the  picture  that comes  naturally to mind  is  one  involving these  long filamentary structures  linking large  clusters to one  other. The  flow of haloes  within these  filaments  can  be responsible  for  the emergence  of preferential directions and alignments. Previous publications showed that the distributions of spin vectors are   not  random.   For  example,  haloes  in simulations tend to  have their spin pointing orthogonally  to the filaments' direction  (\\citet{Falten}).   Furthermore,  down  to galactic  scales,  the angular  momentum remains  mainly aligned  within  haloes (\\citet{Bullock}). Combined with the results suggesting that haloes' spins are mostly sensitive to recent  infall (\\citet{Haarlem}),  these  alignment properties  fit well with accretion  scenarii  along special  directions~:  angular  momentum can  be considered as a good marker to test this picture.  Most  of  these  previous studies  focused  on  the  fact  that  alignments  and preferential directions are consequences of the formation process of haloes. However, the effects of such preferential directions on the inner properties of  galaxies have  been  less addressed.   It  is widely  accepted that  the properties  of galaxies  partly result  from their  interactions  with their environments.   While the  amplitude  of the  interactions  is an  important parameter, some  issues cannot  be studied without  taking into  account the spatial  extension  of these  interactions.   For  example,  a warp  may  be generated  by   the  torque  imposed   by  infalling  matter  on   the  disk (\\citet{OsBin}, \\citet{Lopez})~: the direction but also the amplitude of the warp  are  a  direct  consequence   of  the  spatial  configuration  of  the perturbation.  Similarly, it is likely  that disks' thickening due to infall is  not   independent  of  the   incoming  direction  of   satellites  (e.g. \\citet{Quinn}, \\citet{Velazquez}, \\citet{Huang}).  Is it possible to observe  the small-scale alignment ? In particular, weak lensing deals with effects as small as the level of detected anisotropy (if  not smaller)  (e.g.   \\citet{Hatton}, \\citet{Croft},  \\citet{Heavens}), hence the  importance to  put quantitative constraints  on the  existence of alignments on small scales. Therefore, the present paper also  addresses the issue of detecting preferential projected orientations on the sky of  substructures  within  haloes.  Our  main  aim  is  to   provide  quantitative  measurements  to  study  the consequences of  the existence of  preferential directions on  the dynamical properties  of  haloes  and  galaxies,  and on  the  observation  of  galaxy alignments.    Hence  our  point   of  view   is  more   galactocentric  (or cluster-centric)  than  previous studies.   We  search  for local  alignment properties on  scales of  a few hundred  kpc.  Using  a large sample  of low resolution  numerical simulations,  we aim  to extract  quantitative results from  a large  number of  halo  environments.  We  reach a  higher level  of statistical significance while reducing the cosmic variance.  We applied two complementary approaches  to study the  anisotropy around haloes:  the first one is particulate and uses a new substructure detection tool ADAPTAHOP, the other  one  is  the  spherical  galactocentric fluid  approach.   Using  two methods, we can assess the self-consistency of our results.  After a brief  description of our set of simulations  ($\\S 2$), we describe the  galactocentric  point of  view  and  study  the properties  of  angular momentum and infall  anisotropy measured at the virial  radius ($\\S 3$).  In ($\\S 4$) we  focus on anisotropy in the  distribution of discrete satellites and substructures and we study the properties of the satellites' proper spin , which provides an explanation for the detected anisotropy.  In ($\\S 5$) we discuss the  level of anisotropy as seen  in projection on the  plane of the sky.  We then  investigate how the anisotropic infall  is propagated inwards and discuss the possible implications of  our results to weak lensing and to the dynamics  of the disk through  warp generation and  disk thickening ($\\S 6$).   Conclusions  and  prospects   follow.   The  appendix  describes  the substructures detection tool ADAPTAHOP  together with the relevant aspects of one    point centered  statistics on  the sphere.   We  also formally derive there the perturbative inward  propagation of infalling fluxes into a collisionless self gravitating sphere. ",
        "conclusions": "\\begin{figure*} \\centering \\resizebox{15cm}{10cm}{\\includegraphics{./sketch5}} \\caption{A schematic representation of all estimate of anisotropic accretion considered in  this paper.   {\\bf 1)}: we  measured the distribution  of the angle  between the  orbital momentum  on the  virial sphere  and  the halo's spin. The average  orbital momentum measured on the  virial sphere is mostly aligned with the  spin of the halo embedded in  the virial sphere (discussed in section \\ref{s:two-points}).  {\\bf 2)}:  we compared the accretion in the plane  orthogonal to  the  direction of  the  halo's spin  with the  average accretion on  the sphere.  On  the virial sphere,  we detected an  excess of ring-like  or  harmonic accretion  in  the  equatorial  plane (discussed  in section \\ref{s:one-point}  and \\ref{s:expansion}). {\\bf  3)}: in projection, we  used  a `synthetic'  halo  to look  at  the  distribution of  satellites detected  with ADAPTAHOP  and at  the orientation  of their  spin  around the direction of the spin.  In  projection, satellites lie preferentially in the projected equatorial plane  (discussed in section \\ref{s:projsatpop}).  {\\bf 4)}: We measured the angle between  the halo's spin and the orbital momentum of  each satellite.  The  orbital momentum  of satellites  is preferentially aligned  with  the  spin  of   their  hosting  core  (discussed  in  section \\ref{s:Corespin_satmomentum}).   {\\bf 5)}:  We compared  the  orientation of each  satellite  velocity  vector  (in  the  core's  rest  frame)  with  the orientation of  their own spin.  The  velocity vector of  satellites (in the core's rest frame)  is orthogonal to the direction  of their spin (discussed in  section  \\ref{s:satvel}).   {\\bf  6)}:  In  the  equatorial  plane,  the projected  orientation  of  satellites  is  more 'radial',  while  near  the direction of  the spin  a 'circular' configuration  of orientation  seems to emerge (discussed in section \\ref{s:projsat}).} \\label{sketch} \\end{figure*}  \\begin{table} \\caption{Summary of the fitting parameters for the angular correlations} \\begin{tabular}{@{}lcccc@{}} \\hline angle & $a_1$ & $a_2$  & $a_3$ & $a_4$\\\\ \\hline $\\theta_{\\rm \\rho L}$ &    2.351$\\pm$0.006     &    -0.178$\\pm$0.002    &     1.343$\\pm$0.002    & 0.669$\\pm$0.000\\\\\\hline   $\\theta_{\\rm  \\rho   vrL}$  &   3.370$\\pm$0.099  & -0.884$\\pm$0.037 &  1.285$\\pm$0.016 & 0.728$\\pm$0.001\\\\  \\hline $\\theta_{\\rm cs}$   &    0.399$\\pm$0.003   &   0.059$\\pm$0.008    &   0.881$\\pm$0.005   & 0.938$\\pm$0.000\\\\   \\hline    $\\theta_{\\rm   vs}$   &    0.295$\\pm$0.004   & 1.544$\\pm$0.001    &     0.804$\\pm$0.005    &    0.914$\\pm$0.001\\\\    \\hline $\\theta_{p}$&0.099 $\\pm$  0.003 &  1.548 $\\pm$ 0.003  & 0.825 $\\pm$  0.013 & 0.973 $\\pm$ 0.000\\\\ \\hline \\end{tabular}  Here $\\theta_{\\rm  \\rho L}$  is the  angle between the  halo's spin  and the angular momentum measured  on the virial sphere; $\\theta_{\\rm  \\rho vrL}$ is the angle between the halo's spin and the accreted angular momentum measured on the virial sphere; $\\theta_{\\rm cs}$ is the angle between the core's spin and the satellite orbital momentum;   $\\theta_{\\rm vs}$ is the angle between the satellite velocity  in the core 's rest frame  and the satellite's spin; $\\theta_p$  is the  projected angle  between  the satellite's  spin and  its direction relative  to the core s'  position.  The fitting model  we used is $1+\\xi(\\theta)=a_1/(\\sqrt{2\\pi}a_3)\\exp\\left[-(\\theta-a_2)^2/(2a_3^2)\\right]+a_4$. \\end{table}  \\begin{table} \\caption{Summary of other quantities related to anisotropic accretion} \\begin{tabular}{@{}lc@{}} \\hline $\\langle  \\delta_m\\rangle(z)$ & 0.0161($\\pm$ 0.0103)  z +0.147 ($\\pm$ 0.005)\\\\  \\hline $S_3(\\delta_m)$  &  0.44\\\\ \\hline  $\\epsilon(R_m)$ &  0.1\\\\ \\hline ${\\tilde  a_{2 0}}$  & -0.65 $\\pm$  0.04\\\\ \\hline${\\tilde  a_{40}}$ & 0.12  $\\pm$  0.02\\\\  \\hline  ${\\tilde  a_{6  0}}$  &  -0.054  $\\pm$  0.015\\\\ \\hline${\\tilde  a_{80}}$ &  0.0145 $\\pm$  0.0014\\\\ \\hline \\end{tabular} \\medskip  $\\langle \\delta_m\\rangle(z)$ is the redshift evolution of the average excess of   accretion  in   the  plane   orthogonal   to  the   direction  of   the spin.  $S_3(\\delta_m)$ is  the skewness  of  the distribution  of excess  of accretion.  $\\epsilon(R_m)$ is  the axis  ratio  $a/b-1$ with  $a>b$ of  the projected  satellite distribution.  ${\\tilde a_{2  0}}$,  ${\\tilde a_{40}}$, ${\\tilde  a_{6 0}}$  and  ${\\tilde  a_{8 0}}$  are  the normalized  harmonic coefficients of the `equatorial' modes. \\end{table}   \\subsection{Conclusion}  Using a set of 500 $\\Lambda$CDM simulations, we investigated the properties of the spatial configuration of the  cosmic infall of dark matter around galactic $\\approx L_{\\star}$  haloes.  The aim of the  present work was to  find out if the existence  of preferential  directions existing on  large scales  (such as filaments) is reflected in the behaviour of matter accreted by haloes, and the answer is a clear quantitative yes.  Two important assumptions  were made in the present  paper.We did not consider different class of haloes' masses  (except for Fig.  \\ref{Imass}), but instead applied normalisations  to includes all haloes in  our measurements (considering e.g. the statistical average of constrasts).   We also did not take in account outflows and focused on accreted quantities.  First  we looked  at the  angular distribution  of matter  at  the interface between the  intergalactic medium and the  inner regions of  the haloes.  We measured the accreted  mass and the accreted angular  momentum at the virial radius, describing these quantities as spherical fields. \\begin{itemize} \\item The  total (resp.  advected)  angular momentum measured at  the virial radius is strongly aligned with the inner spin of the halo with a proportion of aligned  configuration $30\\%$ (resp.  $50\\%$) more frequent than  the one expected  in   an  isotropic  distribution  of   accreted  angular  momentum ($1+\\xi_\\mathrm{LS}(0)  \\sim  1.5$).   This  result reflects  the  importance  of accreted angular momentum in the building of the haloes' inner spin. \\item The accretion of mass measured at  the virial radius in the ring-like region perpendicular to the direction of the halo's spin  is $\\sim 15 \\%$ larger than the one expected in the case of an isotropic infall of matter.\\\\ We also detected the excess of  accretion at the same level in the equatorial  plane using a spherical  harmonic  expansion  of  the  mass density  flux. \\item  In  the  spin's frame,  the  average  of  the harmonic  $a_{\\ell  0}$ coefficients  does not  converge toward  zero,  indicating that  there is  a systematic accretion structured in rings parallel to the equatorial plane.  Using  the substructure  detection  code ADAPTAHOP,  we  confirmed that  the existence of a preferential plane for the infalling mass is reflected in the distribution of satellites around  haloes. \\item Investigating the degree of alignment between the orbital momentum of satellites and the central spin of the halo, it is shown that the aligned configuration is present in excess of $\\sim  12\\%$. Satellites  tend to  revolve in  the plane  orthogonal  to the direction of the halo's spin.   The two methods (using spherical fields and satellites' detection) yield consistent results and suggest  that  the  image  of   a  spherical  infall  on  haloes  should  be reconsidered at the  quoted level. \\\\ We studied the distribution of the angle between the direction of accretion of satellites and their own spin. \\item An  orthogonal configuration is $5  \\%$ more frequent  than what would one  expect  for  an  isotropic  distribution  of  spin  and  directions  of accretion.  Satellites  tend to be  accreted in the direction  orthogonal to their own spin.  These findings  are interpreted as the  results of the  filamentary flows of structures,  where  satellites  and  haloes  are  accreted  along  the  main direction  of filaments  with their  spins orthogonal  to  this preferential direction. The flow along filaments also explains why the matter is accreted preferentially in the equatorial plane at the virial radius. The halo points its spin  perpendicular to the  flow and sees  a larger flux in  the regions normal to the flow direction, i.e.  near the equator.  Thus, it appears that the existence of preferential directions on large scales is still relevant on galactic  scales and  should have  consequences for  the inner dynamics of the halo.  We  addressed the issue  of  observing  these  alignments in projection. \\item The distribution  of satellites projected onto the  sky is flattened, with an axis ratio of $1.1$ at the virial radius. \\item  It seems  that the orientation of satellites  around their haloes is not random, even  if the  two dimensional representation  dilutes the  effects of alignments.  The `radial' orientation, where the  satellites main axis is aligned with  the line joining  the satellite to  the halo centre,  is $\\sim 5\\%$ more frequent  than the one expected in  a completely random distribution of  orientation.  The  `circular' configuration,  where the  satellites main axis is perpendicular to the line  joining the satellite to the halo centre, is also present  in excess compared to an random  distribution near the pole of  the  host  galaxy. \\end{itemize} All  corresponding  fits  are  summarized  in  Table~1  and~2,  while  Fig. \\ref{sketch} gives a schematic view of the measurements we carried out.  We investigated how the self-consistent dynamical response of the halo would propagate anisotropic infall down to  galactic scales. In particular we gave the corresponding polarisation operator in  the context of an opened system. We  have  shown in  appendix~\\ref{s:response} that  accounting  for  dark matter  infall required the knowledge  of the first three moments of  the flux densities, $ \\varpi_{\\rho}(\\bO,\\tau),\\varpi_{\\rho \\bfv }(\\bO,\\tau), \\varpi_{\\rho \\sigma_i \\sigma_j }(\\bO,\\tau) $.  It is suggested  that the existence of a preferential  plane of accretion of matter,  and thus  of angular  momentum, should  have an  influence  on warp generation and disk thickening.   If the anisotropic properties of infalling matter measured  in the  outer parts  of haloes are  conserved in  the inner region  of galaxies, there  may exist  a 'typical'  warp amplitude  and this anisotropic accretion of  matter may explain the properties  of warp line of nodes.  In the same spirit, the efficiency of the thickening of the disk may be enhanced  or reduced  by equatorial accretion.   Finally, our  finding of intrinsic alignments  on small  scales as well  as specific  orientations of structures should be  relevant for cosmic shear studies  on wide and shallow surveys.  \\subsection{Prospects} \\label{s:prospect}  The  main   purpose  of  our  investigation  was   to  provide  quantitative measurements of  the level  of anisotropy involved  in the infall  on scales $\\leq  500$  kpc.   The  next   step  should  clearly  involve  working  out quantitatively  their implications  for warp,  disk heating  etc...   as was discussed in section~\\ref{s:applications}.  Our measurements were carried out at $R_{200}$, which on galactic scales is a long way from the inner region  of the galaxy.  One should clearly propagate the infall (and  its anisotropy) towards the centre of  the galaxy, and more radial infalling  components will play a  more important role  and should be weighted accordingly.  It should also be  stressed that we did not take into account the extra polarisation induced  by the presence of an embedded disk, which will undoubtedly reinforce the  polarisation and the anisotropy of the infall.  We also concentrated on  mass accretion, as the lowest order moment of the underlying ``fluid''  dynamics.  Clearly higher moments involving the anisotropically accreted momentum, the  kinetic energy etc.  are dynamically relevant  for  the  evolution of  the  central  object  as is  discussed  in section~6 and in the appendix. The time evolution of the statistics of these flux  densities is  also essential  for the  inner dynamic  of the  halo and should  be addressed  systematically  as  well.  It  will  be worthwhile  to explore  different  cosmologies  and   their  implications  on  small  scale dynamics, and  on the  characteristics of infalling  clumps, though  we hope that the qualitative results sketched here should persist.  It should be emphasized that some  aspects of the present work are exploratory only,  in that  the  resolution achieved  ($M_{\\mathrm{halo}} >  5\\cdot10^{12} M_\\odot$) is  somewhat high  for $L_{\\star}$ galaxies.   In fact, it  would be interesting  to  see  if the  properties  of  infall  changes for  lower  mass ($M_{\\mathrm{halo}}  <5\\cdot10^{10}  M_\\odot$)  together  with  the  intrinsic properties of galaxies. In addition a systematic study of biases induced  by  the  estimators  of  angular correlations  should  be  conducted, e.g. the mass weighted errors we introduced in section~\\ref{s:two-points}.  Observationaly, the synthetic  halo described in section \\ref{s:projsatpop} could be  compared to stacked  satellite distributions relying  on galactic surveys such as  the SDSS.  Once the anisotropy has  been propagated to the inner  regions  of the  galactic  halo  following  the method  sketched  in section~\\ref{s:applications},  we should  be  in a  position  to compile  a synthetic edge-on  galactic disk and compare  the flaring of  the disk with the  corresponding  predictions.   The  residual  preferred  orientation  of galactic    disks    around    more    massive   objects    discussed    in section~\\ref{s:projsat}  should be observed  on the  scales $\\leq  500 {\\rm kpc}$.  Using larger simulations  will allow us to combine  high resolution with the statistics required to detect the  anisotropic accretion of mass and angular momentum. A  wide range of halo  masses will become accessible  and the halo mass dependency of our findings will be constrained without suffering from the lack of statistics. Better angle determinations will naturally follow from a better resolution and will improve the accuracy of our quantitative results. Resimulations  (zoom simulations)  should give  access to  a larger  range of satellite  masses,  while we  were  here  mostly  sensitive to  the  biggest substructures.  Large  infalling objects are likely to  feel differently the effects  of tidal  forces  or dynamical  friction  than smaller  satellites. Resimulated haloes  allow us  to investigate the  dependency on  the spatial distribution  of  satellites with  their  masses  corresponding  to a  given cosmological environment. However using only  a few resimulations may not be sufficient to overcome cosmic variance and, given the difficulty to produce a large number of high resolution haloes, such a project remains challenging.    The inclusion  of gas  physics in  these simulations and  their impact  on the results is the natural following step. For example, gas filaments are known to be narrower than dark matter filaments, thus we  expect to  see a higher  level of  anisotropy in the  distribution of accreted gas by the haloes.  Furthermore, the transmission of angular momentum from one parcel  of gas to another  (or to the underlying dark  matter) may be highly effective and would lead to higher homogeneity of the properties of the accreted   angular  momentum   direction,   enhancing  the   effect  of   spin alignments. The loss of angular momentum from the gaz to the halo will lead to a modification  of our pure dark  matter findings.  Yet, the  inclusion of gas physics  in  simulations  would  force  us  to  address  issues  such  as  the over-cooling,  the requirement  to take  star formation  and  related feedback processes into account.  It remains that  in the longer term, the inclusion of gas physics cannot be avoided and  will give new insights into the anisotropic accretion of matter by haloes.   \\parn{\\bf Acknowledgements}  {\\sl We  are grateful to J.~Devriendt,  J.~Heyvaerts, A.~Kalnajs, D.~Pogosyan, E.~Scannapieco, F.~Stoehr, R.~Teyssier, E.~Thi\\'ebaut, for useful comments and helpful  suggestions.  DA  would  like  to thank  C.~Boily  for reading  early versions of  this paper. CP would  like to thank  F.~Bernardeau for stimulating discussions during the premises of this work.  We would like to thank D.~Munro for freely distributing his  Yorick programming language (available at {\\em\\tt ftp://ftp-icf.llnl.gov:/pub/Yorick}),   together   with   its   {\\em\\tt   MPI} interface, which we used to implement our algorithm in parallel.}"
    },
    "0402/hep-ph0402240_arXiv.txt": {
        "abstract": "{ If the axino is the lightest superpartner and satisfies cosmological bounds, including a preferred range of the relic abundance of cold dark matter, then the usual stringent constraints on the parameter space of the CMSSM become greatly relaxed. The lightest superpartner of the usual CMSSM spectrum will appear to be stable in collider experiments but will not necessarily obey relic abundance constraints. It may be either neutral (lightest neutralino) or charged (typically a stau). With the axino as cold dark matter, large regions of the CMSSM, often corresponding to heavy superpartners, become allowed, depending on the axino mass and the reheating temperature. }  \\preprint{CERN--PH--TH/2004--023\\\\ hep-ph/0402240 } ",
        "introduction": "\\label{sect:intro}  The lightest supersymmetric particle (LSP) that arises in models featuring supersymmetry (SUSY) at low energy with the conservation of $R$--parity can be a cold dark matter (DM) candidate~\\cite{MSSM-dm}. The requirement that the LSP is neutral and produced in the required abundance in the early Universe sets strong constraints on the minimal supersymmetric standard model (MSSM)~\\cite{MSSM-dm} or its constrained version (CMSSM)~\\cite{CMSSM-dm}. In the CMSSM, the LSP is often the lightest neutralino with a large bino component~\\cite{na92,rr93,kkrw94}. In particular, the cases of a neutralino nearly degenerate with the lighter stau~\\cite{efo98} or stop~\\cite{eos03} are often of cosmological importance, but parameter choices for which such charged particles themselves are the LSP are thought to be excluded on astrophysical grounds.  In fact, very strong bounds on the presence of electrically charged and colored relics have been obtained using the unsuccessful searches for exotic nuclei: strong bounds on electrically charged relics have been obtained up to masses of the order $10^8$ GeV, while masses greater than a few $\\tev$ for colored thermal relics remain consistent with observation~\\cite{ch-lsp}.\\footnote{ Note that in any case for stable charged particles not excluded by searches the overclosure bound applies, thus strongly limiting the possibility of such heavy thermal relics, except if some mechanism for suppressing its number density is at work. }  However, if the LSP does not belong to the usual MSSM spectrum, other possibilities arise. One such attractive scenario involves supersymmetric models implementing the Peccei--Quinn mechanism for solving the strong CP problem~\\cite{pq}.  In this class of models the axino, the fermionic superpartner of the axion, can naturally be the LSP and constitute the dominant component of DM in the form of warm~\\cite{kmn,bgm,rtw,ckkr,ay00} or cold dark matter (CDM)~\\cite{ckr,ckkr,crs1}. (See also~\\cite{ckl00,kk02}.) This is because, unlike for the neutralino, the mass of the axino is not directly determined by the soft SUSY--breaking terms and can be much smaller~\\cite{axinomass}.  Axinos can be efficiently produced in the early Universe through several possible processes. A class of {\\em thermal production} (TP) processes involves scatterings and decays of particles in the primordial plasma. Alternatively, in {\\em non--thermal production} (NTP) the next--to--lightest supersymmetric particle (NLSP) decays after first freezing out from the plasma. In addition, one can think of other possible production mechanisms, \\eg\\ from inflaton decay, but they are much more model dependent and not necessarily as efficient. The above processes are supposed to re--generate the axinos, after their primordial population has been diluted as a result of inflation and subsequent reheating with $\\treh\\lsim\\fa$ (with $\\fa\\sim10^{11}\\gev$ being the Peccei--Quinn scale) in order to avoid overclosure. Otherwise, the axinos would have to be very light, $\\maxino<0.2\\kev$~\\cite{rtw}, with an update in~\\cite{ckkr}, thus constituting warm DM.  In~\\cite{ckr,ckkr,crs1} we conducted a detailed investigation of the axino LSP as CDM in the KSVZ--type (hadronic) axion models~\\cite{ksvz} coupled with the MSSM. As summarized in Figs.~8 and~9 of~\\cite{crs1}, we showed that, for $\\maxino\\gsim 100\\kev$ and $\\treh\\lsim 5\\times 10^{6}\\gev$, the axino can constitute {\\em cold} DM, while at a lower mass range and larger $\\treh$ it could be a warm or even hot DM relic. At $\\treh\\gsim 10^{4}\\gev$ NTP is typically subdominant while TP processes allows for a narrow region of $\\maxino$ to satisfy the relic abundance constraint $\\abunda\\sim1$. On the other hand, in the region of $\\maxino\\gsim 10\\mev$ and $\\treh\\lsim 10^{4}\\gev$, the NTP mechanism of neutralino (or other NLSP) decay after freeze-out often plays a dominant role in producing enough axinos, while at larger $\\treh$ TP does this too efficiently. Contrary to the case of a gravitino, the other popular alternative SUSY candidate for the LSP and CDM, the lifetime of the NLSP is much shorter than $10^4 \\sec$ and therefore avoids the very strong bound coming from photo--destruction of light elements after nucleosynthesis. For this reason, only the much weaker constraints from hadronic showers play a role for the axino LSP scenario. However, if the squarks are much lighter than the gluino, their decays in TP processes can reduce an upper bound on $\\treh$ further~\\cite{crs1}.  In this paper, we turn our attention to investigating the implications on the allowed spectra of the CMSSM of assuming the axino as the LSP and CDM, with the right amount of the cosmological relic abundance and with other constraints, especially from nucleosynthesis, satisfied. This is of interest because the standard paradigm is highly constrained, especially at not too large values of $\\tanb$ on which we will concentrate here. While this makes the CMSSM highly predictive, a question arises as to whether viable alternative scenarios exist which would lead to very different predictions and to possibly relaxing the stringent cosmological bounds on the CMSSM without abandoning a supersymmetric explanation for CDM.  As we will see, very different regions of the CMSSM parameter space will often become allowed.  In particular, it will be possible to have either an electrically neutral NLSP (the lightest neutralino) or a charged one (the lighter stau). The NLSP will appear in a collider detector as a stable lightest ordinary supersymmetric partner (LOSP). Since axino couplings are suppressed by the Peccei--Quinn scale, the NLSP lifetime is long compared to the timescales of relevance for accelerator searches. Therefore, the usual mass bounds arising from the decay of the sparticle into the lightest neutralino are not applicable. Depending of the case of a chargino or slepton L(O)SP, the bounds become a bit relaxed or more stringent. For the case of a ``stable\" charged L(O)SP, a general lower bound of about $99.4 \\gev$ has been obtained at LEP~\\cite{lep-ch-lsp}.  This may, once superpartners are discovered at the LHC, allow one to distinguish the case of the axino from the usually assumed case of the lightest neutralino as the true LSP.  In this context it might be worth commenting on the gravitino as the LSP and CDM. In a recent study~\\cite{eoss04-gravitino} it has been concluded that one can find limited regions of the CMSSM where the gravitino LSP would be allowed but that, in these regions, NLSP decays typically provide too little relic abundance for the gravitino to fall into the favored range, if one does not want to spoil nucleosynthesis predictions.  However, thermal production processes involving gluino decay and scattering may rectify this \\cite{thermal-grav}.  In the following, we will first summarize the effective axino--neutralino couplings and compute effective axino--stau couplings and corresponding NLSP lifetimes in the MSSM coupled with the KSVZ axion models. We will then discuss bounds on the NLSP number density from primordial nucleosynthesis constraints and summarize the calculation of axino relic abundance. We will next concentrate on the CMSSM and will first illustrate the cosmologically allowed parameter space in the CMSSM in the standard scenario with the neutralino LSP. Next we will re--compute the relic abundance in the CMSSM with the axino as the LSP and will investigate the resulting changes in the cosmological bounds on the ($\\mhalf,\\mzero$) plane for several values of the axino mass and $\\treh$. ",
        "conclusions": "The Constrained MSSM provides a popular and predictive framework for analyzing properties of low--energy SUSY. While theoretical assumptions and experimental bounds from LEP and $b\\to s\\gamma$ provide significant constraints on the parameter space, it is the cosmological relic density of the LSP as cold DM and the requirement of its electric neutrality that rule out most of the ($\\mhalf,\\mzero$) plane and, for $\\tanb\\lsim50$, allow for only a tiny region in the plane.  In this paper, we have considered the cosmological bounds on the scenario in which the LSP is not the usual neutralino but instead the axino. Our purpose was not to provide an exhaustive scan of the CMSSM but to illustrate with a few examples the very different patterns that arise.  We concentrated on two cases of a fairly light axino $\\maxino=1\\gev$ and of $\\maxino\\sim m_{\\rm NLSP}$ where the NLSP was either the neutralino or the lighter stau. In computing the axino relic abundance we included both thermal and non--thermal production processes.  Our results are a function of $\\treh$ which in all cases had to be rather low, below a few hundred $\\gev$ in order to give the right amount of axino CDM. We further included constraints from nucleosynthesis which are important for lighter axinos and staus.  The cosmologically favored regions consistent with axino CDM relic abundance are in general very different from the usual scenario. Depending on $\\maxino$ and/or $\\treh$, basically nearly any point in the ($\\mhalf,\\mzero$) plane can become cosmologically allowed. This is true for both the regions where the NLSP is the neutralino or the lighter stau.  This significantly relaxes the cosmological bounds obtained in the absence of the axino.   From the point of view of collider phenomenology, with the lifetime of ${\\cal O}(1\\sec)$, the NLSP would appear stable in a detector. However, if the NLSP is the neutralino, then very likely one would find its relic abundance (calculated under the assumption that it is the true LSP) to significantly exceed $0.13$. The NLSP could also be electrically charged (stau), or even colored (stop) -- the possibility we have not investigated here. In conclusion, in collider studies of the CMSSM parameter space, both the apparently cosmologically excluded bulk of the $(\\mhalf,\\mzero)$ plane where the neutralino is the NLSP and the wedge at its bottom where the lighter stau is the NLSP, should be given equal attention as the narrow regions preferred by the standard paradigm."
    },
    "0402/nucl-th0402077_arXiv.txt": {
        "abstract": "The inelastic neutral reaction of neutrino on $^{4}\\mathrm{He}$ is calculated microscopically, including full final state interaction among the four nucleons. The calculation is performed using the Lorentz integral transform (LIT) method and the hyperspherical-harmonic effective interaction approach (EIHH), with a realistic nucleon-nucleon interaction. A detailed energy dependent calculation is given in the impulse approximation. With respect to previous calculations, this work predicts an increased reaction cross-section by $10\\%-30\\%$ for neutrino temperature up to $15$ MeV. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0402/astro-ph0402627_arXiv.txt": {
        "abstract": "In this work we aim at addressing the still open question about the nature of Narrow-Line Seyfert 1 (NLS1) galaxies: are they really active nuclei with lower mass Black-Holes (BHs) than Seyfert 1 (S1) and quasars? Our approach is based on the recently discovered physical connections between nuclear supermassive BHs and their hosting spheroids (spiral bulges or ellipticals). In particular we compare BH masses of NLS1s and S1s analyzing the properties of their hosts by means of spectroscopic and photometric data in the optical wavelength domain. We find that NLS1s fill the low BH mass and bulge luminosity values of the $\\rm {\\cal M}_{BH}-M_{B}$ relation, a result strongly suggesting that NLS1s are active nuclei where less massive BHs are hosted by less massive bulges. The correlation is good with a relatively small scatter fitting simultaneously NLS1s, S1s and quasars. On the other hand, NLS1s seem to share the same stellar velocity dispersion range of S1s in the $\\rm {\\cal M}_{BH}-\\sigma_*$ relation, indicating that NLS1s have a smaller BH/bulge mass ratio than S1s. These two conflicting results support in any case the idea that NLS1s could be young S1s. Finally we do not confirm the significantly non linear BH--bulge relation claimed by some authors. ",
        "introduction": "Narrow-Line Seyfert 1 (NLS1) galaxies belong to the family of active galactic nuclei (AGNs) and owe their name to the peculiar properties that distinguish them from the other type 1 AGNs. In general a Seyfert 1 (S1) is classified as NLS1 when its nuclear spectrum shows the permitted lines only slightly broader (FWHM(H$\\beta) < 2000$ km s$^{-1}$) than the forbidden ones \\citep{op85}. But NLS1 galaxies are characterized also by a strong Fe II emission, a soft X-ray slope steeper than the slope typical for S1 galaxies, rapid and large soft X-ray variability \\citep*{bolet96}, a weak Big Blue Bump in the optical/UV range, likely shifted toward higher energies \\citep{pou87}, a bright IR emission and a nuclear super-solar metallicity \\citep{mathur00,km01}.  Even though NLS1s are known since almost twenty years \\citep{op85}, their nature is still a matter of debate. For example, at present the reason of the narrowness of broad permitted lines in NLS1s is not clear. \\citet{bolet96} suggested that if the gravitational force from the central Black-Hole (BH) is the main cause of motion of the Broad Line Region (BLR) clouds, narrower optical emission lines will result from smaller $\\cal M_{\\rm BH}$ ($\\sim 10^6-10^7 {\\cal M}_{\\odot}$) provided the BLR distance from the central source does not change strongly with $\\cal M_{\\rm BH}$. These BHs with smaller masses are expected to accrete matter at near or super-Eddington rates in order to maintain the relatively normal observed luminosities. Recently \\citet{matet01} proposed that NLS1s may be relatively young AGNs hosting BHs still in a growing phase. Nevertheless other authors suggested that if the BLR clouds were largely confined to a plane, as it seems to happen in radio-loud AGNs \\citep[see e.g.][and references therein]{mcdun02}, NLS1 galaxies could be simply a case in which their BLR is observed more face-on than in S1s \\citep[see e.g.][]{op85}. \\citet{smi02} pointed out that also a partly obscured BLR could justify narrow permitted lines, but these last two scenarios should produce polarization. They claimed that \\ha\\, polarization properties of the NLS1s are indistinguishable from those of S1s, excluding obscuration of the inner regions of BLR or a face-on orientation of a disk-like BLR as the explanation for the relatively narrow broad-line profiles.   The connection between physical properties of nuclear supermassive BHs and their host galaxies, on which several works have focused in the last years, might turn out to be a powerful tool to understand the nature of NLS1s and settle the above cited controversies. \\citet{koric95} first found out a correlation between BH mass ($\\cal M_{\\rm BH}$) and the absolute B magnitude of the spheroidal component ($M_{B}$). \\citet{maget98} determined  $\\cal M_{\\rm BH}$ values for a sample of 32 nearby galaxies and suggested that $\\cal M_{\\rm BH}$ is proportional to the $\\cal M_{\\rm bulge}$ such that on average $\\cal M_{\\rm BH}/\\cal M_{\\rm bulge}\\sim$ 0.005. Other authors estimated a similar mass ratio, $\\sim$ 0.002 (see e.g.,  Ho 1999). Recently, studying a sample of nearby galaxies \\citet{gebet00a}, \\citet{fermer00} and \\citet{tre02}  have shown that $\\cal M_{\\rm BH}$ is tightly correlated with the velocity dispersion of the bulge stellar component ($\\sigma_*$), although they disagreed about the value of the slope.  In AGNs as well, BHs are also expected to correlate with their host bulges. This possibility was explored on a sample of PG quasars by \\citet{lao98} who found agreement with the relation of \\citet{maget98}. Later \\citet{wan99} claimed that Seyfert galaxies show on average a $\\cal M_{\\rm BH}/\\cal M_{\\rm bulge}$ ratio systematically lower than that for normal galaxies and quasars. Conversely \\citet{gebet00b} included in their work seven AGNs for which the $\\cal M_{\\rm BH}$ were obtained by means of the reverberation mapping technique. They found that these objects were in agreement with their previously found $\\cal M_{\\rm BH}$--$\\sigma_*$ correlation. Further support came from \\citet{mcdun01}, and from \\citet{Wuhan01} who studied samples of quasars and Seyfert galaxies, and did not find any evidence that Seyfert galaxies follow a different $\\cal M_{\\rm BH}- \\cal M_{\\rm bulge}$ relation from quasars or nearby galaxies.\\\\ Until now no agreement has been found about NLS1s. \\citet*{matet01} and very recently \\citet{bz03} and \\citet{gm04} showed that the $\\cal M_{\\rm BH}/\\cal M_{\\rm bulge}$ ratio in NLS1s is significantly smaller than that for Seyfert galaxies. Conversely \\citet{walu01}, studying a sample of 59 NLS1s observed spectroscopically by \\citet*{vvg01}, found that there is non clear difference in the $\\cal M_{\\rm BH}$--$\\sigma_*$ relation (where $\\sigma_*$ is represented by the \\oiii\\ emission line width) between NLS1s, Broad-Line AGNs and nearby galaxies.  Our purpose in this work is to investigate the nature of NLS1s by exploring the physical properties of their bulges, namely the luminosity and the nuclear stellar velocity dispersion. Our approach makes use of both new observational data and data from the literature. The necessary corrections are applied to the latter in order to obtain as homogeneous as possible a dataset. Furthermore, as a main difference with respect to several other works on this topic, when determining the bulge properties of the host galaxies we take into account the influence of the AGN, which, as we show, can be non-negligible and affect the results considerably. The structure of the paper is as follows: in \\S\\,2 we present our dataset and derive and compare BH mass values for a sample of NLS1s and S1s. In \\S\\,3 we estimate the blue absolute magnitudes and the stellar velocity dispersions of their bulges. Our results are summarized and discussed in \\S\\,4. ",
        "conclusions": "In this work we have investigated the nature of NLS1s by following an indirect way, that is by using the host galaxy properties to compare black hole masses in NLS1s, S1s and also quasars.  Starting from the assumption that the emission line clouds of BLR are gravitationally bound to the BH and in random motion, we have calculated BH masses of NLS1s obtaining a typical logarithmic value of 6.65$\\pm$0.64 M$_\\odot$, which is almost one order of magnitude lower than the value obtained for S1s (7.37$\\pm$0.62 M$_\\odot$). Simultaneously we have confirmed that NLS1s and S1s have quite similar nuclear luminosities: (in logarithm), 43.51$\\pm$0.65 L$_\\odot$ and 43.08$\\pm$0.45 L$_\\odot$ respectively.  The physical properties of the bulges were investigated in NLS1s and S1s by means of photometric and spectroscopic data. \\\\ Published total apparent B magnitudes were corrected for extinction, inclination, redshift and AGN contribution, and converted into absolute magnitudes. Then, the bulge magnitudes were calculated taking advantage of the empirical B/T -- morphology relation given by \\citet{side86}, and used by several authors. We plotted these values against BH masses for NLS1s and S1s, adding quasars extracted from literature, and we found that NLS1s are mostly confined in the lower ranges of the ${\\cal M}_{BH}- M_{B}$ plane. This result suggests that NLS1s are characterized by less massive BHs hosted in less massive bulges than S1s. This is in agreement with previous findings of \\citet{walu01}, who explored only the ${\\cal M}_{\\rm BH}-\\sigma_*$ relation, and in contrast with \\citet{matet01} and \\citet{bz03}, who claimed for NLS1s with lower ${\\cal M}_{\\rm BH}/{\\cal M}_{\\rm bulge}$ ratios. It is not straightforward to justify the different results found by these authors. After a careful inspection of the effects introduced by each correction we applied to the photometric data of our targets, we can observe that the slightly different way we calculated the correction terms for AGN contribution and galaxy inclination seem to cause changes which roughly compensate each other. Therefore we guess that the main source of disagreement is the morphological classification, which strongly affects the resulting bulge magnitudes by quantities in the range 0.25-0.5 dex for uncertainties discussed in \\S\\, 3.1.   A fit of NLS1s, S1s and quasar values together shows a strongly correlated relation, $\\rm M_{B}~=-2.32(\\pm0.18)~log({\\cal M}_{BH})-3.40(\\pm1.36)$, which leads to $\\rm {\\cal M}_{BH}\\propto {\\cal M}_{bulge}^{0.85\\pm0.09}$. Since the slope is close to unity, we do not confirm the results by \\citet{lao01} and \\citet{bz03} who found a significantly non-linear BH--bulge correlation.  The velocity dispersion of the bulge stellar component was also explored to check whether NLS1s have also $\\sigma_*$ values typically lower than those measured in S1s. Since few values of $\\sigma_*$ were available in literature, we had to calculate the others measuring the \\oiii\\ emission line widths, and assuming that the gas kinematics is dominated by the bulge potential. As done before, we plotted these values against BH masses, adding quasars and also nearby non-active galaxies, and obtaining a good correlation ${\\cal M}_{\\rm BH}\\propto\\sigma_*^{3.70\\pm0.37}$. Contrary to the previous result, NLS1s are not clearly separated from S1s in the ${\\cal M}_{\\rm BH}-\\sigma_*$ plane. In particular they span similar ranges of $\\sigma_*$, suggesting that their bulges are in this respect identical. Moreover, all NLS1s of our sample fall below the fit, showing $\\sigma_*$ values higher than expected, and also below the relations found by \\citet{gebet00a}, \\citet{merfer01} and \\citet{tre02}, which on the contrary are well in agreement with our S1s values. This discrepancy between the two results could be caused by the assumption of \\oiii\\ widths as representative of the stellar kinematics. Indeed the \\oiii-$\\sigma_*$ conversion has a large scatter and is not yet proved to be valid for NLS1 galaxies. Therefore we stress the necessity to directly measure $\\sigma_*$, for example observing NLS1s in spectral ranges different form the optical one.  Since both relations are based on important assumptions, which can introduce significant errors, at the moment it is not possible to favor one over the other. The fact that, according to our first result, NLS1 galaxies seem to have less massive bulges harboring equally less massive BHs than S1 galaxies, strongly indicates that the hypothesis of an inclination effect of a disk-like BLR on the narrowness of Balmer emission lines should be rejected, at least for most NLS1s. A pure selection effect is not expected to be related to the physical properties of the bulges, and therefore NLS1s should be just those S1s with intrinsecally smaller bulges to justify the difference we observe between NLS1s and S1s. On the other hand, our second result confirms the smaller $\\rm {\\cal M}_{BH}/{\\cal M}_{bulge}$ ratio in NLS1s, which led \\citet{matet01} and \\citet{gm04} to suggest an evolutionary scenario for these AGNs toward a S1 stage. In particular, NLS1s would be AGNs in a phase when BHs are growing independently from their hosting environment.  However, both cases support the general idea of NLS1s in terms of AGNs with less massive BHs accreting matter at high rates in order to maintain nuclear luminosities comparable to those of S1s, as we said above. Moreover, our results are not in conflict with the evolutionary scenario and sustain the idea that NLS1s are likely to be young S1s, even in case of a joined evolution of BH and bulge. In fact calculations by \\citet{bz03} suggest time scales of the order of some $\\sim10^8$ yr for a NLS1 to become a S1. This time scale is in agreement with the growth time of a spiral bulge through minor merger phenomena. Indeed \\citet{wmh96} demostrated through N-body simulations that a minor merger makes significant disturbances to the morphology of a larger galaxy in less than 1 Gyr of the onset of the merger. Moreover \\citet{agu01} showed that the accretion of small satellites is an effective mechanism for the growth of bulges in spiral galaxies."
    },
    "0402/astro-ph0402282_arXiv.txt": {
        "abstract": "Pulsar motions in our Galaxy from their birth until 2 Gyr are studied statistically via Monte-Carlo simulation of $2\\times 10^5$ pulsars with the best available representation of the Galactic potential.  We find that the distribution of height above the Galactic plane for pulsars with characteristic ages less than about 8~Myr could be well fitted by a Gaussian function. For older pulsars, an extra exponential function is necessary to fit the distribution. The scale-height of the Gaussian component increases linearly with time until about 40~Myr. The height distribution becomes stabilized after about 200~Myr.  These results are not sensitive to initial height or radial distributions. Taking the relationship between the initial velocity and height distribution, we found from the latest pulsar catalog that the height distribution of pulsars younger than 1~Myr directly implies the mean initial velocity of $280\\pm96$~km~s$^{-1}$. Comparison of simulated sample of pulsars with the current available millisecond pulsars shows that their 1D initial velocity dispersion should be most probably $60\\pm10$~km~s$^{-1}$. ",
        "introduction": "Pulsars are high velocity objects in our Galaxy \\cite{ll94}.  They were born in supernova explosions near the Galactic plane where their progenitors live, but move away very fast from the plane.  The evolution of pulsar heights was considered \\cite{bwh+92,hbw+97,mk97} but not explicitly described in most pulsar population syntheses, therefore, the picture for pulsar height evolution was not very clear. We want to demonstrate in this paper by 3D simulations how pulsars move in our Galaxy.  Birth velocities of both normal and millisecond pulsars (MSPs) have been investigated by many authors. For normal pulsars, Lyne \\& Lorimer~\\shortcite{ll94} studied transverse speeds of 29 pulsars younger than 3~Myr and obtained the mean birth velocity $\\bar{V_{\\rm B}}=$450$\\pm$90~km~s$^{-1}$. Taking into account selection effects and using all available pulsar proper motion data in population synthesis, Lorimer, Bailes \\& Harrison~\\shortcite{lbh97}  found the mean birth velocity $\\bar{V_{\\rm B}} \\sim$500~km~s$^{-1}$. Hansen \\& Phinney~\\shortcite{hp97} considered the distribution of proper motions, including all available upper limits, and concluded that the mean birth velocity $\\bar{V_{\\rm B}}$ should be smaller, only 250-300~km~s$^{-1}$. Cordes \\& Chernoff~\\shortcite{cc98} performed a detailed analysis for measured velocities of 49 young pulsars. They favoured a two-component Gaussian model in 3-dimension with characteristic velocities\\footnote{In the Maxwellian distribution, velocity in each dimension is a Gaussian distribution.  The mean velocity in 3-D and velocity dispersion in 1-D are connected by $\\bar{V_{\\rm B}} \\sim \\sigma_v \\sqrt{8/\\pi}\\simeq 1.6 \\sigma_v$.} $\\sigma_v$ of 175~km~s$^{-1}$ and 700~km~s$^{-1}$.  Arzoumanian, Chernoff \\& Cordes~\\shortcite{acc02} infer the velocity distribution of radio pulsars detected in 400-MHz surveys, taking into account beaming, selection effects and luminosity evolution. They also favour a two-component Gaussian model with $\\sigma_v \\sim 90$~km~s$^{-1}$ and 500~km~s$^{-1}$.  For MSPs, Lorimer~\\shortcite{lor95} accounted for known survey selection effects and obtained the local surface density of MSPs, birthrate and the lower limit of the mean birth velocity that is $\\sigma_v > 50$km~s$^{-1}$, i.e. $\\bar{V_B}\\geq80$km~s$^{-1}$.  The results were confirmed by Cordes \\& Chernoff~\\shortcite{cc97}, who obtained the velocity perpendicular to the Galactic plane, $v_z=52^{+17}_{-11}$km~s$^{-1}$, i.e.~a 3-D velocity dispersion $\\sigma_v= 84$~km~s$^{-1}$ using likelihood analysis on previous surveys for MSPs plus selection effects. Later, Lyne et al.~\\shortcite{lml+98} found that the population syntheses of the MSPs with the Maxwellian initial 1D velocity dispersion of $80\\pm20$~km s$^{-1}$ agrees the best with the data of newly discovered MSPs in Parkes southern sky survey together with previous MSPs. The 3D mean birth velocity is correspondingly $130\\pm30$~km~s$^{-1}$.  These results were later confirmed by Toscano et al.~\\shortcite{tsb+99} using more MSP proper motions obtained from timing measurements.  Previously, Narayan \\& Ostriker~\\shortcite{no90} tried to model the observed pulsar populations. They deduced an analytic formula for the height distribution by assuming a Gaussian distribution at all times and solving the energy conservation equation in $z$-direction (perpendicular to the Galactic plane). However, the height distribution changes from time to time when pulsars run away from the Galactic plane. The integrated $z$-distribution of pulsars with different ages may not be a Gaussian. We will show by our Monte-Carlo simulation that the scale-height evolution is much more complicated than a Gaussian.  Analytic solutions obtained from the over-simplified form of a Galactic acceleration model can hardly represent the realistic situation, especially over long time with non-linear evolution of pulsar heights. By simulations, we are able to track the height distribution of a pulsar sample, for any form of acceleration.  In our work, the initial velocity dispersion for MSPs is obtained by comparison of the height distribution between simulated and observed pulsar samples. We also tried to derive the initial velocity dispersion of normal pulsars by comparing their scale-heights and characteristic ages. In Section 2 we will describe the simulation procedures together with input parameters and governing equations. Several possibilities of parameter choice are also discussed. In Section 3, we show the simulation results, mostly regarding how the pulsar scale-height evolves. We demonstrate in Section 4 two applications of our simulations, i.e.~to determine the birth velocity dispersions of normal pulsars and MSPs. The influence of the selection effects is also discussed. The conclusions are given in Section 5. ",
        "conclusions": "We presented a generalized statistical picture of how pulsars move in our Galaxy. The potential given by Paczy\\'nski~(1990) can be a good representation of mass distribution of our Galaxy, and has a simple analytic formula. We found that the final height distributions are not sensitive to the forms of initial $z$ and $R$ distributions. The height distribution can be well fitted by a Gaussian function within $\\sim$8~Myr, and the scale-height increases linearly with time. After that an extra exponential function is required to fit the height distribution. The height distribution gets stabilized after about 200~Myr.  The height distribution of pulsars younger than 1~Myr implies directly that the mean initial velocity of 280$\\pm$96~km~s$^{-1}$. Comparison of the simulated sample of millisecond pulsars and observed sample suggests the 1D initial velocity dispersion of MSPs to be most probably 60$\\pm$10~km~s$^{-1}$, consistent with estimates given by previous authors."
    },
    "0402/gr-qc0402040_arXiv.txt": {
        "abstract": "Axisymmetric numerical simulations of rotating stellar core collapse to a neutron star are performed in the framework of full general relativity. The so-called Cartoon method, in which the Einstein field equations are solved in the Cartesian coordinates and the axisymmetric condition is imposed around the $y=0$ plane, is adopted. The hydrodynamic equations are solved in the cylindrical coordinates (on the $y=0$ plane in the Cartesian coordinates) using a high-resolution shock-capturing scheme with the maximum grid size $(2500,2500)$. A parametric equation of state is adopted to model collapsing stellar cores and neutron stars following Dimmelmeier et al. It is found that the evolution of central density during the collapse, bounce, and formation of protoneutron stars agree well with those in the work of Dimmelmeier et al. in which an approximate general relativistic formulation is adopted. This indicates that such approximation is appropriate for following axisymmetric stellar core collapses and subsequent formation of protoneutron stars. Gravitational waves are computed using a quadrupole formula. It is found that the waveforms are qualitatively in good agreement with those by Dimmelmeier et al. However, quantitatively, two waveforms do not agree well. Possible reasons for the disagreement are discussed. ",
        "introduction": "Rotating stellar core collapse is among the most promising sources of gravitational waves. To date, there has been no systematic work for computation of the collapse to a neutron star and of emitted gravitational waves in full general relativity (but see \\cite{Siebel}). Gravitational waves associated with the formation of rotating neutron stars have been widely computed in the Newtonian gravity \\cite{Newton,Newton1,Newton2,Newton3,Newton4,Newton45,Newton5,Newton6} or in an approximate general relativistic gravity \\cite{HD} using the so-called conformal flatness approximation (or Isenberg-Wilson-Mathews approximation \\cite{IWM}). As indicated in \\cite{HD}, the general relativistic effects modify the dynamics of the collapse and corresponding gravitational waveforms significantly. This implies that simulation in full general relativity is the best approach for accurate computation of gravitational waves.  During the stellar core collapse to a neutron star, the characteristic radius changes from initial stellar core radius $\\sim 2000$ km to neutron star radius $\\sim 10$ km. Adopting a uniform and fixed grid with a grid spacing as $\\sim 1$ km, the required grid number for the simulation is more than 2000 for one direction. With current computational resources, it is very difficult to take such a huge number of grid points in three-dimensional simulations. If the progenitor of the neutron star is not very rapidly rotating, nonaxisymmetric instabilities will not set in and the collapse will proceed in an axisymmetric manner. By restricting our attention to axisymmetric spacetimes, the grid resolution can be improved significantly for a given computational resource. Thus, as a first step, it is better to perform axisymmetric simulations than to do nonaxisymmetric ones for a well-resolved and convergent computation of the collapse, the bounce, and corresponding gravitational waveforms, focusing only on the moderately rapid rotation case.  In this paper, we study gravitational waves from rotating stellar core collapses to a neutron star assuming the axial symmetry. Dynamics of the collapse is followed by fully general relativistic simulation. Gravitational waves are approximately computed using a quadrupole formula adopted and tested in \\cite{SS}. Necessity of adopting quadrupole formulas arises from the fact that the amplitude of gravitational waves is too small ($< 10^{-5}$ in the local wave zone) to accurately extract the waveforms from raw data sets of metric. Although exact gravitational waveforms cannot be computed, the quadrupole formula is a useful tool for approximate computation of gravitational waves associated with matter motion such as oscillations of neutron stars as indicated in \\cite{SS}.  Recently, gravitational waves from axisymmetric rotating stellar core collapses have been extensively computed in a relativistic manner by Dimmelmeier et al. \\cite{HD}. As mentioned above, they determine the gravitational fields adopting an approximate formulation of the Einstein equation. The approximation is likely to be applicable to a moderately relativistic and stationary spacetime such as that for a rapidly rotating neutron star \\cite{CST96}. However, no one has clarified whether this is the case for dynamical spacetimes. To confirm that their treatment is indeed appropriate, it is necessary to compare their solutions with fully general relativistic ones for a calibration. One of the purposes in this paper is to examine whether the numerical solution for the stellar core collapse computed in \\cite{HD} is a well approximated one for a fully general relativistic solution.  In \\cite{HD}, gravitational waveforms were computed in terms of a quadrupole formula. In general relativity, there is no unique definition of the quadrupole moment, nor is the quadrupole formula, for axisymmetric dynamical spacetimes. Accuracy of gravitational waveforms depends on the choice of the quadrupole formula and the gauge conditions. Thus, to know how accurately the approximate gravitational waveforms can be computed by the chosen quadrupole formula, a calibration is required by comparing the resulting waveforms with those computed from the metric as we did in \\cite{SS}. Unfortunately, such calibration is not possible in \\cite{HD} since they adopted an approximate general relativistic formulation for the gravitational field in which the metric does not carry any information of gravitational waves. Consequently, it is not clear whether the quadrupole formula they adopted can actually yield accurate approximate gravitational waveforms and how large the magnitude of the error is. On the other hand, in the previous paper \\cite{SS}, we did such a calibration for a quadrupole formula which is different from that in \\cite{HD}, and showed it possible to compute gravitational waves from oscillating and rapidly rotating neutron stars of high values of compactness fairly accurately, besides possible systematic errors for the amplitude due to neglecting post-Newtonian corrections. Computing gravitational waveforms by the calibrated quadrupole formula and comparing the results with the previous ones, we estimate how accurate the waveforms computed in \\cite{HD}.   This paper is organized as follows. In Sec. II, our numerical implementations for general relativistic simulation in axial symmetry are briefly reviewed. In Sec. III, the initial condition and computational setting are described. Sec. IV presents the numerical results. Sec. V is devoted to a summary. Throughout this paper, we adopt the geometrical units in which $G=c=1$ where $G$ and $c$ are the gravitational constant and the speed of light, respectively. ",
        "conclusions": "We performed axisymmetric numerical simulations of rotating stellar core collapses to a neutron star in full general relativity, paying particular attention to gravitational waveforms and to comparison of our results with previous results \\cite{HD}. The Einstein field equations are solved in the Cartesian coordinates imposing an axisymmetric condition by the Cartoon method \\cite{alcu}. The hydrodynamic equations are solved in the cylindrical coordinates (with the Cartesian coordinates restricting on the $y=0$ plane) using a high-resolution shock-capturing scheme with the maximum grid size $(2500,2500)$. A parametric equation of state is adopted to model collapsing stellar cores and formed protoneutron stars following Dimmelmeier et al. \\cite{HD}. Gravitational waveforms are computed using a quadrupole formula proposed in \\cite{SS}.  We choose moderately rapidly rotating stars as the initial conditions for which the value of $T/W$ is between 0.005 and $0.01$. Simulations are performed changing three parameters ($\\Gamma_1$, $\\Gamma_2$, and $\\Gamma_{\\rm th}$) which characterize the equation of state. The dynamics of the collapse depends on the three parameters as well as $T/W$ and $\\hat A$ of the initial condition. The dependence of the evolution of the system and gravitational waveforms on these five parameters is studied. The value of $\\Gamma_1$ mainly determines the duration of the infall phase and the coherence of the early phase of the collapse. For the smaller value of $\\Gamma_1$, the infall time becomes shorter and the collapse is accelerated more in the central region. This results in that the core mass at the bounce is smaller and that the magnitude of $\\Phi_c$ (which may be regarded as the depth of the gravitational potential) at the bounce is smaller for the smaller value of $\\Gamma_1$. The amplitude of gravitational waves becomes also smaller for the smaller value of $\\Gamma_1$.  The value of $\\Gamma_2$ determines the equation of state for formed protoneutron stars. Thus, it does not affect the evolution during the infall phase. It determines the final value of the central density of the formed protoneutron star and the gravitational waveforms emitted during the ring-down phase in which the eigen oscillation modes of the protoneutron stars are excited. The value of $\\Gamma_{\\rm th}$ determines the strength of shock waves. We choose this value as 1.35, 1.5, and 5/3 extending the work by Dimmelmeier et al. \\cite{HD}. It is found that for the smaller value of this parameter, the shock heating becomes weaker and the amplitude of gravitational waves smaller.  The values of $T/W$ and $\\hat A$ play a significant role in determining the dynamics of collapse and the corresponding gravitational waveforms in particular in the bounce and ring-down phases. For the rigidly rotating case ($\\hat A \\rightarrow \\infty$), the maximum value of $T/W$ is $\\approx 0.009$ which we choose in this paper. Even in this maximum case, the collapse leads to a neutron star irrespective of the values of $\\Gamma_1$, $\\Gamma_2$, and $\\Gamma_{\\rm th}$. This indicates that for rigidly rotating initial conditions, the neutron stars are formed soon after the collapse, irrespective of the angular velocity of the initial condition, with our choice of the equations of state. For the differentially rotating case with $\\hat A = 1/4$, the collapse does not lead to a neutron star but an oscillating star of subnuclear density is formed for $T/W \\agt 0.01$ since the centrifugal force is strong enough near the rotational axis. As shown in \\cite{HD}, more rapidly rotating initial conditions with $T/W \\gg 0.01$ may be constructed. For such high values of $T/W$, a neutron star will not be formed soon after the collapse.  With a slight change of $\\hat A$ from 0.25 to 0.32 for the initial condition, the angular velocity at the rotational axis is changed by a large factor even if $T/W$ is approximately identical. As a result of this change, the subsequent evolution of the collapse and gravitational waveforms in the bounce and ring-down phases are modified significantly. This implies that the dynamics rotating stellar core collapses and the corresponding gravitational waveforms are sensitive not only to the equation of state but also to the initial angular velocity profile.   Several simulations are performed setting the same initial conditions as those adopted in \\cite{HD}. It is found that the dynamics of the collapse and the bounce for such initial conditions are very similar to those found in \\cite{HD}, in which an approximate general relativistic gravity (the conformal flatness approximation) is assumed. This indicates that such approximate relativistic formulation is appropriate for computing axisymmetric rotating stellar core collapses and the subsequent formation of protoneutron stars. (Note that this is the conclusion for the formation of neutron stars. This may not be the case for the black hole formation.)  Gravitational waveforms are compared with a previous result \\cite{HD}. It is found that the waveforms are qualitatively in good agreement but not quantitatively with those in \\cite{HD}. Either of two plausible elements could explain this disagreement. One is that the grid resolution and computational setting are different between two groups. This could modify the waveforms slightly. However, the main reason seems to be that the adopted quadrupole formulas by two groups are different. As mentioned in the previous section, there is no unique definition of the quadrupole formula for dynamical spacetimes in general relativity. This implies that when one attempts to use a quadrupole formula in a relativistic simulation, one needs to calibrate the formula in advance by performing a fully general relativistic simulation and by comparing the waveforms computed by the quadrupole formula with those computed from the metric. The quadrupole formula adopted in our study has been calibrated in simulations for highly relativistic, highly oscillating, and rapidly rotating neutron stars \\cite{SS}. Thus, we believe that the quadrupole formula adopted in this paper is appropriate and that the numerical results presented here are approximate solutions of high quality.  In this paper, we have focused on the neutron star formation and on the comparison with a previous work \\cite{HD}. If more massive progenitor is chosen as the initial condition, a black hole instead of a neutron star will be formed. Formation of black holes and corresponding gravitational waves have been studied by several groups \\cite{N,SP,shiba2000}. However, the initial conditions in their previous works are not very realistic for modeling a rotating stellar core collapse in nature. Namely, the stellar core collapse to a black hole from a realistic initial condition in fully general relativistic simulation is an unsolved issue. We are currently working in this subject and will present the numerical results in a subsequent paper.  \\vspace{4mm}  \\begin{center} {\\bf Acknowledgments} \\end{center}  We thank Toni Font for discussion and careful reading of this manuscript, and Harald Dimmelmeier for comments. We also thank Harald Dimmelmeier, Toni Font, Jose-Maria Ib\\'a\\~nez, and Eward M\\\"uller for suggesting this work. Numerical computations were performed on the FACOM VPP5000 machine in the data processing center of National Astronomical Observatory of Japan. This work is in part supported by Japanese Monbu-Kagakusho Grants (Nos. 14047207, 15037204, and 15740142)."
    },
    "0402/astro-ph0402231_arXiv.txt": {
        "abstract": "{ The discovery of deep spectral features in the X-ray spectrum of 1E1207.4-5209 focussed the attention of the astronomical community on this radio-quiet NS, making it the most intensively observed INS ever. The harvest of X-ray photons, collected mainly by XMM-Newton, unveiled a number of unique characteristics, raising questions on this source very nature. ",
        "introduction": "Neutron star atmosphere models predicted the presence of absorption features depending on  atmospheric composition, but high quality spectra, collected both by \\cha and by \\xmm, did not yield  evidence for any such feature (see Pavlov et al. 2002a and Becker and Aschenbach 2002 for recent reviews). INS spectra are well fitted by one or more black-body curves with, possibly, a power law contribution at higher energies, but with no absorption or emission features. \\\\ The spectrum of 1E1207.4-5209, on the contrary,  appeared from the start as dominated by two broad absorption features seen, at 0.7 and 1.4 keV, both by \\cha (Sanwal et al, 2002)  and \\xmm~(Mereghetti et al, 2002). To better understand the nature of such features, \\xmm~devoted two orbits to 1E1207.4-5209, for a total observing time of 257,303 sec. In the two MOS EPIC cameras the source yielded  74,600 and 76,700 photons in the energy range 0.2 - 3.5 keV, while the pn camera recorded 208,000 photons, time-tagged to allow for timing studies. Analysis of this long observation, while confirming  the  two phase-dependent absorption lines at 0.7 and 1.4 keV, unveiled a statistically significant third line at $\\sim$2.1 keV, as  well as  a  possible fourth  feature at  2.8 keV.  The nearly perfect 1:2:3:4 ratio  of the line centroids,  as well as  the phase variation, naturally following the pulsar B-field rotation, strongly suggest that such lines are due to cyclotron resonance scattering (Bignami et al. 2003).  A year after the \\xmm~observation, \\ee was deeply scrutinized by Chandra for $\\sim$300 ksec, thus logging an overall effective exposure slightly larger than that devoted to the most popular NSs so far, such as Crab, Vela and RX J1856.5-3754.  The release of a new, and improved, instrument calibration software prompted us to revisit the \\xmm~ data set. Here we shall briefly report on some highlights of our new spectral and temporal analysis (see De Luca et al. 2004 for details). In addition, the ``best'' X-ray positional information is used, in conjunction with a new optical observation, to investigate the source optical behaviour. ",
        "conclusions": "\\ee is a radio-quiet NS at the center of a well known SNR. However, the longer we observe such an object, the harder it becomes to make it fit into a known class of the NS family. The long observations devoted to \\ee both by \\cha and by \\xmm~ have unveiled a number of unique and somewhat contradictory characteristics that, at the moment, defy standard theoretical interpretations.  The first and outmost problem seems to be the interpretation of the period evolution. The usual, simplistic, monotonic fit yields a $\\dot{P}$=(1.4$\\pm$0.3)$\\times$10$^{-14}$ s s$^{-1}$ that, combined with the period value, makes \\ee much too old ($>$ 50 times) for its SNR. Moreover, the value of the star's magnetic field inferred  on the basis of the classical dipole braking formula (B=2.6$\\pm$0.3$\\times10^{12}$ G) is significantly different from the value obtained from the cyclotron absorption lines interpreted both in terms of electrons (B$\\sim8\\times10^{10}$ G) as well as protons (B$\\sim1.6\\times10^{14}$ G).  More creative interpretations of the period evolution (e.g. De Luca et al. 2004, Zavlin et al. 2004) require scenarios such as bynary system or peculiar glitching. However, such types of behaviour  are already tightly constrained by the lack of an optical counterpart.  Are the unique characteristics of \\ee pointing towards a new class of NSs?"
    },
    "0402/astro-ph0402188_arXiv.txt": {
        "abstract": "{ We present the results of an extensive search for periodic and irregular variable pre-main sequence (PMS) stars in the young (2--4 Myr) open cluster \\object{NGC\\,2264}, based on photometric monitoring using the Wide Field Imager (WFI) on the 2.2\\,m telescope on La Silla (Chile). In total, about 10600 stars with $I_{\\mathrm{C}}$ magnitudes between 9.8\\,mag and 21\\,mag have been monitored in our $34\\arcmin \\times 33\\arcmin$ field. Time series data were obtained in the $I_{\\mathrm{C}}$ band in 44 nights between Dec. 2000 and March 2001; altogether we obtained 88 data points per star. Using two different time series analysis techniques (Scargle periodogram and CLEAN) we found 543 periodic variable stars with periods between 0.2 days and 15 days. Also, 484 irregular variable stars were identified using a $\\chi^2$-test. In addition we have carried out nearly simultaneous observations in $V$, $R_{\\mathrm{C}}$ and a narrow-band H$\\alpha$ filter. The photometric data enable us to reject background and foreground stars from our sample of variable stars according to their location in the $I_{\\mathrm{C}}$ vs $(R_{\\mathrm{C}}-I_{\\mathrm{C}})$ colour-magnitude and $(R_{\\mathrm{C}}-H\\alpha)$ vs $(R_{\\mathrm{C}}-I_{\\mathrm{C}})$ colour-colour diagrams. We identified 405 periodic variable and 184 irregular variable PMS stars as cluster members using these two different tests. In addition 35 PMS stars for which no significant variability were detected could be identified as members using an $H\\alpha$ emission index criterion. This yields a total of 624 PMS stars in NGC\\,2264, of which only 182 were previously known. Most of the newly found PMS stars are fainter than $I_{\\mathrm{C}} \\simeq 15\\,\\mathrm{mag}$ and of late spectral type ($\\ga$M2). We find that the periodic variables, as a group, have a smaller degree of variability and smaller $H\\alpha$ index than the irregular variables. This suggests that the sample of periodic variables is biased towards weak-line T~Tauri stars (WTTSs) while most of the irregular variables are probably classical T Tauri stars (CTTSs). We have quantified this bias and estimated that the expected fraction of WTTSs among PMS stars in the cluster is $77\\%$. This is relatively close to the fraction of WTTSs among the periodic variables which is $85\\%$. We also estimated the total fraction of variables in the cluster using only two well selected concentrations of PMS stars called NGC\\,2264\\,N \\& S in which we can easily estimate the total number of PMS stars. We find that at least $74\\%$ of the PMS stars in the cluster with $I_{\\mathrm{C}}\\leq 18.0\\,\\mathrm{mag}$ were found to be variable (either periodic or irregular) by our study. This number shows that our search for PMS stars in NGC\\,2264 through extensive and accurate photometric monitoring is very efficient in detecting most PMS stars down to at least $I_{\\mathrm{C}}=18.0\\,\\mathrm{mag}$. ",
        "introduction": "The angular momentum problem is one of the longest standing conundrums in astronomy. Briefly stated, molecular cloud cores have four to six orders of magnitude too much angular momentum to be incorporated into a star. Without shedding their excess angular momentum, stars could not form. When and how they rid themselves of angular momentum is still an open question. It is intimately connected to the issue of disk formation and evolution, to the formation of binary systems and, during the pre-main sequence (PMS) phase, to the magnetic interaction between stars and disks which likely regulates accretion and drives outflows and jets (e. g. Bodenheimer, \\cite{bodenh}, Bouvier et al. \\cite{bouvier1997}, Stassun et al. \\cite{stassun1999}, Mathieu \\cite{bob}).  Observationally, our best hope of constraining the angular momentum evolution of PMS stars is to obtain stellar rotation rates for objects at a variety of masses and ages. Fortunately, it is possible to do this photometrically because PMS stars have strong surface magnetic fields which produce large cool spots (e.~g. Feigelson \\& Montmerle \\cite{feigelson}). It is often the case that the spot pattern is sufficiently asymmetric and stable that the rotation period can be found by photometric monitoring (e.~g. Rydgren \\& Vrba \\cite{rydgren1983}, Herbst et al. \\cite{herbst1994}, Stassun et al. \\cite{stassun1999}). The advantage of this photometric method compared with $v\\,sin\\,i$ measurements is that it yields directly the rotation period independent of the inclination of the star and it can be very accurate even for slow rotating stars. Herbst et al. (\\cite{herbst2001}, \\cite{herbst2002}) have used this method and obtained rotation periods for 369 T Tauri stars (TTSs) in the $\\sim\\!1\\,\\mathrm{Myr}$ old open cluster Orion Nebular cluster (\\object{ONC}). They found that the period distribution for higher mass stars (i.~e. $M \\geq 0.25 M_{\\sun}$) is bimodal with peaks at 2 and 8 days. The bimodality is interpreted as an effect of disk-star interactions in PMS stars: Slow rotators have been magnetically locked to their disks which prevent them spinning up with increasing time or equivalently with decreasing radius. Magnetic coupling to an accretion disk was first proposed more than a decade ago as a dominant braking mechanism for PMS stars (Camenzind \\cite{camenzind}; K\\\"onigl \\cite{koenigl}; Shu et al. \\cite{shu}).  However, we also mention a principal limitation of the photometric method for studying the rotational periods among TTSs. As we will show the photometric method can much more easily measure the periods among the weak-line T Tauri stars (WTTSs) but only for a fraction of the classical T Tauri stars (CTTSs) due to their higher irregular variability. This ``photometric noise'' often prevents the detection of a periodic signal in the photometric data. Therefore any photometric monitoring of a young open cluster (like the ONC or NGC\\,2264) will be biased towards WTTSs in the sense that for these stars the fraction of measurable periods will be higher.  From the results of the rotational study in the ONC several questions naturally arise: 1) Is the period distribution similar in other clusters (i.~e. does environment play a role)? 2) How many PMS stars interact with their disks and how strongly does this affect the angular momentum evolution? 3) How does the period distribution evolve with time (i.~e. do most stars decouple from their disks on time scales of the age of an ONC star and spin up with conserved angular momentum or do disks and disk-locking persist for longer time scales)? To answer these questions it is necessary to measure rotation periods of large samples of PMS stars in clusters with different ages.  Aside from the ONC, the open cluster \\object{NGC\\,2264} is perhaps the best target for a detailed rotational study, since it is sufficiently nearby (760~pc, Sung et al. \\cite{sung}), fairly populous, and with an estimated age of 2 -- 4\\,Myr (Park et al. \\cite{park}) it is about a factor of 2 -- 4 older than the ONC.  Prior to our study a few monitoring programs have been carried out in \\object{NGC\\,2264} (e.g. Kearns \\& Herbst \\cite{KH}) which  altogether yielded  only about 30 published rotation periods for PMS stars. Vogel \\& Kuhi \\cite{vogel1981} and Soderblom et al. (\\cite{soderblom1999}) reported rotational velocities ($v\\,\\sin\\,i$) for in total about 60 low mass NGC\\,2264 stars. An early variability study has been carried out by Nandy \\& Pratt (\\cite{nandy1972}) who identified 26 optical variable stars in the cluster. Neri, Chavarr\\'{\\i}a, \\& de Lara (\\cite{neri1993}) reported optical and near-infrared photometry for 50 stars in NGC\\,2264 and detected for $70\\%$ of these stars optical variability. Due to the limiting sensitivity of these studies most of the stars have spectral types earlier than K5.  About 200 PMS stars and some 20 PMS candidates have been identified in NGC\\,2264 prior to our study using a variety of methods including $H\\alpha$ spectroscopy (Herbig \\cite{herbig}; Ogura \\cite{ogura}; Marcy \\cite{marcy1980}; Rydgren \\cite{rydgren1979}), $H\\alpha$ narrow band photometry (Sung et al. \\cite{sung}; Park et al. \\cite{park}) or ROSAT X-ray flux measurements (e.g. Flaccomio et al. \\cite{flaccomio}). A PMS membership catalogue that summarises (most) these results is available in Park et al. (\\cite{park}).  Vasilevskis, Sanders, \\& Balz (\\cite{vasil1965}) have carried out a proper motion study of 245 stars in the region of the cluster. However, the suggested membership probabilities must be handled with caution because several stars with strong H$\\alpha$-emission which clearly must be members have low membership probability while other (non-emission) stars with high membership probability are clearly no members (see e.~g. the H$\\alpha$-emission study by Rydgren \\cite{rydgren1979}). King (\\cite{king1998}) and Soderblom et al. (\\cite{soderblom1999}) have in total determined lithium abundances of 35 cluster members. Recently, circumstellar disk candidates have been identified in NGC\\,2264 by Rebull et al. (\\cite{rebull2002}).  In this paper we present the results of an extensive photometric monitoring program in NGC\\,2264, in which we monitored $\\sim 10600$ stars over a very broad magnitude range ($9.8 \\leq I_{\\mathrm{C}} \\leq 21$). This extensive monitoring program allowed us to discover about 400 new PMS stars in the cluster and in addition we could measure the rotation periods of 405 PMS stars down to the substellar limit.  In the following section we describe the observational and data reduction strategy. Sect.~\\ref{photometrie} describes the methods we have used for obtaining absolute and relative photometry. In Sect.~\\ref{timeseries} we describe the time series analysis. In Sect.~\\ref{member} we discuss in detail how we selected the PMS stars among the periodic and irregular variables we have found. The nature of the stellar variability is examined in Sect.~\\ref{natvari}. The identification of additional PMS members for which no variability could be detected is described in Sect.~\\ref{PMScand}. In Sect.~\\ref{comleteness} we estimate the fraction of variable PMS stars and the completeness level of our PMS sample. In Sect.~\\ref{fractions} we investigate if our sample of periodic variables is representative for all cluster members. Our major conclusions are summarised in Sect.~\\ref{summary}. The obtained extensive rotational period data and possible consequences for the disk-locking scenario will be discussed in a separate paper (see Lamm et al. \\cite{lamm}, hereafter Paper II). ",
        "conclusions": "\\label{summary}  We have carried out an extensive search for rotation periods and variability in NGC\\,2264. The main results of this investigation are as follows:  \\begin{enumerate}  \\item We have obtained absolute photometry in the Cousins $V$, $R_{\\mathrm{C}}$, and $I_{\\mathrm{C}}$ bands and instrumental $H\\alpha$ magnitudes for about 10600 stars with $9.8\\,\\mathrm{mag} \\leq I_{\\mathrm{C}}\\leq 21\\,\\mathrm{mag}$ in NGC\\,2264. Relative light curves with more than 80 data points each were obtained for each of these stars in the $I_{\\mathrm{C}}$ band. Two different periodogram analysis techniques and a $\\chi^2$ test yielded a sample of 543 periodic and 484 irregular variables with $11.4\\,\\mathrm{mag} \\leq I_{\\mathrm{C}} \\leq 19.7\\,\\mathrm{mag}$ in the observed field.  \\item In order to check the PMS nature of the periodic and irregular variables two different PMS selection criteria were applied. These criteria are their locations in the $I_{\\mathrm{C}}$ vs $(R_{\\mathrm{C}}-I_{\\mathrm{C}})$ colour-magnitude diagram and their location in the $(R_{\\mathrm{C}}-H\\alpha)$ vs $(R_{\\mathrm{C}}-I_{\\mathrm{C}})$ colour-colour diagram. In this way 405 periodic and 184 irregular variable PMS stars could be selected from the sample. This is a enormous increase in the number of known rotation periods compared to the about 30 published periods prior to this study. The rotation periods of the periodic variables are typically between a quarter of a day and 15 days.  \\item Using an enhanced $H\\alpha$ emission index as a PMS indicator we could identify an additional 35 PMS stars which are \"non-variable\" according to our study. Together with the 589 variable PMS stars we found a total of 624 PMS stars in NGC\\,2264. Before our study only 182 PMS members and candidates were known. Most of our newly found PMS stars are fainter than $I_{\\mathrm{C}}\\simeq 15\\,\\mathrm{mag}$ and of late spectral type ($\\ga$M2).   \\item Comparison of the distribution of the standard deviation $\\sigma$ in the light curves as well as the distribution of the $H\\alpha$ emission index for periodic variables with the corresponding distributions for the irregular variables strongly suggests that our method for determining the rotation periods preferentially selects weak-line T~Tauri stars (WTTSs) while classical T~Tauri stars (CTTSs) are underrepresented. This interpretation is confirmed by the calculated fractions of periodic variables among WTTSs and CTTSs which, however,  show that the bias of the measured rotation periods towards WTTSs is relatively small. We have estimated that $23\\%$ of the PMS stars in the cluster are CTTSs but only $15\\%$ of the periodic variables are CTTSs. The corresponding numbers for WTTS are $77\\%$ and $85\\%$ respectively.  \\item The different properties of the two subclasses of T Tauri stars (CTTSs and WTTSs) are supposed to be due to the different mechanisms which produce the variability. While the variability of WTTSs is likely due to large (magnetically) cool spots the variability of CTTSs is believed to be mainly caused by hot spots resulting from accretion flows onto the star. On CTTSs the periodic brightness variations resulting from cool spots (which are believed to exist on their surface too) are often undetectable due to the ``noise'' from the irregular variability. Determining the rotational periods for a larger fraction of CTTSs by photometric monitoring would therefore require more extensive observations (e.~g. larger time base and dense sampling).  \\item We have estimated that we have detected variability among at least $70\\%$ of the PMS stars (defined according to our PMS tests) in NGC\\,2264 with $I_{\\mathrm{C}}\\leq 18.0\\,\\mathrm{mag}$ and $(R_{\\mathrm{C}}-I_{\\mathrm{C}})\\leq 1.8\\,\\mathrm{mag}$. From these numbers it is evident that extensive photometric monitoring with a similar accuracy as in our study is a powerful tool for finding most PMS stars. We emphasise that this method is highly efficient only if there is enough data with a proper sampling to allow a search for periodic variables because many PMS members were identified by their periodicity and not just by their variability. We conclude that the sample of variable PMS stars obtained in this study is a representative subset of the cluster members and it is not expected that additional monitoring programs will substantial increase the number of known PMS stars at least for stars with $I_{\\mathrm{C}}\\leq 18.0\\,\\mathrm{mag}$.  \\end{enumerate}"
    },
    "0402/astro-ph0402141_arXiv.txt": {
        "abstract": "We present new photometric observations of EG\\,And, Z\\,And, BF\\-Cyg, CH\\,Cyg, CI\\,Cyg, V1329\\,Cyg, TX\\,CVn, AG\\,Dra, RW\\,Hya, AG\\,Peg, AX\\,Per, IV\\,Vir and the peculiar M giant V934\\,Her, which were made in the standard Johnson $UBV(R)$ system. QW\\,Sge was measured in the Kron-Cousin $B,~V,~R_{\\rm C},~I_{\\rm C}$ system and for AR\\,Pav we present its new visual estimates. The current issue gathers observations of these objects to December 2003. The main results can be summarized as follows: {\\bf EG\\,And}: The primary minimum in the $U$ light curve (LC) occurred at the end of 2002. A 0.2 -- 0.3\\,mag brightening in $U$ was detected in the autumn of 2003. {\\bf Z\\,And}: At around August 2002 we detected for the first time a minimum, which is due to eclipse of the active object by the red giant. Measurements from 2003.3 are close to those of a quiescent phase. {\\bf BF\\,Cyg}: In February 2003 a short-term flare developed in the LC. A difference in the depth of recent minima was detected. {\\bf CH\\,Cyg}: This star was in a quiescent phase at a rather bright state. A shallow minimum occurred at $\\sim$\\,JD~2\\,452\\,730, close to the position of the inferior conjunction of the giant in the inner binary of the triple-star model of CH\\,Cyg. {\\bf CI\\,Cyg}: Our observations cover the descending branch of a broad minimum. {\\bf TX\\,CVn}: At/around the beginning of 2003 the star entered a bright stage containing a minimum at $\\sim$\\,JD~2\\,452\\,660. {\\bf AG\\,Dra}: New observations revealed two eruptions, which peaked in October 2002 and 2003 at $\\sim$\\,9.3 in $U$. {\\bf AR\\,Pav}: Our new visual estimates showed a transient disappearance of a wave-like modulation in the star's brightness between the minima at epochs E = 66 and E = 68 and its reappearance. {\\bf AG\\,Peg}: Our measurements from the end of 2001 showed rather complex profile of the LC. {\\bf RW\\,Hya}: Observations follow behaviour of the wave-like variability of quiet symbiotics. {\\bf AX\\,Per}: In May 2003 a 0.5\\,mag flare was detected following a rapid decrease of the light to a minimum. {\\bf QW\\,Sge}: CCD observations in $B~,V,~R_{\\rm C},~I_{\\rm C}$ bands cover a period from 1994.5 to 2003.5. An increase in the star's brightness by about 1\\,mag was observed in all passbands in 1997. Less pronounced brightening was detected in 1999/2000. {\\bf V934\\,Her}: Our observations did not show any larger variation in the optical as a reaction to its X-ray activity. ",
        "introduction": "The symbiotic stars are currently understood as interacting binary systems consisting of a cool giant and a hot compact star, which is in most cases a white dwarf. Typical orbital periods are between 1 and 3 years, but they can be significantly larger. The mass loss from the giant represents the primary condition for appearance of the symbiotic phenomenon. A part of the material lost by the giant is transferred to the more compact companion via accretion from the stellar wind or Roche-lobe overflow. This process generates a very hot ($T_{\\rm h} \\approx\\,10^5$\\,K) and luminous ($L_{\\rm h} \\approx\\,10^2 - 10^4\\,\\rm L_{\\odot}$) source of radiation. On the basis of the way in which the generated energy is being liberated, we distinguish two phases of symbiotic binary. {\\em Quiescent phases} during which the hot component releases its energy approximately at a constant rate and spectral distribution. Generally, we observe a wave-like variation in their LCs as a function of the orbital phase. During {\\em active phases} the hot component radiation changes significantly, which leads to a 2-3\\,mag brightening of the object in the optical. A common feature of active phases is a high-velocity mass ejection.  Generally, the hot radiation ionizes a fraction of the neutral circumbinary material, which gives rise to a strong nebular emission. This component of radiation is physically displaced from the hot star and its optically thick part can be very complex in its shape. In addition, its location and shaping in the binary depend on the level of the activity. As a result we often observe unexpected variation in the LCs as, for example, flares, drops in brightness, the effect of eclipses and outbursts. A very interesting feature of variability in this respect is the effect of eclipses, which is very sensitive to physical displacement and radiative contributions of individual components in the system. In the case that a significant fraction of radiation at the wavelength under consideration comes from the region which is subject to eclipse, a minimum in the LC is well observable. In the opposite case, the eclipse effect is very faint. Therefore the eclipse effect can be observed only at specific brightness phases, at which the radiative contribution from a pseudophotosphere in the optical rivals that from the nebula.  Accordingly, to reveal the above mentioned peculiarities in the LCs of symbiotic stars, it requires a very careful long-term monitoring programme. In this paper we present the recent observational results of such our programme obtained during the period December 2001 to December 2003. We note that this paper continues the work of Skopal {\\it et al.} (2002, hereafter S+02, and references therein). ",
        "conclusions": ""
    },
    "0402/astro-ph0402377_arXiv.txt": {
        "abstract": "The currently available empirical evidence on the star formation processes in the extreme, high-pressure environments induced by galaxy encounters, mostly based on high-resolution {\\sl Hubble Space Telescope} imaging observations, strongly suggests that star {\\it cluster} formation is an important and perhaps even the dominant mode of star formation in the starburst events associated with galaxy interactions. The production of young massive star clusters (YSCs) seems to be a hallmark of intense star formation, particularly in interacting and starburst galaxies. Their sizes, luminosities, and mass estimates are entirely consistent with what is expected for young Milky Way-type globular clusters (GCs). YSCs are important because of what they can tell us about GC formation and evolution (e.g., initial characteristics and early survival rates). They are also of prime importance as probes of the formation and (chemical) evolution of their host galaxies, and of the initial mass function in the extreme environments required for cluster formation. Recent evidence lends support to the scenario that Milky Way-type GCs (although more metal rich), which were once thought to be the oldest building blocks of galaxies, are still forming today. ",
        "introduction": "Stars rarely form in isolation. In fact, it is well known that the vast majority of stars in the Galaxy, and also in nearby galaxies, are found in groups ranging from small associations, containing some 100 M$_\\odot$, to compact, old ``globular'' and young massive clusters. The nearest examples of these latter objects include the Galactic star-forming regions NGC 3603 and Westerlund 1, and the giant starburst region 30 Doradus with its central star cluster R136 in the Large Magellanic Cloud.  Although the older Galactic {\\it open} clusters (with ages of several Gyr) are undoubtedly gravitationally bound objects, their lower masses compared to the globular cluster population, and more diffuse structures make them much more vulnerable to disk (and bulge) shocking when they pass through the Galactic disk (or close to the bulge), thus leading to enhanced cluster evaporation. These objects are therefore unlikely globular cluster progenitors. It appears that the conditions for the formation of compact, massive star clusters -- that have the potential to eventually evolve into globular cluster-type objects by the time they reach a similar age -- are currently not present in the Galaxy, or at best to a very limited extent (e.g., Westerlund 1; Hanson 2003).  The production of luminous, massive yet compact star clusters seems to be a hallmark of the most intense star-forming episodes, or starbursts. The defining properties of young massive star clusters (YSCs; with masses often significantly in excess of $M_{\\rm cl} = 10^5 {\\rm M}_\\odot$) have been explored in intense starburst regions in several dozen galaxies, often involved in gravitational interactions of some sort (e.g., Holtzman et al. 1992, Whitmore et al. 1993, O'Connell et al. 1994, Conti et al. 1996, Watson et al. 1996, Carlson et al. 1998, de Grijs et al.  2001, 2003a,b,c,d,e).  An increasingly large body of observational evidence suggests that a large fraction of the star formation in starbursts actually takes place in the form of such concentrated clusters, rather than in small-scale star-forming ``pockets''. YSCs are therefore important as benchmarks of cluster formation and evolution. They are also important as tracers of the history of star formation of their host galaxies, their chemical evolution, the initial mass function (IMF), and other physical characteristics in starbursts.  In a detailed study of the young star cluster population associated with the fading starburst region ``B'' in the prototype nearby starburst galaxy M82 (de Grijs et al. 2001), we concluded that the last tidal encounter between M82 and its large neighbour spiral galaxy M81, which occurred about $500 - 800$ Myr ago (Brouillet et al. 1991, Yun 1999) had a major impact on what was probably an otherwise normal, quiescent, disk galaxy. It caused a concentrated starburst, resulting in a pronounced peak in the clusters' age distribution, roughly 1 Gyr ago (de Grijs et al. 2001, 2003c, Parmentier et al.  2003). The enhanced cluster formation decreased rapidly, within a few hundred Myr of its peak. However, general star formation activity continued in the galactic disk of M82's region B, probably at a much lower rate, until $\\sim 20$ Myr ago.  The evidence for the decoupling between cluster and field star formation is consistent with the view that star cluster formation requires special conditions, such as large-scale gas flows, in addition to the presence of dense gas (cf.  Ashman \\& Zepf 1992, Elmegreen \\& Efremov 1997).  Such conditions occur naturally in the extreme environments of gravitationally interacting galaxies.  Using optical observations of the ``Mice'' and ``Tadpole'' interacting galaxies (NGC 4676 and UGC 10214, respectively) -- based on a subset of the Early Release Observations obtained with the {\\sl Advanced Camera for Surveys} on board the {\\sl Hubble Space Telescope (HST)} -- and the novel technique of pixel-by-pixel analysis of their colour-colour and colour-magnitude diagrammes, we deduced the systems' star and star cluster formation histories (de Grijs et al. 2003e).  In both of these interacting systems we find several dozen YSCs (or, alternatively, compact star-forming regions), which overlap spatially with regions of active star formation in the galaxies' tidal tails and spiral arms (from a comparison with H$\\alpha$ observations that trace active star formation; Hibbard \\& van Gorkom 1996). We estimate that the main gravitational interactions responsible for the formation of these clusters occurred $\\sim 150 - 200$ Myr ago.  In fact, we show that star cluster formation is a major mode of star formation in galactic interactions, with $\\ge 35$\\% of the active star formation in encounters occurring in star clusters (de Grijs et al. 2003e). The tidal tail of the Tadpole system is dominated by star forming regions, which contribute $\\sim 70$\\% of the total flux in the {\\sl HST} {\\it I}-band filter (decreasing to $\\sim 40$\\% in the equivalent {\\it B}-band filter). If the encounter occurs between unevenly matched, gas-rich galaxies then, as expected, the effects of the gravitational interaction are much more pronounced in the smaller galaxy. For instance, when we compare the impact of the interaction as evidenced by star cluster formation between M82 (de Grijs et al. 2001, 2003b,c) and M81 (Chandar et al.  2001), or the star cluster formation history in the ``Whirlpool Galaxy'' M51 (Bik et al. 2003), which is currently in the process of merging with the smaller spiral galaxy NGC 5194, the evidence for enhanced cluster formation in the larger galaxy is minimal if at all detectable.  The NGC 6745 system represents a remarkably violently star-forming interacting pair of unevenly matched galaxies.  The optical morphology of NGC 6745, and in particular the locations of the numerous bright blue star-forming complexes and compact cluster candidates, suggest a recent tidal passage by the small northern companion galaxy (NGC 6745c; nomenclature from Karachentsev et al.  1978) across the eastern edge of the main galaxy, NGC 6745a.  The high relative velocities of the two colliding galaxies likely caused ram pressure at the surface of contact between both galaxies, which -- in turn -- is responsible for the triggering of enhanced star and cluster formation, most notably in the interaction zone in between the two galaxies, NGC 6745b (cf.  de Grijs et al.  2003a).  The smaller galaxy, however, does not show any significant enhanced cluster formation, which is most likely an indication that it contains very little gas.  For the NGC 6745 young cluster system we derive a median age of $\\sim 10$ Myr.  Based on the age distribution of the star clusters, and on the H{\\sc i} morphology of the interacting system, we confirm the interaction-induced enhanced cluster formation scenario once again. NGC 6745 contains a significant population of high-mass clusters, with masses in the range $6.5 \\le \\log( M_{\\rm cl}/{\\rm M}_\\odot ) \\le 8.0$. These clusters do not have counterparts among the Galactic globular clusters (e.g., Mandushev et al. 1991, Pryor \\& Meylan 1993), but are similar to or exceed the spectroscopically confirmed mass estimates of the ``super star clusters'' (SSCs) in M82 (e.g., M82 F and L; see Smith \\& Gallagher 2001) and the Antennae galaxies (Mengel et al.  2002). We caution, however, that these massive SSC candidates may not be gravitationally bound objects, but more diffuse star forming regions or aggregates of multiple unresolved clusters instead. Nevertheless, we measure a very compact effective radius for the most massive object ($M_{\\rm cl} \\simeq 5.9 \\times 10^8 {\\rm M}_\\odot$) of only $R_{\\rm eff} \\sim 16$ pc. However, this object appears very elongated, or may in fact be a double cluster. We should also keep in mind that this high mass estimate is a strong function of the (low) metallicity assumed; if we had assumed (higher) solar metallicity for this object, the derived age would have been significantly smaller ($\\sim 10-20$ Myr vs. $\\sim 1$ Gyr assumed in our study), and the mass could be smaller by a factor of $\\ga 10$. Even so, if we could confirm this mass estimate spectroscopically, either of the subcomponents would be the most massive cluster known to date, significantly exceeding cluster W3 in NGC 7252, which has a mass of about $(3-18) \\times 10^7 {\\rm M}_\\odot$, depending on the age, metallicity and IMF assumed (Schweizer \\& Seitzer 1998; Maraston et al.  2001, 2004). Our detection of such massive SSCs in NGC 6745, which are mostly located in the intense interaction zone, supports the scenario that such objects form preferentially in the extreme environments of interacting galaxies. ",
        "conclusions": ""
    },
    "0402/cond-mat0402227_arXiv.txt": {
        "abstract": "We determine an asymptotic expression of the blow-up time $t_{coll}$ for self-gravitating Brownian particles or bacterial populations (chemotaxis) close to the critical point. We show that $t_{coll}=t_{*}(\\eta-\\eta_{c})^{-1/2}$ with $t_{*}=0.91767702...$, where $\\eta$ represents the inverse temperature (for Brownian particles) or the mass (for bacterial colonies), and $\\eta_{c}$ is the critical value of $\\eta$ above which the system blows up. This result is in perfect agreement with the numerical solution of the Smoluchowski-Poisson system. We also determine the asymptotic expression of the relaxation time close but above the critical temperature and derive a large time asymptotic expansion for the density profile exactly at the critical point. ",
        "introduction": "\\label{sec_introduction}  Recently, several papers have focused on the blow-up properties of a cluster of self-attracting particles \\cite{wolansky,biler1,csr,herrero,dolbeault,rosier, crs,sc,biler2,anomalous,sc2,guerra}. This concerns in particular the ``isothermal collapse'' \\cite{aa} of self-gravitating Brownian particles and the ``chemotactic aggregation'' \\cite{ks} of bacterial populations in biology. These systems have a similar mathematical structure and they are described, in a first approximation, by the Smoluchowski-Poisson system \\cite{crrs}. The Smoluchowski equation is a particular Fokker-Planck equation in physical space involving a diffusion due to Brownian motion and a drift. In the standard Brownian theory developed by Einstein, the drift is due to an external potential \\cite{risken}. Alternatively, we can consider the more complicated situation in which the potential is produced by the particles themselves. If we assume a Newtonian type of interaction, we have to couple the Smoluchowski equation to the Poisson equation.  It can be shown that the Smoluchowski-Poisson system decreases a Lyapunov functional $F[\\rho]$ which can be interpreted as a Boltzmann free energy \\cite{gtpre}. Furthermore, the type of evolution depends on the value of a dimensionless parameter $\\eta$. For $\\eta\\leq\\eta_{c}=2.51755132...$, the Smoluchowski-Poisson system evolves to an equilibrium state which minimizes the free energy at fixed mass. It corresponds to an isothermal distribution of particles similar to isothermal stars and isothermal stellar systems \\cite{bt}. On the contrary, for $\\eta>\\eta_{c}$, there is no possible equilibrium state and the system blows up. This is the case for self-gravitating Brownian particles below a critical temperature $T_{c}$ and for bacterial populations above a critical mass $M_{c}$. It is found that the collapse is self-similar and that it develops a finite time singularity, {\\it i.e.} the central density becomes infinite in a finite time $t_{coll}$ \\cite{crs,sc}. Then, a Dirac peak is formed in the post-collapse regime \\cite{sc2}.  It is clear that the collapse time $t_{coll}$ depends on the distance to the critical point $\\eta_{c}$ and that it should diverge as $\\eta\\to \\eta_{c}$.  More precisely, by using heuristic arguments, it is argued in \\cite{crs} that $t_{coll}\\sim t_{*} (\\eta_{c}-\\eta)^{-1/2}$. This scaling is consistent with numerical simulations of the Smoluchowski-Poisson system. However, the approach of \\cite{crs} is qualitative and does not provide the numerical value of $t_{*}$.  After some introductory material presented in Section 2, we develop a systematic procedure in Section 3, inspired from \\cite{kavallaris}, which confirms the scaling law $t_{coll}\\sim t_* (\\eta_{c}-\\eta)^{-1/2}$ and leads to the explicit value of $t_{*}$. In section 4, we study the relaxation time $\\tau$ to equilibrium below $\\eta_c$ and show that it diverges like $\\tau\\sim c_\\eta (\\eta-\\eta_{c})^{-1/2}$, where $c_\\eta$ is given explicitly. In Section 5, we derive a systematic large time expansion of the density profile exactly at $\\eta_c$. In particular, we show that the central density approaches its equilibrium value according to $\\rho_0-\\rho(r=0,t)\\sim \\frac{c_\\rho}{t}$, where $c_\\rho$ is a universal constant which is explicitly computed. Finally, in Section 6, we determine the pulsation period of bounded isothermal spheres described by the Euler-Jeans equations close to the critical point. ",
        "conclusions": "In this paper, we have obtained the asymptotic expressions of the blow-up time and relaxation time of self-gravitating Brownian particles and biological populations close to the critical point of collapse in $d=3$. An excellent agreement is obtained by solving numerically the Smoluchowski-Poisson system. This study confirms and improves the results obtained in \\cite{crs} on the basis of heuristic arguments. A possible extension of this work would be to consider the case of polytropic distributions arising in case of anomalous diffusion \\cite{anomalous}. More generally, we could consider the generalized Smoluchowski equation (81) proposed in \\cite{gtpre} which is valid for an arbitrary equation of state. This equation can provide (among other applications) a generalized chemotactic model valid when the diffusion coefficient or the drift term depends on the density. These extensions will be considered in future works."
    },
    "0402/astro-ph0402606_arXiv.txt": {
        "abstract": "We introduce a fluid dynamics algorithm that performs with nearly spectral accuracy, but uses finite-differences instead of FFTs to compute gradients and thus executes 10 times faster.  The finite differencing is not based on a high-order polynomial fit. The polynomial scheme has supurb accuracy for low-wavenumber gradients but fails at high wavenumbers. We instead use a scheme tuned to enhance high-wavenumber accuracy at the expense of low wavenumbers, although the loss of low-wavenumber accuracy is negligibly slight. A tuned gradient is capable of capturing all wavenumbers up to 80 percent of the Nyquist limit with an error of no worse than 1 percent. The fact that gradients are based on finite differences enables diverse geometries to be considered and eliminates the parallel communications bottleneck. ",
        "introduction": "The spectral fluid algorithm (Canuto et. al. 1987) uses Fast Fourier Transforms (FFTs) to compute gradients, the most precise means possible. Finite-difference gradients based on a polynomial fit execute faster than FFTs but with less accuracy, necessitating more grid zones to achieve the same resolution as the spectral method. The loss of accuracy outweighs the gain in speed and the spectral method has more resolving power than the finite-difference method.  We introduce an alternative finite-difference formula not based on a polynomial fit that executes as quickly but with improved accuracy, yielding greater resolving power than the spectral method.  In section \\ref{finite} we derive high-wavenumber finite difference formulas and exhibit their effect on the resolving power of a turbulence simulation in section \\ref{flop}. In section \\ref{mhd} we apply finite differences for the purpose of mimicking the spectral algorithm, and then proceed to other applications in section \\ref{applications}. ",
        "conclusions": ""
    },
    "0402/astro-ph0402430_arXiv.txt": {
        "abstract": "In this paper we report the results of the first ever time-dependent general relativistic magnetohydrodynamic simulations of the magnetically dominated monopole magnetospheres of black holes. It is found that the numerical solution evolves towards a stable steady-state  solution which is very close to the corresponding force-free solution found by Blandford and Znajek. Contrary to the recent claims, the particle inertia does not become dynamically important near the event horizon and the force-free approximation provides a proper framework for magnetically dominated magnetospheres of black holes. For the first time, our numerical simulations show the development of an ultra-relativistic particle wind from a rotating black hole. However, the flow remains Poynting dominated all the way up to the fast critical point. This suggests that the details of the so-called ``astrophysical load'', where the electromagnetic energy is transferred to particles, may have no effect on the efficiency of the Blandford-Znajek mechanism. ",
        "introduction": "It is now widely believed that the relativistic jets generated in active galactic nuclei, galactic microquasars and, presumably, during gamma ray bursts are powered by rapidly rotating black holes.  This paradigm is largely based on the theoretical results of Blandford and Znajek (1977), who argued that, under typical conditions of astrophysical black holes, their rotational energy can be efficiently extracted in the form of magnetically dominated relativistic wind.  In fact, the original model of Blandford and Znajek was based on the approximation of steady-state force-free degenerate electrodynamics (FFDE, see also Macdonald \\& Thorne 1982).  In this approximation the dynamical role of magnetospheric plasma is reduced to providing perfect conductivity in the space surrounding the black hole, and the particle inertia is totally neglected.  Not all mathematical and physical aspects of this electrodynamic model are well understood. Similar classical systems, e.g. the Faraday disc, involve the poloidal current driven over the surface of a rotating conductor by a non-electrostatic electromotive force. Moreover, the inertia of this conductor is a key factor in the origin of this force. Obviously, there is no such conductor in the Blandford-Znajek model. Often, this missing element of the model is artificially introduced in the form of the so-called ``membrane'', or the ``stretched horizon'', located somewhat above the real event horizon \\cite{MT82,TPM}. Mathematically, this is consistent with the so-called ``horizon boundary conditions'' employed in the Blandford-Znajek model \\cite{Z77}.   However, the ``Membrane paradigm'' cannot always provide correct insights into the black hole electrodynamics (this would be a great mystery if it did.)  Since the membrane is very close to the real horizon, the outgoing fast waves emitted by the membrane are highly red-shifted. Moreover, the membrane cannot at all communicate with the outer space by means of Alfv\\'en waves because the inner Alfv\\'en surface is quite far away from the horizon.  As the result, these waves cannot transport angular momentum and adjust the poloidal electric current of the outgoing wind as they do so, for example, in stellar winds of magnetised stars. Punsly and Coroniti (1990) used these causality arguments to expose the lack of physical clarity in the electrodynamic model of Blandford and Znajek (1977) and its representation in the membrane paradigm.  In fact, they concluded that there is no physically meaningful electromotive force in the Blandford-Znajek model but only the artificial one which has been effectively introduced via Znajek's horizon condition.  If so, then the Blandford-Znajek model and, hence, its steady-state solutions have to be nonphysical.  Since these FFDE solutions are, in fact, proper mathematical solutions their nonphysical nature has to give itself away via instability \\cite{Pun-Cor,Pun01}.  In fact, more or less the same criticism was applied to the GRMHD model of black hole magnetospheres developed by Phinney (1982) as it also utilized Znajek's horizon condition.  In addition to this criticism, Punsly and Coroniti developed a rather attractive alternative GRMHD theory of black hole magnetospheres where particle inertia played the key role in the extraction of rotational energy of black holes \\cite{Pun-Cor1,Pun01}.  While almost everywhere in their models the magnetic field was completely dominating, in some regions within the black hole ergosphere the plasma particles were accelerated to such a high Lorentz factor that their inertia could no longer be ignored.  This led to strong inertially driven electric currents flowing across the magnetic field lines, so that the system resembled the famous Faraday disc rather closely.  In spite of the great astrophysical importance of the electrodynamic/MHD mechanism of extraction of rotational energy of black holes and admirable efforts of B.Punsly, the dramatic clash between the Blandford-Znajek and Punsly-Coroniti models seemed to remain largely unnoticed by the astrophysical community for a whole decade.  Only very recently, following the developments in the theory of force-free electrodynamics \\cite{Uch,Gru,Kom02a} and the impressive progress in numerical methods for relativistic astrophysics, e.g. \\cite{Pons,Kom99,Koi99,Koi03,Kol,Gam,Del-Zan,Dev-Haw}, things began to change. One of most interesting recent results is concerned with the conjecture of Punsly and Coroniti on instability of the Blandford-Znajek monopole solution. In \\cite{Kom01b} this conjecture was tested by means of time-dependent general relativistic FFDE simulations and was found to be incorrect. Contrary to the appealing Punsly-Coroniti causality arguments, this steady-state monopole solution of Blandford and Znajek is asymptotically stable.  There is, however, the question of validity of the FFDE approximation which has to be fully investigated before reaching any final conclusion.  There are many examples in physics and astrophysics where particular approximations fail to describe certain natural phenomena because one or another physical factor ignored in the approximation becomes important. These examples taught us always to check the limits of applicability of our theoretical models and never be satisfied by their mere self-consistency. The very relevant example can be found in the theory of magnetically dominated pulsar winds where particle inertia becomes an important factor as a result of acceleration by electromagnetic field (Mestel 1999, Section 13.2.3) Other potential examples in the theory of black hole magnetospheres are discussed by Punsly (2001, Sections 8 and 9).  Recently, Punsly (2004) argued that FFDE is deficient near the event horizon where ``the plasma attains an infinite inertia in a global sense''. Macdonald and Thorne \\shortcite{MT82} also anticipated a breakdown of the FFDE approximation near the event horizon, though no detailed explanations were given.  Thus, in the problem of electrodynamic/MHD mechanism of extraction of rotational energy of black holes we have to verify that the particle inertia can indeed be ignored everywhere without oversimplifying this problem.  This can be done via GRMHD simulations of black hole magnetospheres.    The main goal of this paper is to explore the role of particle inertia in the pair-filled monopole magnetospheres of black holes by means of such simulations. In particular, we need to know whether incorporation of initially small particle inertia will eventually lead to strong deviations from the corresponding electrodynamic solution both locally, e.g. near the event horizon, and globally. By latter we mean changes of the global system of electric current, the angular velocity of magnetic field lines, and, hence, the efficiency of the energy extraction.  Since, at present only codes for the perfect relativistic MHD are readily available, the monopole magnetic configuration appears to be the most suitable one for this purpose. Indeed, this configuration allows a dissipation free FFDE solution \\cite{BZ,Kom01b}.  Dipolar magnetospheres, like the one considered by Koide \\shortcite{Koi03}, are less suitable because the corresponding electrodynamic solution involves a strong current sheet located within the ergosphere \\cite{Kom02b,Kom04}.  This indicates that the approximation of perfect MHD is also likely to break down in this configuration. ",
        "conclusions": "In this paper we described the first ever GRMHD simulations of monopole magnetospheres of rotating black holes. Given the central role played by the monopole problem in the development of the general theory of black hole magnetospheres it is not surprising that the results of these simulations allow to re-examine a number of important issues of the theory.  The main conclusion which follows from these results is the validity of the FFDE approximation at least in the case of the monopole magnetic configuration considered by  Blandford and Znajek(1977). Inertia of initially rarefied magnetospheric plasma does not grow dynamically important neither in any localized  regions near the black hole horizon nor in a global sense. The system of poloidal electric current and the efficiency of energy extraction are basically the same as in our GRMHD solution as in the FFDE solution found earlier \\cite{Kom01b}.  This means that in the magnetospheres of black holes there exists an electromotive force that has nothing to do with particle inertia.  In a separate paper \\cite{Kom04} we show that, in great contrast to the Faraday disc or a magnetized stellar wind, the poloidal currents in the Blandford-Znajek model are driven by the so-called ``gravitationally induced'' electric field, which was first discovered by Wald (1974). Contrary to the conclusion reached in \\cite{Pun-Cor}, this field cannot be screened within the black hole ergosphere by any static distribution of electric charge. We also show there that Znajek's horizon condition is not a boundary condition after all, but a regularity condition imposed at the fast critical point of the ingoing wind in the limit of vanishing particle inertia. This proves the legitimacy of its utilization in the steady-state solutions by Blandford-Znajek(1977) and Phinney(1982).  In addition, the numerical solution shows a number of other interesting features that deserve discussing.  Our GRMHD solution remains very close to the FFDE one all the way up to the fast critical surface of the outgoing wind and even at the fast critical surface the wind is still Poynting flux dominated. Further away the electromagnetic energy may be transferred to the wind particles. However, the details of this energy transfer as well as the details of the interaction between the wind and its surrounding which is responsible for the observational phenomena like superluminal jets, radio galaxies etc. cannot effect the wind solution in the sub-fast region. This makes us wonder whether the key global properties of the Blandford-Znajek mechanism, such as its efficiency, are sensitive at all to the nature of the so-called ``astrophysical load'' \\cite{MT82,TPM}.  Given the potential importance of this finding let us show that a RMHD flow may indeed stay Poynting flux dominated at the fast critical point. For this purpose we consider a one dimensional flow in flat spacetime. In the limit of cold MHD, the ratio kinetic energy flux to the Poynting flux is $$ \\kappa = \\rho W^2/B_t^2, $$ where $B_t$ is the tangential component of magnetic field. At the fast point the flow velocity $$ v^2=b^2/(b^2+\\rho), $$ where $b^2 = B_n^2 + (B_t/W)^2$, $B_n$ being the normal component of the magnetic field. From these two equations one finds that at the fast point \\begin{equation} \\kappa = \\frac{1}{W^2} + \\frac{B_n^2}{B_t^2}. \\end{equation} Thus, the flow is Poynting dominated provided the fast speed is ultra-relativistic and the magnetic field is predominantly tangential.  These are the first GRMHD simulations where an ultrarelativistic outflow is produced by a black hole. In all previous simulations of this sort \\cite{Koi99,Koi00,Koi03,Kom01a,Gam,Dev-Haw} no such high velocity outflows were reported,  which was a bit worrying given the well known results of various astrophysical observations. In more realistic simulations of the future, which will include both the accretion disc and its magnetised corona,  it may be possible to obtain not only ultra-relativistic but also well collimated outflows.  Since the fast and the Alfven surfaces of the outgoing wind are relatively far away from the event horizon, one may expect the outgoing wind to be more or less accurately described by the flat space monopole solution. According to \\cite{Bes97} the ratio of the fast suface radius, $r_f$, and the Alfv\\'en surface radius, $r_a$, in the equatorial plane of this solution is $$ r_f/r_a \\approx \\sigma^{1/3}, $$ whereas the Lorentz factor at the fast surface $$ W_f=\\sigma^{1/3}, $$ where $\\sigma$ is the magnetization papameter. In our simulations $\\sigma^{1/3} \\approx 4.08$, $r_f/r_a=6.60$, and $W_f=4.35$. Thus, the numerical and the analytical results agree quite well.  The presented results also suggest a way of dealing with such more realistic problems of computational astrophysics of black holes which involve both magnetically dominated ultrarelativistic components and  particle dominated slow components.  Namely, one may use the same GRMHD approximation to model all components of the system but impose an upper limit on the magnetization of the ultrarelativistic component, via appropriate floor values for both the pressure and density of matter. Although, there is nothing particularly new in such an approach and utilization of similar lower/upper limits is widely spread in computational practice, the accurate representation of magnetically dominated component in our simulations is reassuring."
    },
    "0402/astro-ph0402576_arXiv.txt": {
        "abstract": "We analyze cross-correlation functions between Gamma-Ray Burst (GRB) hosts and surrounding galaxies. We have used data obtained with the Very Large Telescope at Cerro Paranal (Chile), as well as public Hubble Space Telescope data. Our results indicate that Gamma-Ray Burst host galaxies do not reside in high galaxy density environments. Moreover, the host-galaxy cross-correlations show a relatively low amplitude. Our results are in agreement with the cross-correlation function between star-forming galaxies and surrounding objects in the HDF-N. ",
        "introduction": "The origin of cosmological Gamma-Ray Bursts (GRBs) is still one of the outstanding problems in modern astronomy \\citep{vanpa97}. Over the past half decade, the discovery of localized transients in the small error boxes of GRBs has led to intense multiwavelength campaigns and many advances have been made in understanding the nature of the bursts and their afterglows throughout the electromagnetic spectrum (see for instance \\citet{vanpa00}).\\\\ The scenario for the origin of GRBs include either the merging of two collapsed objects such as black holes or neutron stars merge \\citep{nara92} or the cataclysmic destruction of massive stars (supernovae or hypernovae) \\citep{woos93,macwoos99}. In the first scenario, GRBs can occur long after the star formation episodes in the nucleus since the merging characteristic time scale is as long as $\\gtapp$ 1 Gyr. In the second scenario the host galaxies are likely to be on-going significant star formation.  \\citet{bloom02} analyzed the observed offset distribution of Gamma-Ray Burst from their host galaxies, and found a strong connection of GRB localization with the UV light of their host galaxies. This provides a significant observational evidence for the correlations between GRBs and star-forming regions.  The characteristics of the physical environment surrounding the bursts may provide strong constrains for the origin of these events. As suggested by \\citet{macwoos99} in the collapsar model, GRBs could be produced by rotating massive stars in which the accretion of a helium core leads to the prompt formation of a black hole. Low metallicity in the stellar envelope reduces the mass loss and inhibits the loss of angular momentum by the star. Star formation in regions of low metallicity are likely to generate GRBs so that dwarf and sub-luminous galaxies would be preferred host of these events (see for instance \\citet{floch03} ).  This model is supported by the Ly$\\alpha$ emission from GRB host galaxies which show a preference to be metal poor and with on-going star formation in a dust poor environment \\citep{Fynbo02,Fynbo03}.  Analyzing the properties of the host galaxies of GRBs, \\citet{floch03} found that GRB host galaxies are sub-luminous in the $K$-band and exhibit very blue colors, not comparable to the luminous star-burst and/or reddened sources observed at high redshifts in the infrared and sub-millimeter deep surveys.  There is new evidence that at least some Gamma-Ray Bursts are related to supernovae: the association of GRB 980425 with the peculiar Type Ic  SN 1998bw \\citep{gala98}, and also the recent spectroscopic discovery of the SN 2003dh associated with GRB 030329 \\citep{hjorthnat,stanek03}. These observations provide a very solid confirmation of the association of some GRBs with the death of massive stars in strong star formation regions.  At the present it is unknown if GRB host galaxies are preferentially located in dense environments, or if there is any correlation between the local density of galaxies and the presence of a GRB. So far, there are only two works on the environment of GRB host galaxies. \\citet{Fynbo02} studying the Ly$\\alpha$ emission of two GRB fields (GRB 000301C \\& GRB 000926), found a number of galaxies at the same redshift of GRB hosts without signs of overdensity at small scales. Moreover, the lack of blank fields at a similar depth prevented these authors to conclude if GRB hosts reside in overdense regions. From an analysis of photometric redshifts of galaxies in the field of GRB 000210, \\citet{goro03} found that there is no obvious concentration of galaxies around the host. Another case, the afterglow of GRB 980613 was located close to a very compact object inside a complex region consisting of star-forming knots and/or interacting galaxy fragments \\citep{Hjorth02,djo03}. If these host galaxies are associated with the underlying large scale structure of the universe, then they should show similar galaxy clustering properties of normal galaxies.\\\\ In this paper we examine the nature of the density enhancement of galaxies around GRB host galaxies by studying the cross-correlation function between GRB hosts and the surrounding galaxy distribution.  This paper is organized as follows: Section 2 describes the sample of GRB host galaxies analyzed, the source extraction and photometry techniques. We analyze galaxy source counts in Section 3. Section 4 deals with the angular-cross correlation analysis of galaxies around GRB host galaxies. Finally we discuss our results in Section 5. ",
        "conclusions": "We have analyzed different data sets corresponding to deep imaging in the field of Gamma-Ray Burst hosts. A remarkable low correlation amplitude at small angular scales is detected from the cross-correlation of  GRB's and the surrounding galaxies in all samples analyzed. The reliability of the results presented in this work can be judged from the lack of correlation of stars in {\\tt ISAAC} and \\hst/STIS images. By contrast, the galaxy-galaxy correlation function obeys the usual power-law shape with a significant signal at small separations. Given that these correlation functions were computed using angular positions, any effect in real space would be diluted by projection of foreground and background objects.  A comparison with a similar data set, centered in USS sources, indicates the very different environment of these two types of objects. While USS sources clearly reside in rich environments, GRBs are likely to reside in a typical or even lower galaxy density environment. Moreover the amplitude of the autocorrelation function of galaxies is larger than that of the GRB-galaxy cross-correlation.  In the \\hst/STIS fields GRB-tracers cross-correlation functions are significantly lower than early-type-tracers cross correlations. In addition, GRB targets have a lower cross correlation amplitude than ISO and early-type galaxy targets. This suggests that the star formation events associated to GRBs occur in particularly low density environments, a result that is supported by the fact that objects formed in global underdense regions are expected to be biased to low luminosity, consistent with GRB hosts characteristics \\citep{floch03}."
    },
    "0402/astro-ph0402095_arXiv.txt": {
        "abstract": "We present a 3-dimensional lensing analysis of the $z=0.16$ supercluster A901/2, resulting in a 3-D map of the dark matter distribution within a $3\\times10^{5}[\\Mpc]^3$ volume. This map is generated from a combined catalogue of 3-D galaxy coordinates together with shear estimates, using $R$-band imaging and photometric redshifts from the COMBO-17 survey. To estimate the 3-D positions and masses of the main clusters in the supercluster from lensing alone, we perform a $\\chi^2$-fit of isothermal spheres to the tangential shear pattern around each cluster as a function of redshift. Motivated by the appearance of a second cluster behind A902 in galaxy number density, we also fit a two-cluster model to A902.  We then present the first 3-D map of the dark matter gravitational potential field, $\\Phi$, using the Kaiser-Squires (1993) and Taylor (2001) inversion methods. These maps clearly show the potential wells of the main supercluster components, including the new cluster behind A902, and demonstrates the applicability of 3-D dark matter mapping and projection free-mass-selected cluster finding to current data. Finally, we develop the halo model of dark matter and galaxy clustering and compare this with the auto- and cross-correlation functions of the 3-D gravitational potential, galaxy number densities and galaxy luminosity densities measured in the A901/2 field. We find significant anti-correlations between the gravitational potential field and the galaxy number density and luminosities, as expected due to baryonic infall into dark matter concentrations. We find good agreement with the halo model for the number densities and luminosity correlation functions, but some disagreement with the shape of the gravitational potential correlation function, which we attribute to finite-field effects. ",
        "introduction": "Gravitational lensing is a valuable method for measuring the matter distribution in the Universe (see e.g. Mellier 1999, Bartelmann \\& Schneider 2001). Light rays are deflected by the gravitational potential along their paths causing distortions in the shapes of background galaxies together with a change in their local number counts. These measurables are sensitive to all forms of matter, whether visible or dark, which brings gravitational lensing to the fore in studies of non-baryonic matter in the Universe. For most regions of the sky, observations of shape distortions (shear) reveal that the phenomenon is very weak ($\\approx 1$\\% change in ellipticity of an object). However we are still able to measure this effect by averaging the ellipticity of many galaxies to overcome the random intrinsic ellipticity of galaxies.    Many studies have used weak shear measurements to measure the two dimensional projected matter distribution of regions of space. This has led to precise understanding of the masses and mass profiles of galaxy clusters (see e.g. Tyson et al 1990, Kaiser \\& Squires 1993, Bonnet et al 1994, Squires et al 1996, Hoekstra et al 1998, Luppino \\& Kaiser 1997, Gray et al 2002). The 2-D shear field is also used to measure the large-scale structure distribution, and conseqeuently to measure cosmological parameters (see e.g. van Waerbeke et al 2001, Hoekstra 2002, Bacon et al 2002, Refregier et al 2002, Jarvis et al 2003, Brown et al 2003, Pen et al 2003).  Recently, there has been much interest in using redshift information with weak gravitational lensing in order to add depth to our picture of the mass distribution in the Universe. Lensing tomography has been developed as a means of studying the growth of the weak lensing power spectrum as a function of redshift (e.g. Seljak 1998, Hu 1999, 2002, Huterer 2002, King \\& Schneider 2002b), and the use of redshift information to remove the impact of intrinsic galaxy alignments on lensing measurements has been explored (Heymans \\& Heavens 2002, King \\& Schneider 2002a,b, Heymans et al 2003). Observationally, Wittman et al (2001, 2002) have used 3-D shear information to directly measure the mass and redshift of two clusters discovered through this shear field.  A further significant development has been the realisation that a full reconstruction of the 3-D gravitational potential and dark matter density field is possible using just weak shear and redshifts (Taylor 2001). A 3-D gravitational lensing analysis is clearly of importance for cosmology. With a 3-D analysis we will have, for the first time, a way of imaging the full dark matter distribution from cluster scales upwards and over a order of magnitude in redshift, independent of the galaxy distribution. With such depth one can hope to visually see the growth of structure. A 3-D lensing survey can also be used to construct a mass-selected galaxy cluster catalogue, which would have many important cosmological uses. A three dimensional lensing survey would remove the projection effects that contaminate standard 2-D projected surveys (White, van Waerbeke \\& Mackey, 2002). In the case of individual clusters this can cause biases in mass estimation due to both foreground and background structures. With a 3-D lensing analysis, both the mass and position of clusters can be found independently of the baryonic content.   Following Taylor (2001), Hu \\& Keeton (2003) subsequently developed a pixelised version of the 3-D reconstruction, while Bacon \\& Taylor (2003) examined the practical implementation of Taylor's (2001) method. Bacon \\& Taylor (2003) demonstrated that one can hope to reconstruct the gravitational potential on cluster mass scales with ground-based shear surveys to a depth of $z=1$ by including photometric redshifts with accuracy $\\Delta z \\simeq 0.05$.  To date the idea of reconstructing the full 3-D gravitational or matter distribution has only been applied to simulated data. However, data of this quality already exist; the COMBO-17 survey (Wolf et al 2001) is a 17-band photometric redshift survey with accuracy  $\\Delta z = 0.05$ throughout $0<z<0.8$ and has already been used for weak shear studies of a supercluster (Gray et al 2002), large-scale structure (Brown et al 2002), star formation efficiency (Gray et al 2003) and galaxy alignment effects (Heymans et al 2003). It is therefore an ideal dataset with which to carry out a first direct reconstruction of the 3-D gravitational potential.    In this paper we carry out a full 3-dimensional analysis of the A901/2 supercluster field from the COMBO-17 survey, including measurements of the 3-D shear, lensing potential and gravitational potential fields. We describe the necessary methodological tools for these measurements in Section 2. In particular we summarise our approach to measurement of tangential shear, the Kaiser-Squires inversion for the lensing potential and reconstruction of the gravitational potential. In Section 2 we also develop the halo model of mass and galaxy clustering allowing an analysis of the correlations between the 3-D gravitational potential and the galaxy number and luminosity densities. In Section 3 we describe the COMBO-17 dataset, discussing the imaging and photometric redshift data, and the means of obtaining a shear catalogue. We describe the A901/2 field in detail, recalling the relevant shear measurements obtained by Gray et al (2002) for this supercluster.  In Section 4 we measure the 3-D tangential shear of the COMBO-17 A901/2 supercluster field, where we fit 3-D shear models to the clusters, following the approach of Wittman et al (2001, 2002), and measure the mass and redshift for the clusters directly from the gravitational distortions. We examine the evidence for a lensing signature of a cluster behind the supercluster at $z=0.48$, and constrain its mass using a two-cluster fit to the shear data. This is compared to the results of a 2-D lensing analysis by Gray et al (2002). In Section 5 we calculate the 3-dimensional lensing potential and reconstruct the 3-dimensional dark matter potential of the A901/2 supercluster field. We find significant potential wells at the expected position of the supercluster in 3-D, and also find a significant potential well at the position of the background cluster. We discuss the measurement significance for each of these detections.  In Section 6 we explore a natural application of 3-D gravity measurements to measure the auto- and cross-correlation of the 3-D gravitational potential with two other quantities: the galaxy number density, and the luminosity distribution. We find that there are anti-correlations for each of these quantities, and describe how this can be understood in terms of baryonic infall into dark matter concentrations. We compare our results with theoretical predictions of the halo model. Finally, we present our conclusions in Section 7. ",
        "conclusions": "In this paper we have for the first time measured the 3-D gravitational potential of a volume of space from gravitational lensing. This has been possible by accurately measuring the gravitational shear on the A901/2 supercluster field of the COMBO-17 survey, combining these shear measurements with photometric redshifts, and applying the reconstruction techniques of Kaiser \\& Squires (1993), Taylor (2001) and Bacon \\& Taylor (2003). We have described the COMBO-17 survey, including our measurements of weak shear (Gray et al 2002) and the photometric redshift accuracy of the survey, $\\Delta z = 0.05$ for $0<z<0.8$.  As a first step towards full 3-D analysis of the A901/2 field, we have measured the tangential shear around the supercluster members as a function of redshift. By fitting a singular isothermal sphere shear model to the shear pattern behind each of these clusters (c.f. Wittman et al 2001, 2002) we estimated the clusters' velocity dispersion and redshifts from weak lensing alone. We found best-fit two-parameter velocity dispersions and redshifts of $\\sigma_v=840^{+100}_{-105} \\kms$ and $z=0.30^{+0.10}_{-0.16}$ for A901a, $\\sigma_v=840^{+170}_{-115}\\kms$, $z_{\\rm lens}=0.37^{+0.14}_{-0.11}$ for A901b and $\\sigma_v=760^{+220}_{-200}$, $z_{\\rm lens}=0.38^{+0.23}_{-0.20}$ for A902. These measurements are consistent with photometrically determined redshifts of the clusters, $z=0.16$. If we fix the redshifts of the cluster to the photometric redshift of $z=0.16$, we find $\\sigma_v=680^{+25}_{-90}\\kms$ for A901a; $\\sigma_v=600^{+40}_{-85}\\kms$ for A901b; and $\\sigma_v=520^{+55}_{-90}\\kms$ for A902. We have also compared with velocity dispersions of the clusters found from a 2-D lensing analysis (Gray et al 2002) and found good agreement. Any differences are easily understood as due to the different weighting of shear data.  Examination of the 3-D number density of galaxies revealed the existence of a cluster, CB1, at $z=0.48$ behind the A902 cluster. This prompted us to investigate a two-cluster shear analysis for A902 which provided results consistent with the presence of this second cluster. Our initial two-cluster analysis failed to improve on the redshift estimates of the two clusters, due to the difficulty a global parametric fit to the data has in dealing with outlying shear values. However, by fixing the positions of the clusters, either with a weak prior by restricting the redshifts of the clusters to lie below $z=1$, or with a strong prior of fixing the redshifts to the known cluster positions, the total mass of the two clusters was well determined, and each cluster velocity dispersion could be marginalised over to give $\\sigma_{v,1}=470_{-280}^{+100}\\kms$ for A902 and $\\sigma_{v,2}=730_{-340}^{+160}\\kms$ for CB1.  We have calculated the 3-D lensing potential, $\\Delta \\phi$, for this volume of space using the methods of Kaiser \\& Squires (1993) generalised to a 3-D grid. We have then used this lensing potential to reconstruct the 3-D gravitational potential, following the reconstruction method of Taylor (2001) and Bacon \\& Taylor (2003). To improve the signal-to-noise in the reconstructed potential field we Wiener filtered the potential field in the redshift direction, preserving the properties of the reconstructed field in the transverse direction.  The recovered gravitational potential was found to include the largest troughs at the positions of the supercluster members, with a peak signal-to-noise ratio at the central pixel of $S/N\\approx 3$. Hence we have demonstrated that 3-D dark matter potential mapping is achievable with currently available shear and photometric redshift data, such as the COMBO-17 survey.  In addition to the A901/2 supercluster potential, a further significant ($3\\sigma$ per pixel) trough was found at $z=0.48$ behind A902, corresponding to the background cluster CB1 seen in the 3-D number density and luminosity. This demonstrates that potential mapping will be a useful tool in the detection of clusters from lensing data, including clusters behind clusters. Indeed this suggests an algorithm for the construction of a mass-selected cluster catalogue, which is free of projection effects. One first constructs a 3-D potential field from the shear and photometric data and searches for the largest potential wells. The positions of these clusters can be accurately determined from this galaxy number density field, while the cluster masses can be estimated from parametric fits.  Finally, we have developed the halo model to predict the auto- and cross-correlation functions of the 3-D potential field, the galaxy number density field and, for the first time, the galaxy luminosity-density field, taking into account the geometry of the survey and pixelisation. We compared these predictions with the measured auto- and cross-correlations between the gravitational potential and the galaxy number density, and between the gravitational potential and the 3-D luminosity distribution in the COMBO-17 A901/2 field. We found that peaks in the baryonic matter concentration were strongly correlated to troughs in the gravitational potential, as one would expect due to baryons preferentially being found at dark matter concentrations. We also found reasonable agreement between the amplitudes of the measured and predicted correlations -- although the statistics of the A901/2 field are dominated by the four clusters. This introduces oscillations in the measured correlation functions as we move in and out of clusters. In particular, the correlation length of the potential field is shorter than we might have expected from the clustering of the galaxies. This is most likely due to the removal of the mean, gradients and parabolic terms in each redshift slice of the potential field, including the supercluster slice. This effect will be alleviated with larger area surveys (see Bacon \\& Taylor, 2003, for a further discussion of this effect).   In conclusion, the success of this study in measuring the 3-D gravitational potential bodes well for future 3-D lensing studies.  We have seen that 3-D potential mapping will be useful for constructing mass-selected galaxy cluster catalogues with minimal projection effects. By extending and comparing with the halo model, we have also shown that the analysis of haloes within this framework can be extended to groups and clusters. With advent of better shear and photometric redshift data, from e.g. the SNAP survey (Massey et al 2003), the prospects look encouraging for 3-D mass mapping large volumes of the universe, to construct mass-selected cluster catalogues, to generate statistical probes of dark matter haloes over a wide range of scales, as a probe of cosmological parameters, and finally to compare the 3-D dark matter distribution with the properties of galaxies."
    },
    "0402/astro-ph0402210_arXiv.txt": {
        "abstract": "We investigate the effect of dark energy on the density profiles of dark matter haloes with a suite of cosmological N-body simulations and use our results to test analytic models.  We consider constant equation of state models, and allow both $\\wQ \\ge -1$ and $\\wQ<-1$. Using five simulations with $\\wQ$ ranging from $-1.5$ to $-0.5$, and with more than $\\sim 1600$ well-resolved haloes each, we show that the halo concentration model of Bullock \\etal (2001) accurately predicts the median concentrations of haloes over the range of $\\wQ$, halo masses, and redshifts that we are capable of probing.  We find that the Bullock \\etal (2001) model works best when halo masses and concentrations are defined relative to an outer radius set by a cosmology-dependent virial overdensity.  For a fixed power spectrum normalization and fixed-mass haloes, larger values of $\\wQ$ lead to higher concentrations and higher halo central densities, both because collapse occurs earlier and because haloes have higher virial densities.  While precise predictions of halo densities are quite sensitive to various uncertainties, we make broad comparisons to galaxy rotation curve data.  At fixed power spectrum normalization (fixed $\\sigma_8$), $\\wQ > -1$ quintessence models seem to exacerbate the central density problem relative to the standard $\\wQ= -1$ model. For example, models with $w \\simeq -0.5$ seem disfavored by the data, which can be matched only by allowing extremely low normalizations, $\\sigma_8 \\la 0.6$.  Meanwhile $\\wQ < -1$ models help to reduce the apparent discrepancy. We confirm that the Jenkins \\etal (2001) halo mass function provides an excellent approximation to the abundance of haloes in our simulations and extend its region of validity to include models with $\\wQ < -1$. ",
        "introduction": "In the prevailing model of galaxy formation, galaxies assemble and evolve in the potential wells established by gravitationally bound haloes of cold and collisionless dark matter (CDM).  Except for some possible difficulties on small scales, the CDM model is remarkably successful in explaining a large number of observations.  However, this success requires an additional ``dark energy'' component, that drives an accelerated cosmic expansion, to be added to the universal energy budget.  While the presence of dark energy is firmly established observationally, measuring its equation of state as well as developing a theoretical understanding of the nature of the dark energy are two of the biggest outstanding problems in cosmology today. Dark energy not only affects the large-scale evolution of the Universe, but also the collapse histories and density structures of dark matter haloes.  Understanding the precise nature of these effects is important for studies that aim to quantify the nature of dark energy using strong (e.g., Sarbu, Rusin, \\& Ma 2001; Huterer \\& Ma 2004; Kuhlen, Keeton, \\& Madau 2004) and weak (e.g.  Hu \\& Jain 2003; Bartelmann et al. 2002) gravitational lensing.  Changing the dark energy model should similarly change expectations for galaxy rotation curves, and could affect one of the main small-scale problems facing CDM -- the central density problem (\\eg Zentner \\& Bullock 2002; McGaugh, Barker, \\& de Blok 2003, and references therein).  In the present paper, we use a suite of N-body simulations to study how halo density profiles change as a function of dark energy equation of state, discuss our results in the context of analytic models, and discuss the observational implications of dark energy on galaxy scales.  The existence of some form of dark energy is supported by a preponderance of data.  Taken together, observations of the magnitude-redshift relation of type Ia supernovae (SNIa; Perlmutter \\etal 1999; Riess \\etal 2001; Knop \\etal 2003; Barris \\etal 2003), the power spectrum of cosmic microwave background (CMB) anisotropy (Spergel \\etal 2003; Tegmark \\etal 2003a), the power spectrum of galaxy clustering (Dodelson \\etal 2003; Tegmark \\etal 2003b), and the luminosity function and baryon fraction of clusters (Allen \\etal 2003) provide nearly unimpeachable evidence for the existence of dark energy.  The most common supposition is that the dark energy takes the form of a cosmological constant or vacuum energy.  In this case, the energy density $\\rho$, and pressure $p$, are related through $p = -\\rho$.  An attractive alternative candidate for the dark energy is the potential energy of a slowly-varying scalar field $\\phi$, or ``quintessence'' (\\eg Ratra \\& Peebles 1988; Caldwell, Dav{\\'e}, \\& Steinhardt 1998).  A convenient parametrization of the dark energy is through an equation of state $\\wQ \\equiv p_{\\phi}/\\rho_{\\phi}$ relating its energy density and pressure.  In general, the equation of state parameter $\\wQ$ is time-varying, but it is useful to model quintessence with a constant equation of state parameter because current observational data sets have limited power to distinguish between a time-varying and constant equation of state (\\eg Kujat et al. 2002).  Useful limits on the equation of state for the dark energy, assuming that it remains constant in time, come from SNIa studies, $-1.67 < \\wQ < -0.62$ ($2\\sigma$; Knop \\etal 2003), and can be refined by combining SNIa data with CMB anisotropy and galaxy clustering statistics yielding $-1.33 < \\wQ < -0.79$ at $2\\sigma$ (Tegmark \\etal 2003b). For our simulations, we adopt an empirical view and study five models with constant $\\wQ$, that span a comparably large range of parameter space: $\\wQ = -1.5, -1.25, -1.0, -0.75, -0.5$.  Our theoretical understanding of halo profiles has improved recently largely through numerical simulations, performed in the context of CDM plus cosmological constant ($\\Lambda$CDM) or standard CDM (SCDM, \\ie $\\Omega_{\\rm M} = 1$, $\\Omega_{\\rm Q}=0$) cosmologies (Navarro, Frenk, \\& White, 1995, 1996, 1997, hereafter NFW; Kravtsov, Klypin, \\& Khokhlov 1997; Ghigna \\etal 1998; Jing 2000; Bullock \\etal 2001; Eke, Navarro, \\& Steinmetz 2001; Wechsler et al. 2002, hereafter W02; Zhao et al. 2003; Hayashi \\etal 2003; Navarro \\etal 2003; for a review, see Primack 2003).  It is generally understood that the final density profiles of haloes are linked closely to their formation histories. Halo central densities are set during an early, rapid-accretion phase, and tend to be proportional to the density of haloes in the Universe at the time of this rapid collapse (W02; Zhao et al. 2003; Tasitsiomi et al. 2003).  Larger values of $\\wQ$ lead to earlier collapse times and also to more rapid collapse of overdensities, thus we expect haloes with higher relative central densities and higher virial densities (see our discussions in \\S~\\ref{sec:Qstruc} and \\S~\\ref{sec:halostruc}).  In \\S\\S~\\ref{sec:sims}-\\ref{sec:DeltaVby2} we quantify these effects and test the expected scalings.  We verify that the analytic technique of Bullock \\etal (2001, B01) for predicting halo concentrations works well when applied directly to constant $\\wQ$ cosmologies, implying that it is fairly generally applicable.   A related issue is the effect of dark energy on the halo mass function.  Although it is not directly observable, theoretical insight into the expected halo mass function is attainable through current N-body simulations.  The effects of dark energy on halo mass functions have been investigated by Bartelmann, Perota, \\& Baccigalupi (2003), Linder \\& Jenkins (2003), Klypin \\etal (2003), Macci\\`o \\etal (2003), and \\L okas, Bode, \\& Hoffman (2003) and the accurate prediction of halo mass functions as well as accretion histories and density structures over a large range of redshifts is necessary in order to take full advantage of the ability of upcoming cluster surveys, such as the Sunyaev-Zeldovich Array (\\texttt{\\small http://astro.uchicago.edu/sza/}), to constrain the dark energy equation of state (see Carlstrom, Holder, \\& Reese 2001 for a review). There appears to be general agreement that halo mass functions can be approximated accurately by the ``universal'' mass function of Jenkins \\etal (2001, hereafter J01) even in models with dark energy.  In \\S~\\ref{sec:mf}, we confirm the results of previous studies of quintessence cosmologies with $\\wQ > -1$, and extend this agreement to models with $\\wQ < -1$.  Klypin \\etal (2003) and Dolag \\etal (2003) have performed previous numerical studies of CDM haloes in quintessence cosmologies.  Where our study overlaps with these, our results are generally in agreement. This study extends, complements, and improves upon previous studies in several ways.  We have investigated the effects of normalizing the power spectrum using clusters vs. CMB, which yield quite different results. We have assembled large catalogues of haloes, with masses and concentrations for more than $1600$ haloes with more than $100$ particles each in each simulation. With this large number of haloes we are able to measure the distribution of concentrations at fixed mass and test the predictions of the B01 model for both the mean relation between concentration and mass and the halo-to-halo scatter.  We also extend the range of quintessence parameter space probed by simulations by exploring two models with $\\wQ<-1$.  The format of this paper is as follows.  After a brief review of the basics of structure formation in models with dark energy in \\S~\\ref{sec:Qstruc}, we go on in \\S~\\ref{sec:halostruc} to describe modifications to the B01 model for halo concentrations in quintessence cosmologies.  In \\S~\\ref{sec:sims}, we describe the set of numerical simulations that we performed.  In \\S~\\ref{sec:nbody_results}, we present the mass functions and concentrations of CDM haloes in these simulations and compare them to the results of the analytic B01 model. In \\S~\\ref{sec:DeltaVby2}, we briefly discuss the central density problem of CDM in light of our results.  In \\S~\\ref{sec:conclusions}, we present a summary of our findings, highlight improvements and differences to previous studies, and discuss the implications of this work.  Throughout this paper, we assume a flat universe with present matter density relative to critical of $\\Omega_{\\rm M}=0.3$, and Hubble parameter $h=0.7$. The quintessence energy density is then $\\Omega_{\\rm Q} = 1 - \\Omega_{\\rm M} = 0.7$.  We refer to a model with a cosmological constant ($\\wQ=-1$) as $\\Lambda$CDM. We use the terms ``quintessence'' and ``$Q$CDM'' to refer to all models which have $\\wQ \\ne -1$. ",
        "conclusions": "\\label{sec:conclusions}  Although the nature of dark energy is unknown, its effects on structure formation can be studied using numerical N-body simulations. We have performed a series of these simulations for a range of dark energy equation of state parameters. Confirming previous findings by Linder \\& Jenkins (2003), Klypin \\etal (2003), Macci\\`o \\etal (2003), and \\L okas \\etal (2003) we show that the J01 formula provides a good fit to halo mass functions even in the presence of non-\\lcdm dark energy. We show that this is true for models with $\\wQ<-1$ as well.  The density structure of dark matter haloes is also affected by dark energy. We have shown how the predictions of the B01 model are modified when dark energy with constant $\\wQ$ is accounted for.  As structure tends to form earlier in models with less negative $\\wQ$, halo concentrations tend to be somewhat higher in these models. These findings are in agreement with the results of Klypin \\etal (2003) and qualitatively agree with Dolag \\etal (2003), although we probe a different range of masses than the latter. The larger number of haloes with NFW profile fits and concentrations in our study allows us to quantitatively test the B01 model. We find that the original ($F=0.01, K=4.0$) over-predicts the concentrations of haloes in our simulations by about $\\sim 12-15\\%$. However, the shape of the mass-concentration relation that we find is the same as in B01, and we find that a slightly modified set of the B01 parameters ($F=0.01, K=3.5$) matches our haloes well. This offset may likely be caused by the lower force resolution of our GADGET simulations compared to the adaptive-refinement code ART used in B01.  For the haloes in our simulations the adopted B01 model accurately reproduces the median concentration-mass relation over a range of masses from $M_{\\rm vir} \\sim 6 \\times 10^{11}$ to $M_{\\rm vir} \\sim 4 \\times 10^{13}\\hinv \\msun$.  We confirm that for a fixed mass halo concentration decrease with redshift as $1/(1+z)$, at least out to $z \\sim 2.5$.  Interestingly, we find that halo concentrations are more easily understood when the halo virial radius is defined in terms of a cosmology-dependent virial overdensity rather than by one that uses a fixed overdensity of $\\rho/\\rho_{\\rm crit}=200$.  The result supports one of the (previously-untested) assumptions of the original B01 model.  Specifically, we argue that halo densities in different cosmological models are influenced both by changes in the overall virialization process of haloes as well as by changes in epoch when the halo cores collapse.  As noted in \\S\\ref{sec:DeltaVby2}, it is important to include both of these physical processes when comparing predictions for galaxy densities to real data, as in Figure~\\ref{fig:DeltaVby2}.  Having confirmed that the B01 model correctly describes the scaling of halo concentrations as a function of mass and redshift even in cosmologies with $w\\ne-1$, we have investigated the effects of dark energy on a comparison between model predictions and observations of central halo densities. Zentner and Bullock (2002) found that the observed distribution of $\\Delta_{\\rm V/2}$ as a function of $V_{\\rm max}$ is inconsistent with the predictions of the B01 model for $\\Lambda$CDM and $\\sigma_8=0.9$. The model predicts haloes that are simply too concentrated (see also Primack 2003). A lower value of $\\sigma_8=0.75$, as used in our simulations, can alleviate this discrepancy, but such models may face other difficulties regarding early reionization (Somerville, Bullock, \\& Livio 2003) and possibly with reproducing the properties of halo substructure (Zentner \\& Bullock 2003).  Including the effects of dark energy, we find that for models with $w>-1$ the problem is exacerbated because haloes collapse earlier and because they have higher virial overdensities.  Note that for the extreme case of $w=-0.5$, even a normalization as low as $\\sigma_8 = 0.6$ seems disfavored by the data. Thus one interesting conclusion is that the rotation curves of galaxies coupled only with a weak prior on the normalization of the power spectrum of density fluctuations seem to disfavor quintessence models with $\\wQ$ significantly larger than $-1$ without measuring the expansion history of the Universe, as is done in SNIa analyses. Models with $w<-1$ do better, and can tolerate higher values of $\\sigma_8 \\sim 0.8$, but not high as high as $\\sigma_8 \\sim 1$, as is suggested the CMB normalization.~\\footnote{Note that using $K=3.5$ instead of $K=4.0$ would not change any of the conclusions regarding galaxy rotation curves and the central density problem. }  As we have pointed out throughout this paper, there have been a number of previous studies of the effects of dark energy ($\\wQ \\ne -1$) on dark matter halo abundances and concentrations. Here we summarize how this work distinguishes itself from these past efforts. We have shown for the first time how results for $\\Delta_{\\rm vir}$ and the linear transfer function are modified for dark energy with $w<-1$. In addition we have expanded on previous studies (Schuecker et al. 2003) of $\\delta_c$ in phantom cosmologies. We have pointed out that there are significant differences in the amplitude and $\\wQ$-dependence of $\\sigma_8$ between normalizations based on the CMB and galaxy cluster abundances. For the first time we have been able to draw a distinction between a definition of concentration based on a cosmology-dependent $R_{\\rm vir}$ and constant $R_{200}$, and found that our simulations prefer the former. Unlike previous studies of the effects of dark energy on halo concentrations (Klypin et al. 2003; Dolag et al. 2003), which analysed a small number of pre-selected haloes, we have performed an analysis of $\\sim1600$ haloes in each simulation. This allowed us to directly test the B01 model for $c(M,z)$ with statistical significance. Furthermore these previous studies focused on group- and cluster-sized haloes, whereas we have extended this study down to galaxy mass haloes. Lastly, we have shown for the first time how different values of $\\wQ$ and power spectrum normalizations affect the theoretical predictions of $\\Delta_{V/2}$ vs. $V_{\\rm max}$. A comparison to observations shows that models with w as low as -0.50 are strongly disfavored for any value of $\\sigma_8$, standard $\\Lambda$CDM models require $\\sigma_8 \\la 0.8$, and models with $\\wQ<-1$ can accomodate higher values of $\\sigma_8$.  As future observations further constrain the nature of dark energy, it will become necessary to extend these types of studies to more realistic models of dark energy.  Future simulations with higher mass and force resolution, including the effects of baryons, and incorporating more realistic dark energy models will further advance our understanding of the interplay between cosmology and dark matter halo structure.  \\medskip  We thank P.~Madau for interesting and helpful discussions.  MK is supported by NSF grant AST-0205738. LES is supported by the Department of Energy grant DE-FG02-91ER40690.  ARZ is supported by The Center for Cosmological Physics at The University of Chicago under NSF PHY 0114422.  JSB is supported by NASA through Hubble Fellowship grant HF-01146.01-A from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Incorporated, under NASA contract NAS5-26555.  JRP is supported by NSF grant AST-0205944. ARZ thanks The Center for Cosmology and Particle Physics at New York University for their hospitality during several visits while this work was in progress.  We thank V. Springel for use of the publicly-available code GADGET and we than U. Seljak and M. Zaldarriaga for use of the publicly-available code CMBFAST."
    },
    "0402/astro-ph0402026_arXiv.txt": {
        "abstract": "Using a Monte Carlo technique, we investigate the center shifts that are expected to occur for the broad H$\\alpha$ wings and the He~II~$\\lambda$~6545 feature as a function of the neutral hydrogen column density in symbiotic stars. These two features are proposed to be formed via Raman scattering of UV continuum around Ly$\\beta$ and He~II~$\\lambda$~1025 emission line by neutral hydrogen. The strengths of these two features are determined by a combination of various physical parameters including the covering factor and the neutral hydrogen column density $(N_{HI})$ of the scattering region. The branching ratio of Raman scattering to Rayleigh scattering for UV radiation around Ly$\\beta$ is a non-linearly increasing function of the wavelength, which results in enhanced Raman optical fluxes redward of the H$\\alpha$ line center as $N_{HI}$ increases.  However, we find that the amount of wing center shift is quite small about $20{\\rm\\ km\\ s^{-1}}$ as $N_{HI}$ increases from $10^{20}{\\rm\\ cm^{-2}}$ to $10^{21}{\\rm\\ cm^{-2}}$. Assuming that He~II~$\\lambda$~1025 emission is characterized by a Gaussian profile, the Raman scattered He~II~6545 feature exhibits near Gaussian profiles with the peak shifted redward for $N_{HI}<10^{22}{\\rm\\ cm^{-2}}$. The redward center shift amounts to $1{\\rm\\ \\AA}$ for $N_{HI}\\sim 10^{20}{\\rm\\ cm^{-2}}$ and decreases as $N_{HI}$ increases up to $N_{HI}\\sim 10^{22}{\\rm\\ cm^{-2}}$, above which no center shift is observed.  The redward peak shift is due to the fact that the incident emission profile is symmetric with respect to He~II~$\\lambda$~1025 line center whereas the Raman conversion rate is increasing towards the line center of Ly$\\beta$. We emphasize that the determination of $N_{HI}$ by locating the exact peak position of the He~II~6545 feature will lift the degeneracy to allow a more accurate estimate of the covering factor of the neutral region, providing strong constraints on the mass loss process occurring in symbiotic stars. ",
        "introduction": "Symbiotic stars are generally believed to be binary systems of a mass losing giant and a hot star, which is usually a white dwarf (Kenyon 1986). The giant component loses a large amount of mass in the form of a slow stellar wind, a significant fraction of which may be ionized by the strong UV radiation emanating from the hot component (e.g. Taylor \\& Seaquist 1984). The physical extent of the ionized region should be sensitively dependent on the strength of the ionizing radiation and the mass loss rate of the giant component. In this respect, studies of ionization structures of symbiotic stars are very important in understanding the mass loss process in symbiotic stars (e.g. Skopal 2003).  Symbiotic stars exhibit very unique spectroscopic features, which are Raman scattered lines by neutral hydrogen. The emission features around 6825 \\AA\\ and 7082 \\AA\\ appearing in more than half of the symbiotic stars are Raman scattered O~VI 1032, 1038 doublets first identified by Schmid (1989). The Raman scattering origin for these features is strongly supported by many spectropolarimetric observations (e.g. Schmid \\& Schild 1994, Schmid et al. 2000, Harries \\& Howarth 1996) and contemporaneous far UV observations (e.g. Espey et al. 1995, Birriel, Espey, Schulte-Ladbeck 1998, 2000). The Raman scattering cross sections for O~VI~$\\lambda\\lambda$ 1032,1038 are of order $10^{-22}{\\rm\\ cm^{2}}$. Therefore, the existence of these features implies that a non-negligible fraction of the scattering region is characterized by a very high H~I column densities $N_{HI}\\ge 10^{22}{\\rm\\ cm^{-2}}$.  Van Groningen (1993) found other Raman scattered features including Raman scattered He~II lines in the symbiotic star RR Telescopii. In particular, Raman scattered He~II line at around 4850 \\AA\\ was found in the spectrum of the young planetary nebula NGC 7027 by P\\'equignot(1997). Because a He~II ion has the same atomic structure with four times large level spacings as a hydrogen atom does, the emission lines associated with the energy levels with even principal quantum numbers of He~II have a slightly shorter wavelengths than those of H~I emission lines.  This implies that the Raman scattering cross sections for these He~II lines are so high that the scattering processes can operate in broader range of objects including young planetary nebulae (Nussbaumer, Schmid \\& Vogel 1989).  In the same line of reasoning, UV continuum radiation around Ly$\\beta$ will be Raman scattered to be redistributed around H$\\alpha$, forming broad wings. In this case, the wing profiles reflect the wavelength dependence of the scattering cross section, which is approximately inversely proportional to the squared wavelength difference of the incident radiation and the Ly$\\beta$ line center. These broad H$\\alpha$ wings with profiles well-fitted by a profile $\\propto\\Delta\\lambda^{-2}$ were found in objects including many symbiotic stars (Lee 2000), some young planetary nebulae such as IC 4997 (Lee \\& Hyung 2000) and M2-9 (Balick 1989), and post AGB stars (van de Steene, Wood, \\& van Hoof 2000). Recently Arrieta \\& Torres-Peimbert (2003) added 12 more objects that are also in late stages of stellar evolution. Considerable mass loss is a common feature for all of these objects, where one can find a large amount of neutral matter surrounding the central hot stars.  It is also interesting that a large fraction of these objects exhibit bipolar nebular morphology (e.g. Schwarz \\& Monteiro 2003, Schwarz \\& Corradi 1992).  Lee et al. (2003) used the 3.6~m Canada-France-Hawaii telescope to perform spectroscopy of the symbiotic star V1016~Cygni. They found the Raman scattered He~II 6545 that is converted from He~II~$\\lambda$~1025 emission line. They adopted a simple scattering geometry to compute the broad H$\\alpha$ wing template profiles and isolate the 6545 feature, from which they were able to put a lower limit of the mass loss rate of the Mira variable. However, the strengths of Raman scattered features are determined by a combination of the covering factor and neutral hydrogen column density ($N_{HI}$) of the scattering region for given incident UV radiation. The degeneracy of the two factors should be lifted in order to put more stringent constraints on the mass loss processes occurring in symbiotic stars.  There exists conspicuous asymmetry in the branching ratio of Raman scattering relative to Rayleigh scattering around Ly$\\alpha$, where the branching ratio is an increasing function of the incident wavelength (Yoo, Bak, \\& Lee 2002). Therefore, continuum radiation redward of Ly$\\beta$ will be more easily Raman scattered than continuum blueward of Ly$\\beta$. This tendency increases for higher scattering optical depths, which may lead to a redward shift of the center location of H$\\alpha$ wings as $N_{HI}$ increases. Furthermore, the Raman conversion rate increases in the wavelength region, where He~II~$\\lambda$~1025 is located. Therefore, if He~II~$\\lambda$~1025 exhibits a symmetric emission profile with the peak at its line center, the peak position of the Raman scattered He~II~6545 will move toward H$\\alpha$, where the amount of shift will be mainly determined by the neutral hydrogen column density.  In this paper, we perform Monte Carlo simulations to compute the center shifts of the H$\\alpha$ wings and the Raman scattered He~II~6545 features for various $N_{HI}$, which may be useful in lifting the degeneracy of the column density and the covering factor that determine the strengths of the Raman scattered features. In the following section we briefly describe the basic ideas of the center shifts. The numerical results from our Monte Carlo calculations are presented in Section 3. Finally in Section 4 we discuss our results and observational ramifications. ",
        "conclusions": "\\subsection{Electron Scattering Wings}  \\begin{figure} \\epsfig{file=fig9.ps, height=11cm, width=8cm} \\caption{Electron scattering wings overplotted to the spectrum of V1016~Cyg obtained with the CFHT. The covering factor of the electron scattering region is assumed to be $C_{el} = 0.05$ and the electron column density $N_e = 10^{24}{\\rm\\ cm^{-2}}$, for which the Thomson optical depth $\\tau_e = 0.67$. The temperature of the electron cloud is chosen to be $T_e = 2.5\\times 10^4{\\rm\\ K}$. The top panel shows the electron scattering wing represented by a solid line, which was fitted using the theoretical profile shown by a dotted line investigated by Weymann (1970).  } \\end{figure}  Electron scattering of H$\\alpha$ line photons is one of physical mechanisms that have been proposed to be the origin of the broad H$\\alpha$ wings in many celestial objects including active galactic nuclei and Wolf-Rayet stars (e.g. Mathis 1970, Weymann 1970, Hillier 1991). We generated electron scattering wings using a similar Monte Carlo technique, which is shown in Fig.~9.  Because the H$\\alpha$ emission line is well-approximated by a single Gaussian $f_\\lambda =1650 e^{-(\\lambda-\\lambda_{H\\alpha})/\\Delta\\lambda_{H\\alpha}}$, we used this Gaussian as an incident radiation. The temperature of the electron cloud is chosen to be $T_e = 2.5\\times 10^4{\\rm\\ K}$. The covering factor of the electron scattering region is assumed to be $C_{el} = 3.5\\times 10^{-2}$ and the electron column density $N_e = 10^{24}{\\rm\\ cm^{-2}}$, for which the Thomson optical depth $\\tau_e = 0.67$.  It is also noted that similar profiles will be obtained by a choice of $\\tau_e$ and $C_{el}$ with their product fixed as far as the electron plasma is optically thin.  In the top panel, we show the electron scattering wing produced by a Monte Carlo technique, which is represented by a thin solid line. According to Weymann (1970), electron scattering wings formed in an optically thin region may be approximated by exponential profiles represented by a thick solid line with a functional form \\begin{equation} f_e(\\lambda)=f_0 e^{-(\\lambda-\\lambda_{H\\alpha})/\\Delta\\lambda_e}. \\end{equation} Here, the profile width $\\Delta\\lambda_e=22{\\rm\\ \\AA}$ coincides with the Doppler width corresponding to the thermal electron velocity with $T_e=2.5\\times 10^4{\\rm\\ K}$, from which the electron temperature is strongly constrained. The apparent discrepancy with the theoretical profile is attributed to our choice of rather high $\\tau_e$ and the fact that the incident H$\\alpha$ emission profile is not a Dirac delta function but a Gaussian with a finite width.  In the middle and bottom panels, the Monte Carlo results are represented by dotted lines and the solid lines show the CFHT spectrum of V1016~Cyg. The fit is quite satisfactory. However, the Raman scattering origin of H$\\alpha$ wings naturally explains the existence of the He~II~1025 feature. A more stringent constraint may be put from careful analyses of the broad H$\\beta$ wings and He~II~4850 feature, because Raman scattering cross sections and branching ratios differ from those corresponding to features around H$\\alpha$. Future spectroscopic observations covering both H$\\alpha$ and H$\\beta$ regions simultaneously may shed much light on this point.  \\subsection{Observational Ramifications}  An interesting possibility is that if the He~II emission is not isotropic but the continuum is, then the equivalent width of He~II~$\\lambda$~1025 measured by a hypothetical observer located in the neutral scattering region does not have to be equal to that measured by an observer on Earth. According to Mastrodemos \\& Morris (1998), some fraction of the stellar wind from the giant component may be captured by the white dwarf component in symbiotic stars to form an accretion disk. If we accept this picture, the line emission region may take a disk-like shape, where more line radiation is expected to the normal direction of the orbital plane or the accretion disk. Using the ORFEUS data, Schmid et al. (1999) pointed out that O~VI~$\\lambda\\lambda$~1032, 1038 doublets exhibit different profiles from those of their Raman scattered counterparts, which appears to support this picture. In this case, for an observer with a large inclination to the orbital plane, the equivalent width that is estimated from the Raman scattered He~II 6545 will be smaller than that estimated from the far UV spectroscopy performed for the same object.   In the work of Lee et al. (2003), the neutral column density $N_{HI}$ was assumed to be the value that gives a unit scattering optical depth, of which the validity has not been pursued in a rigorous way. Therefore, if this choice of $N_{HI}$ is lower than the true value, then the covering factor of the scattering region should be larger than they estimated, and vice versa. However, in order to locate the peak position very accurately, we have to subtract [N~II]~$\\lambda$~6548 feature reliably using the 3 times stronger [N~II]~$\\lambda$~6584, which requires very careful data reductions. [N~II] lines are weak or absent in the higher nebular density S-type systems, for which case the line center of the Raman scattered He~II$\\lambda$~6545 can potentially be determined without this complication. As is shown by this study, an accurate value of the mass loss rate of the giant can be obtained by a more realistic model equipped with a range of $N_{HI}$ with independent information about the equivalent width of He~II~$\\lambda$~1025.  It is particularly notable that the continuum photon number flux is diluted by a factor of $(\\lambda_i/\\lambda_o)^2$ due to the incoherency of the Raman scattering process. Because the equivalent width is a measure of the line photon number flux normalized by the continuum photon number flux, the computation of the equivalent width of the He~II~$\\lambda$~1025 presented by Lee et al. (2003) is erroneous and overestimated by a factor of 6.4. The correct equivalent width should be about 0.51 \\AA\\ which is similar to the value  adopted in the current work.  It is very interesting to note that the asymmetric branching ratio of the Raman scattering cross section gives rise to asymmetric wing profiles with the stronger red part around H$\\alpha$. Regarding this, spectropolarimetric studies of Raman scattered O~VI 6825 and 7082 features by Harries \\& Howarth (1996) show the flip of polarization direction occurring in the far red part, implying the existence of a third emission (or scattering) component moving away from the main scattering region (or emission region). Therefore, when examining this kind of component that is moving away, the effect from the asymmetry of atomic physics should be corrected first. More spectropolarimetry will be very useful to investigate the existence of similar phenomenon in H$\\alpha$ wings.  Less scattering cross sections and more branching possibilities may lead to more diverse features in Raman scattered lines around H$\\beta$ and higher Balmer series lines. Therefore, simultaneous spectroscopy for the Raman scattered features around H$\\beta$ and higher Balmer series lines will shed more light on the mass loss processes and the ionization structures of symbiotic stars. Similar center shift phenomena are expected to Raman scattered He~II features formed blueward of higher Balmer series lines. Because center shifts are also sensitive to relative kinematics between the UV emission source region and the scattering region, careful data analyses will be required. We conclude that Raman features provide very important and interesting diagnostic tools for investigating the mass loss processes and ionization structures in symbiotic stars."
    },
    "0402/astro-ph0402356_arXiv.txt": {
        "abstract": " ",
        "introduction": "Cosmic-ray particles --- fully ionized atomic nuclei --- are the only matter which reaches the earth from outside the solar system and are accessible for direct investigations. The energy of these particles can exceed even $10^{20}$~eV, but their acceleration mechanisms and the sites of origin are still under discussion.  The energy spectrum is of non-thermal origin and follows a power law $dN/dE\\propto E^{\\gamma}$ over many orders of magnitude. The spectrum steepens at energies around 4~PeV from a spectral index $\\gamma\\approx-2.7$ to $\\gamma\\approx-3.1$. This feature is commonly called the \\knie\\ and its explanation is generally believed to be a corner stone in understanding the origin of cosmic rays, providing answers to one of the key questions of astroparticle physics.  Many approaches have been discussed in the literature to describe the origin, acceleration, and propagation of cosmic rays and several mechanisms have been proposed to explain the \\knie. In this article the basic ideas of theoretical models are briefly mentioned and the resulting energy spectra and mass compositions are described in section~\\ref{theosect}.  The steeply falling cosmic-ray flux as function of energy requires increasing apertures and exposure times. Presently, above 1~PeV cosmic rays are experimentally accessible in ground-based installations only. These detectors do not measure the primary particles, instead secondary particles are measured, produced in high-energy interactions in the atmosphere, forming an extensive air shower.  This makes the interpretation of the indirect measurements very difficult and the results obtained depend on the understanding of high-energy interactions in the atmosphere.  The actual experimental status is briefly discussed in section~\\ref{pgsect}. ",
        "conclusions": "\\begin{table*} \\caption{Synopsis of all models discussed.} \\label{sumtab} \\renewcommand{\\arraystretch}{1.0} \\begin{tabular}{ll}\\hline Model&Author(s)\\\\ \\hline Source/Acceleration:&\\\\ ~~~Acceleration in SNR       & Berezhko \\& Ksenofontov \\cite{berezhko}   \\\\ ~~~Acceleration in SNR + radio galaxies & Stanev et al. \\cite{stanev}  \\\\ ~~~Acceleration by oblique shocks  & Kobayakawa et al. \\cite{kobayakawa}\\\\ ~~~Acceleration in variety of SNR& Sveshnikova \\cite{sveshnikova}\\\\ ~~~Single source model & Erlykin \\& Wolfendale \\cite{wolfendale}  \\\\ ~~~Reacceleration in the galactic wind& V\\\"olk and Zirakashvili \\cite{voelk}\\\\ ~~~Cannonball model & Plaga                 \\cite{plaga}  \\\\ Propagation/Leakage from Galaxy:&\\\\ ~~~Minimum pathlength model & Swordy  \\cite{swordy}  \\\\ ~~~Anomalous diffusion model & Lagutin et al. \\cite{lagutin}  \\\\ ~~~Hall diffusion model & Ptuskin et al. \\cite{ptuskin}, Kalmykov \\& Pavlov \\cite{kalmykov}  \\\\ ~~~Diffusion in turbulent magnetic fields& Ogio and Kakimoto \\cite{ogio}  \\\\ ~~~Diffusion and drift                   & Roulet et al \\cite{roulet} \\\\ Interactions with background particles:&\\\\ ~~~Diffusion model + photo-disintegration & Tkaczyk \\cite{tkaczyk} \\\\ ~~~Interaction with neutrinos in galactic halo & Dova et al. \\cite{dova}\\\\ ~~~Photo-disintegration (optical and UV photons)& Candia et al. \\cite{candia} \\\\ New interactions in the atmosphere:&\\\\ ~~~Gravitons, SUSY, technicolor& Kazanas \\& Nicolaidis \\cite{kazanasall,kazanasgrav}\\\\ \\hline \\end{tabular} \\end{table*}  All models discussed in this article are summarized in Table~\\ref{sumtab}. Most of them yield similar all-particle spectra. On the other hand, the individual element spectra are quite different in some models. Some predict a very weak \\knie\\ structure, like e.g. the acceleration in cannon balls by Plaga and the anomalous diffusion described by Lagutin et al, while others obtain a steep falling flux at the \\knie, an example is the diffusion model by Kalmykov and Pavlov. The results from indirect measurements seem to favor a relative steep cut-off, see section~\\ref{pgsect}.  Most models predict featureless spectra with a simple shape, however, some scenarios yield distinct features, which may have different theoretical origins.  Kalmykov and Pavlov as well as Swordy consider a diffusion model and a Leaky Box model, respectively and obtain bumps at the individual \\knie\\ positions. Consequently, they find a dip at \\knie\\ energies in the mean logarithmic mass.  The approaches by Sveshnikova (acceleration in supernova remnants) as well as V\\\"olk and Zirakashvili (reacceleration in the galactic halo) lead both to a twofold shape of the element spectra.  Yet, also common features can be recognized, like the relative sharp kinks in the spectra as found by Erlykin and Wolfendale as well as by Stanev et al., both models describe acceleration in supernova remnants. The distinct shape leads in both cases to structures in the $\\lna$ distribution.  The review has shown further that even similar conceptions lead to $\\lna$ values covering the whole range depicted in Figure~\\ref{lnamod}.  As an example, acceleration in supernova remnants yields $\\lna$ values at the lower edge of the experimental values (Berezhko and Ksenofontov) and also at the upper border (Kobayakawa et al.).  Unfortunately, the experiments give no conclusive energy spectra for elemental groups (Figures~\\ref{kascade} - \\ref{hegra}) and the mean logarithmic mass covers a relatively wide range, see Figure~\\ref{lnamod}.  Most astrophysical informations available from air shower measurements depend on the model used to describe the interactions in the atmosphere.  Hence, an interpretation and evaluation of the data should be treated with care and it is very challenging to characterize in a quantitative way the best model to describe the origin of the \\knie.  In models which introduce new type of interactions in the atmosphere, like the one just discussed by Kazanas and Nicolaidis or the approaches by Nikolsky et al. \\cite{nikolsky} or by Petrukhin \\cite{petrukhin}, the threshold for a new type of interaction depends on the energy per nucleon. Thus, such models ultimately arrive at a mass dependent cut-off for individual element spectra. The investigations of the KASCADE group however seem to indicate a rigidity dependent cut-off for individual elements \\cite{ulrich,ulrichisvh,rothisvh,roth,taup}.  Also with the phenomenological approach of the \\modell\\ a rigidity dependent approach is favored \\cite{polygonato}.  In addition, the simultaneous observations of a \\knie\\ in the electromagnetic, muonic and hadronic component by the KASCADE group \\cite{ulrich,emhknie,hknie} yield no major inconsistencies between the different air shower components. Therefore, a proposed new interaction would have to \"hide\" the energy in such a way, that no inconsistencies occur, when the measurements of different shower components are interpreted with standard particle physics. From the present experimental results new particle physics for high-energy interactions in the atmosphere seems not to be a likely explanation for the \\knie.  The interaction of cosmic rays with background particles like neutrinos or photons during their propagation causes fragmentation of heavy nuclei and the production of a large amount of secondary light nuclei, preferable protons. Models based on such an approach, like the models by Tkaczyk, Dova et al.\\footnote{Dova et al give their results only for a very limited range in primary energy, but above $10^7$~GeV a trend towards an extremely light mass composition can be recognized.}, and Candia et al. yield a very light mean logarithmic mass above \\knie\\ energies.  On the other hand, the measurements indicate an increase of $\\lna$ with energy, see Figure~\\ref{lnamod}, thus disfavoring interactions with background particles in the Galaxy as an explanation for the \\knie. In addition, the cross-sections for such interactions depend on the energy per nucleon. This yields a mass dependent cut-off for individual element spectra, which seems not to be favored by the measurements as mentioned above.  Numerous theories consider diffusive propagation of cosmic rays through the interstellar medium and leakage from the Galaxy as an explanation of the \\knie. The escape probability from the Galaxy in these models depends on the nuclear charge $Z$ and a rigidity dependent cut-off behavior is obtained for the individual element spectra. These models yield a more or less strong increase of the mean logarithmic mass above the \\knie\\ as consequence of the leakage starting with light nuclei.  The change of the spectral slope at the \\knie\\ is not very distinct in the Leaky Box model by Swordy as well as in the anomalous diffusion model by Lagutin et al. Consequently, only a modest increase of the mean logarithmic mass is obtained for energies above the \\knie, at the highest energies the predicted $\\lna$ values are at the lower edge of the measurements.  Propagation of cosmic rays in the Galaxy, taking into account contributions from the regular magnetic field, a random component and antisymmetric diffusion are described by Kalmykov and Pavlov, Ogio and Kakimoto, as well as Roulet et al. Although they consider similar scenarios, quite different energy spectra are obtained.  Ogio and Kakimoto derive a very smooth transition in the \\knie\\ region, while Kalmykov and Pavlov as well as Roulet et al. calculate a cut-off behavior similar to the measurements.  The mean logarithmic mass obtained for the three models covers a wide range,  and around $10^6$~GeV a significantly different behavior for the individual models is found.  On the other hand, they all result in similar values of $\\lna\\approx3$ around $10^{7.5}$~GeV, well compatible with the measurements.  Several models attribute the \\knie\\ to the acceleration process of cosmic rays. In the cannonball model described by Plaga, the spectral slopes change only by about $\\Delta\\gamma\\approx0.3$ at the \\knie\\ and protons are the dominant nuclei even above the \\knie. Hence, the resulting mean logarithmic mass is almost a constant function of energy. It is the weakest energy dependence of all models discussed.  In the single source model by Erlykin and Wolfendale heavy nuclei dominate above the \\knie\\ and the $\\lna$ values are in fair agreement with the measurements.  This can also be stated for the concept of V\\\"olk and Zirakashvili \\cite{voelk}, which allows for (re)acceleration of cosmic rays up to energies of the ankle.  Diffusive shock acceleration in supernova remnants is considered by Berezhko and Ksenofontov, Stanev et al., Kobayakawa et al., and Sveshnikova.  Different variations of the basic theory of Fermi acceleration have been applied.  Due to the finite lifetime of the shock front the maximum energy reached is limited. Typical values of $Z\\cdot (0.1 - 3)$~PeV are attained.  Kobayakawa et al. could increase the maximum energy reached in their calculations by the investigation of special magnetic field configurations.  Sveshnikova takes into account recent supernova observations, which indicate the existence of very energetic explosions (hypernovae). In their high-velocity shock fronts particles can effectively be accelerated to high energies.  The spectra obtained by the latter two scenarios are similar to the measurements and the \\modell.  The \\modell\\ is purely phenomenological. The spectra obtained are most compatible with the acceleration in supernova remnants as described by Sveshnikova, especially the heavy enriched case, and diffusive propagation through the Galaxy as discussed by Kalmykov and Pavlov as well as Roulet et al. Hence, these models could be a good theoretical foundation of the phenomenological ansatz.  Most models discussed represent principle ideas, dealing with individual aspects of cosmic-ray acceleration, propagation, or interactions.  Most probably, combinations of this basic processes are needed to describe nature realistically and more precise data will be necessary for a fine tuning of the free parameters in the models.  These parameters have been adjusted to different sets of experimental results in the individual models.  Most likely some of their predictions change, if a different set of experimental results is used.  At present, the experimental results do not allow to select a single approach as the best model to describe nature. The origin of the \\knie\\ in the all-particle energy spectrum is not yet explained.  However, within the boundaries given by the experiments, acceleration in supernova remnants and diffusive propagation through the Galaxy seem to be very attractive models to understand the origin of the \\knie."
    },
    "0402/astro-ph0402160_arXiv.txt": {
        "abstract": "We present a statistical study of substructure within a sample of $\\Lambda$CDM clusters and galaxies simulated with up to 25 million particles. With thousands of subhalos per object we can accurately measure their spatial clustering and velocity distribution functions and compare these with observational data. The substructure properties of galactic halos closely resembles those of galaxy clusters with a small scatter in the mass and circular velocity functions. The velocity distribution function is non-Maxwellian and flat topped with a negative kurtosis of about -0.7.  Within the virial radius the velocity bias $b=\\sigma_{\\rm sub}/\\sigma_{\\rm DM}\\sim 1.12 \\pm 0.04$, increasing to $b > 1.3$ within the halo centers. Slow subhalos are much less common, due to physical disruption by gravitational tides early in the merging history. This leads to a spatially anti-biased subhalo distribution that is well fitted by a cored isothermal.  Observations of cluster galaxies do not show such biases which we interpret as a limitation of pure dark matter simulations - we estimate that we are missing half of the halo population which has been destroyed by physical overmerging. High resolution hydrodynamical simulations are required to study these issues further. If CDM is correct then the cluster galaxies must survive the tidal field, perhaps due to baryonic inflow during elliptical galaxy formation. Spirals can never exist near the cluster centers and the elliptical galaxies there will have little remaining dark matter.  This implies that the morphology-density relation is set {\\it before} the cluster forms, rather than a subsequent transformation of disks to S0's by virtue of the cluster environment. ",
        "introduction": "Early simulation work that attempted to follow the merging hierarchy produced a final dark matter structure that was nearly entirely smooth (\\citealt*{White1976}; \\citealt*{White1987}; \\citealt*{Carlberg1994}; \\citealt*{Summers1995}; \\citealt*{Tormen1997}). The reason for this behaviour was debated in the literature as being due to physical or numerical overmerging (\\citealt*{White1978}; \\citealt*{Carlberg1994}; \\citealt*{vanKampen1995}; \\citealt*{Moore1996}). The development of fast algorithms to accurately integrate the orbits of millions of particles overcame this problem. The first halos simulated with sufficient resolution contained of the order a thousand substructure halos with properties that resembled galaxies within clusters \\citep*{Moore1998}. These simulations took many months using parallel gravity codes running on hundreds of processors.   Ongoing research in this area has given many interesting results and we list some of the main conclusions here (\\citealt*{Ghigna1998}; \\citealt*{Okamoto1999}; \\citealt*{Klypin1999a}; \\citealt*{Klypin1999b}, \\citealt*{Moore1999}; \\citealt*{Ghigna2000}; \\citealt*{Springel2001}; \\citealt*{DeLucia2004}): (i) Subhalos make up a fraction of between 5 and 10\\% of the mass of virialised halos. (ii) Halos on all mass scales have similar substructure populations. (iii) The mass and circular velocity function of subhalos are power laws with slopes -1 and -3. (iv) Velocity bias between the subhalos and smooth dark matter background may be significant. (v) The radial number density profile of subhalos is shallower than the dark matter background. (vi) Subhalos are significantly rounder than field halos (vii) The orbits of subhalos are close to isotropic with apo:peri approximately 4:1. (viii) Subhalos suffer mass loss from tidal stripping which modifies their outer density profiles. (ix) The tidal radii of subhalos decreases with cluster-centric position. (x) Most of the surviving population of subhalos entered the parent halo late.  Several of these statements remain controversial and further work is necessary to clarify certain issues. In this paper we re-address conclusions (i)-(v) and attempt to answer some of the remaining questions, including: What is the scatter in the mass and circular velocity distributions? Is there a positive or negative velocity bias and if so what is its origin? \\citet*{Ghigna2000}; \\citet*{Colin2000} claim a positive velocity bias whilst \\citet*{Springel2001} report a negative velocity bias. Have we converged in the properties of subhalos, including their radial distribution and mass functions? The inner regions ($r<$ 0.2 $r_{virial}$) of clusters and galaxy dark matter simulations are nearly smooth but is numerical overmerging still occurring in these very high density regions? Does the spatial distribution of galaxies in clusters resemble that of the subhalos in simulations? On galaxy scales the observed distribution of satellites is more concentrated than the simulations. Theory can be reconciled with the observations if it is assumed that the visible satellites are a biased subset of the total population (\\citealt*{Taylor2003}; \\citealt*{Kravtsov2004}). On cluster scales we do not expect to find ``dark galaxy halos'' therefore it is interesting to compare the observed distribution of galaxies with the distribution of substructure.  In this paper we analyse a sample of six high resolution simulations of clusters containing between 5 and 25 million particles integrated with high force accuracy. We compare the mass functions with a sample of galactic mass halos with slightly lower resolution. These new simulations are presented in Section 2 and the general properties of the subhalos are given in Section 3. ",
        "conclusions": "We analyse the substructure within six very high resolution cold dark matter simulations of galaxy clusters and four simulations of galaxies. We have addressed several open issues raised in the introduction regarding the results of high resolution simulations of individual halos within the concordance CDM model. Our conclusions can be summarised as follows:  \\begin{enumerate} \\item The spatial distribution of subhalos in cold dark matter simulations of galaxies and clusters is antibiased with respect to the mass. Although this behavior was found by other groups, we demonstrate that this result is robust and does not change as we increase the resolution. We show that this antibias most likely results from a population of early halos that are tidally destroyed in the dense protocluster environment and within the central regions of the final cluster. \\item The surviving population of subhalos have a positive velocity bias that increases towards the center of the halos. The subhalo velocity distribution functions are non-Gaussian, they are 'flat topped', especially in the inner region: There the kurtosis is $k=-0.7$ and the fourth Gauss-Hermite coefficient $h_4 = -0.068$. \\item The spatial anti-bias and the positive velocity bias of the subhalos are consistent with a steady-state solution of the Jeans equation. Subhalos are a hot, more extended component in equilibrium with the potential generated by the smooth particle background. \\item The mass and circular velocity distributions of subhalos in our highest resolution simulation show the same power law slopes as in lower resolution versions, but are steeper at the low mass end. It is not clear that convergence in the number of subhalos has been reached below a scale of a few hundred particle masses. \\item Cluster and galaxy mass halos simulated at the same resolution have similar substructure abundances. The scatter in the circular velocity and mass functions is a factor of three at the high mass end, but falls to just 1.7 at lower masses. \\item An observational comparison with CNOC cluster data and the Coma cluster shows that the galaxy population traces the smooth dark matter background, but not the predicted halo population. This is most likely due to overmerging in the central region of the simulations and we are probably missing a factor of two in the subhalo population. The baryonic cores of these disrupted subhalos may survive intact if dissipational processes increase their densities sufficiently. Also a greatly truncated dark matter halo may survive in this case.   \\end{enumerate}  This latter statement is the most profound conclusion of this work. The spatial distribution of cluster galaxies is significantly different from the distribution of subhalos in dark matter simulations. Either the model is incorrect or we have reached a fundamental limit to this type of pure dark matter simulation. Here we explore the latter possibility and the implications for the morphology density relation.  It is likely that disk galaxies do not significantly modify the overall potential provided by the baryons and dark matter. Whereas a disk-disk merger would funnel gas to the central region, forming an elliptical galaxy with a significantly deeper potential and a effective rotation curve that is at least isothermal, or possibly Keplerian in the center \\citep{Rom2003}. Thus we expect that an elliptical galaxy would most likely survive at any position within the cluster, albeit with a greatly truncated dark matter halo. Late type spiral galaxies are unlikely to survive within the central regions of clusters (or their progenitors) and will become physically overmerged to form the cD halo of diffuse light.  If the CDM paradigm is correct then we are missing close to a factor of two of the `galaxy' population as associated with subhalos, increasing to a factor of five within the inner 10\\% of the cluster. It is possible that simulations with more than $10^9$ particles per system may resolve more central subhalos and calculations this large will be possible in the future. In this case, the velocity bias should decrease as we resolve more halos/galaxies in the central regions.  However, from our convergence study we find very few new halos in the central cluster regions as we increase the resolution by a factor of ten. This implies that we have reached a physical limit to DM-only simulations and that any loss of subhalos in current simulations is due to physical overmerging (\\citealt{White1978}; \\citealt{Moore1996}). In this case progress in this area can only be made by including a realistic treatment of hydrodynamics and star-formation such that realistic disks and elliptical galaxies can be followed within the appropriate cosmological context.  The survival or disruption of a galaxy depends on an intricate balance between the progenitors dark halo structure and the effects of dissipation. Sa-Sb galaxies must lie on the borderline between survival and disruption in the cluster environment. The morphology-density relation may simply reflect the fact that the disks are preferentially destroyed in the central regions of clusters. However if the CDM model is correct one needs to preferentially form ellipticals in high density regions before the cluster forms. The fact that the observed galaxy distribution follows the dark matter distribution implies that no overmerging of galaxies has taken place. It is insufficient to take disks and destroy them in the cluster cores since this would give rise to a cored galaxy distribution.  The fact that 40\\% of halos identified at z=7 can not be associated with a subhalo at z=0, or have not merged with the central cD, implies that they have merged into the smooth particle background. If these objects can be associated with surviving galaxies, it implies a strong age-radius dependence for galaxies within clusters. At the cluster centres over 80\\% of the galaxies must have formed prior to z=7."
    },
    "0402/astro-ph0402483_arXiv.txt": {
        "abstract": "{We present NOT optical observations of a clump (l = 127\\fdg9435 , b = $+$1\\fdg8298), embedded in an extended, irregularly shaped, diffuse optical nebula. This condensation shows an emission-line spectrum typical of classic \\ion{H}{ii} regions. Although its location on the sky coincides with a nearby extended photoionized region recently identified by \\citet{ci03} in radio data from the Canadian Galactic Plane Survey (CGPS), the systemic velocity of this $\\approx$ 1\\arcmin-sized \\ion{H}{ii} region, V$_{\\rm {LSR}}$ = $-$71$\\pm$12 km s$^{-1}$, poses it far out in the Galaxy, beyond the Perseus arm. The location of this region in the Galaxy is supported by \\ion{H}{i} structures visible at comparable radial velocity on CGPS data. We argue that this \\ion{H}{ii} region might belong to an outer Galactic arm. The emission line ratios of the surrounding extended nebula, whose radial velocity is consistent with that of the small \\ion{H}{ii} region, are typical of photoionized gas in the low density limit. Smaller clumps of comparable surface brightness are visible within the optical boundaries of the extended, faint nebula. After comparison of the optical data with far infrared and radio observations, we conclude that this nebula is an \\ion{H}{ii} region, $\\sim$ 70 pc in size, probably photoionized by an association of OB stars and surrounded by a ring of neutral hydrogen. ",
        "introduction": "Galactic \\ion{H}{ii} regions offer the opportunity of investigating and understanding the details of star formation processes. Furthermore, the identification of \\ion{H}{ii} regions is of interest to trace the spiral structure of our Galaxy, which is still a matter of debate \\citep[e.g.][ and references therein]{mcg04,ru03}.  During our searches for galaxies beyond the Galactic plane -- based on visual inspection of Palomar Observatory Sky Survey plates of both first and second generation \\citep[e.g., ][]{ssw96,wgz99} -- we have identified an extended and irregularly shaped nebula, not previously catalogued, within (or in projection onto) which a number of brighter condensations are visible. The whole object is extended over about half a degree on the sky, while the bright condensations are much more limited in size, the brightest and largest one, located at l = 127\\fdg9435 , b = 1\\fdg8298, having a major axis of $\\approx$ 1\\arcmin. Its morphology happens to be similar to that of the giant ($\\sim$ 2\\degr$\\times$2\\degr\\ in extent) \\ion{H}{ii} region \\object{IC 1805} (also known as W4), ionized by the association \\object{Cas OB6} and located at a distance of 2.35 kpc \\citep{mjd95}.   Two IRAS peaks of emission, listed in the IRAS Point Source Catalogue as \\object{IRAS 01342+6358} and \\object{IRAS 01330+6351}, are present in the same region of the sky and were identified and measured as actually extended FIR sources by \\citet{ft96}. Observations of one of these sources, \\object{IRAS 01342+6358}, during the CS(2-1) survey of ultracompact (UC) \\ion{H}{ii} regions \\citep{bnm96} left the source undetected, therefore its radial velocity and kinematic distance could not be determined. Recently, \\citet{ci03} gave evidence that the extended FIR emission is associated with an extended thermal source of radio emission. They argued that this structure is physically related to a ``hole'' in the \\ion{H}{i} distribution observed at a systemic velocity of $\\sim$ $-$29.5 km s$^{-1}$, and traces gas photoionized by the nearby O star \\object{HD 10125}. However, they do not provide any direct measurement of the radial velocity (and therefore distance) of the ionized material. To our knowledge, no spectroscopic observations of this ionized region have been published to date. In this paper we present for the first time $R$-band photometry and optical spectroscopy of a portion of the optical extended nebula spatially coincident with the extended FIR and radio emission, including the $\\sim$ 1\\arcmin-size brighter condensation cited above. Beside investigating the physical properties of this ``condensation'', that is revealed as a classic \\ion{H}{ii} region, we derive its radial velocity and try to establish whether this region is related to the surrounding extended nebula and associated to the extended photoionized region recently identified by \\citet{ci03}. ",
        "conclusions": "We have performed $R$-band photometry and optical spectroscopy of an emission nebula located $\\sim$ 2\\degr\\ above the Galactic plane with the main aims to examine the physical conditions of its gas and establish its location within the Galactic structure. It was also our intention to establish whether the observed optical structures in this region of the sky are physically connected one to the others or are rather distinct structures seen in projection along the line-of-sight.  Our data indicate that the bright condensation $\\sim$ 1\\arcmin\\ in size has typical emission-line ratios of a classic \\ion{H}{ii} region. The radial velocity of the observed H$\\alpha$ line suggests that the \\ion{H}{ii} region is located far out in the Galaxy, well beyond the Perseus arm. We speculate that it might belong to the outer arm of the Galaxy. Actually, indications of the presence of spiral arm tracers in the second Galactic quadrant at distances up to 11 kpc have been reported by several authors \\citep[see, e.g.][ and references therein]{kw89}. In particular, \\citet{kw89}, collecting a number of different spiral arm tracers, gave evidence of the existence of an outer spiral arm, bending above the Galactic plane in the second quadrant \\citep[see also the structure of the Galactic warp shown e.g. by ][]{w01}, at distances comparable to the kinematic distance of the \\ion{H}{ii} region discussed in this work.  The presence of a depression in the \\ion{H}{i} profile along the line-of-sight to the \\ion{H}{ii} region at velocities in agreement with those measured from our optical spectra further supports our interpretation. Additionally, we observe fainter emission lines also from the extended optical nebula that is completely filling our field-of-view. The radial velocity of these emission-lines is in agreement  with that of the identified \\ion{H}{ii} region, within the measurement errors. Once again the diagnostic emission-line ratios are typical of classic \\ion{H}{ii} regions. These facts support the idea that the bright clump and the irregularly shaped, extended optical nebula are physically related and constitute a center-brightened \\ion{H}{ii} region. Considerations on the emission measure, apparent size, and estimated distance of the nebula lead to conclude that its gas is ionized by an association of a small number of OB stars. A comparison between the DSS2 image of this extended nebula and the 60 and 100 $\\mu$m IRAS images reveals a striking spatial coincidence of the structures and, even more strikingly, its optical substructures are well reproduced by substructures visible in the 1420 MHz continuum. This is in contrast with the recent work of \\citet{ci03}, who give some evidence of the association of these FIR sources and their radio-continuum counterparts with a nearby (distance $\\sim$ 3 kpc) region photoionized by the O star \\object{HD 10125}. Although the arguments given by \\citet{ci03} look convincing, they are not supported by direct radial velocity measurements of the ionized gas. Indeed, the correspondence we find between a ``hole'' in the \\ion{H}{i} distribution at velocities around $-$67 km s$^{-1}$ and the extended radio-continuum emission appears much tighter than the correspondence with an \\ion{H}{i} ``hole'' at $-$29.5 km s$^{-1}$ suggested by \\citet{ci03}. In the light of our observations, we argue that both the FIR sources and the radio-continuum emission are rather associated to the more distant ($d \\sim$ 8.2 kpc) photoionized nebula we have identified."
    },
    "0402/astro-ph0402497_arXiv.txt": {
        "abstract": "Large-amplitude, high-luminosity soft X-ray flares were detected by the \\ro\\ All-Sky Survey in several galaxies with no evidence of Seyfert activity in their ground-based optical spectra.  These flares had the properties predicted for a tidal disruption of a star by a central supermassive black hole. We report \\chandra\\ observations of three of these galaxies taken a decade after their flares that reveal weak nuclear X-ray sources that are from 240 to 6000 times fainter than their luminosities at peak, supporting the theory that these were special events and not ongoing active galactic nucleus (AGN) variability. The decline of \\rxjb\\ by a factor of 6000 is consistent with the $(t-t_D)^{-5/3}$ decay predicted for the fall-back phase of a tidal disruption event, but only if \\ro\\ was lucky enough to catch the event exactly at its peak in 1990 October. \\rxja\\ has declined by a factor of 240, also consistent with $(t-t_D)^{-5/3}$. In the \\ion{H}{2} galaxy \\ngc\\ we find only resolved, soft X-ray emission that is undoubtedly associated with starburst activity. When accounting for the starburst component, the \\ro\\ observations of \\ngc, as well as the \\chandra\\ upper limit on its nuclear flux, are consistent with a $(t-t_D)^{-5/3}$ decay by at least a factor of 1000. Although we found weak Seyfert~2 emission lines in {\\it Hubble Space Telescope} spectra of \\ngc, indicating that a low-luminosity AGN was present prior to the X-ray flare, we favor a tidal disruption explanation for the flare itself. ",
        "introduction": "\\label{intsec}  Dormant, supermassive black holes, suspected to be present in the centers of many normal galaxies, should reveal themselves by a UV/X-ray flare when they tidally disrupt a star and some fraction of the stellar debris is accreted.  Tidal disruption flares were proposed by \\cite{li79} and Rees (1988, 1990) as a probe for supermassive black holes in the centers of inactive galaxies.  As argued by \\cite{ul99}, the spectrum of such a flare will be characterized by the blackbody temperature of a thick disk or spherical envelope at the tidal radius, $T_{\\rm eff} \\approx (L_{\\rm Edd}/4\\pi\\sigma R_T^2)^{1/4} = 3.7 \\times 10^5 M_8^{1/12}$~K for a solar-type star. The flare would begin when the most tightly bound portion of the tidal debris returns to the pericenter of the star's orbit and accretes onto the black hole. This first return would occur at a time $t_0$ following the disruption at $t_D$, where $(t_0 - t_D) \\sim 1.1 M_8^{1/2}$ yr. This estimate applies to a non-rotating star; in the likely case that the star is spun up to near break-up before disruption, $t_0 - t_D$ is reduced by the factor $3^{-3/2}$ (e.g., Li, Narayan, \\& Menou 2002). The maximum return rate of debris according to numerical simulations (Evans \\& Kochanek 1989) is $\\dot M_{\\rm max} \\sim 0.14\\,M_{8}^{1/2}\\,M_{\\sun}$ yr$^{-1}$ and occurs at a time $(t_{\\rm max} - t_D) \\sim 1.5\\,(t_0 - t_D)$. After the peak of the flare, material returns at the declining rate $\\dot M(t) = 0.3\\,M_8\\,[(t-t_D)/(t_0-t_D)]^{-5/3}\\,M_{\\sun}$ yr$^{-1}$, which is an important factor that controls the decay of the luminosity over the next few years. For $M_{\\rm BH} < 10^{7}\\ M_{\\sun}$, the maximum return rate is super-Eddington, resulting in a flare with $L_{\\rm flare} \\geq \\eta \\dot M_{\\rm Edd} c^{2} > 1.3 \\times 10^{45} M_7$ ergs s$^{-1}$. Thus, for low-mass central black holes ($M_{\\rm BH} < 10^7 M_{\\sun}$), tidal disruption theory predicts luminous flares of up to $10^{45}$ ergs s$^{-1}$, peaking in the soft X-ray domain, with time scales on the order of months. For $M_{\\rm BH} > 2 \\times 10^{7}\\ M_{\\sun}$, stellar debris takes a longer time to fall back than the time scale on which it can circularize and radiate; in this case accretion probably proceeds through a thin disk \\citep{ul99}.  The \\ro\\ All-Sky Survey (RASS; Voges et al. 1999) conducted in 1990--1991 was an ideal experiment to detect these flares since it sampled hundreds of thousands of galaxies in the soft X-ray band. \\ro\\ detected soft X-ray outbursts from several galaxies with no previous evidence of Seyfert activity (see Komossa 2002 for a review).   From the statistics of the RASS, \\cite{do02} calculated a rate of $\\approx 1 \\times 10^{-5}$ yr$^{-1}$ for X-ray flares from these non-active galactic nucleus (non-AGN) galaxies, consistent with expected rates \\citep{ma99,sy99,wm03}. In order to test in an independent way whether these flares were in fact tidal disruption events, as opposed to some other form of extreme AGN variability, \\cite{ge03} obtained optical spectra of three of these non-AGN galaxies with the {\\it Hubble Space Telescope} (\\hst ) through narrow slits to search for persistent Seyfert activity. These spectra were up to a factor of 100 more sensitive to nuclear activity than previously obtained ground-based data. Two of the galaxies, \\rxja\\ \\citep{ko99} and \\rxjb\\ (Grupe, Thomas, \\& Leighly 1999), showed no evidence of emission lines or a non-stellar continuum in their \\hst\\ nuclear spectra, consistent with their ground-based classification as inactive.  On the other hand, \\ngc\\ \\citep{ba96}, a starburst galaxy with strong emission lines, was found by \\cite{ge03} to have in its inner $0.\\!^{\\prime\\prime}1$ a nucleus with narrow emission line ratios indicating a Seyfert~2 classification. This weak Seyfert~2 nucleus requires a low level of prior non-stellar photoionization powered by accretion, which raises some doubt about whether its X-ray flare must have been a tidal disruption event but does not rule it out. In this paper, we report on follow-up X-ray observations with \\chandra\\ of \\ngc, \\rxja, and \\rxjb.  The superb spatial resolution of \\chandra\\ enables an even more sensitive search for nuclear X-ray activity in these now very weak X-ray sources, and resolves any non-nuclear sources of X-ray emission, both of which are needed to test more rigorously the tidal disruption hypothesis. ",
        "conclusions": "\\label{dissec}  \\chandra\\ observations of \\ngc, \\rxja, and \\rxjb\\ in 2001 and 2002 show that their X-ray fluxes are continuing to decline at a rate that is consistent with the predicted accretion rate as a function of time in the fall-back phase of a tidal disruption event. \\ngc\\ and \\rxjb\\ are observered to be factors of 1000 and 6000 fainter, respectively, than their peak luminosities detected in the RASS.   Only \\rxja\\ still has an identifiable nuclear X-ray source, whose luminosity is nevertheless a factor of 240 less than its peak in 1992.  Since \\rxja\\ and \\rxjb\\ were confirmed to be inactive galaxies from the absence of broad or narrow emission lines, or non-stellar continuum in their \\hst\\ spectra, the most natural interpretation of their presently weak X-ray emission is the continuing decline following the tidal disruption of a star by an otherwise dormant central supermassive black hole.  The relatively low peak luminosity and fluence of the X-ray flare in \\ngc\\ is most simply interpreted as the accretion of only $\\sim 10^{-4}\\,M_{\\sun}$ of stellar debris, while the short duration of the peak is consistent with $M_{\\rm BH} \\leq 1 \\times 10^8\\,M_{\\sun}$.  The much higher luminosities of the flares in \\rxja\\ and \\rxjb\\ do not necessarily imply larger black hole masses for those objects, since a super-Eddington infall rate onto a smaller black hole can last for of order a year and maintain a constant luminosity for that time.  A longer peak duration would make it easier to understand why the RASS caught all of these events at their maximum flux.  \\cite{wm03} argued from theory that the rate of tidal disruption in galaxy nuclei should actually increase with decreasing $M_{\\rm BH}$; if so, this should be considered seriously as a selection effect.  In the case of \\ngc, the \\hst\\ detection of weak, Seyfert~2 emission lines in its nucleus raises additional uncertainty about whether its X-ray flare was the result of a tidal disruption or just exceptionally high-amplitude variability of its low-luminosity Seyfert nucleus that is not yet explained. The ability of \\chandra\\ to resolve the starburst source in \\ngc\\ enabled a further order-of-magnitude decrease in the X-ray luminosity of the nucleus to be detected with respect to the faintest \\ro\\ measurement.  In fact, there was no definite nuclear X-ray source to be seen in 2002, and the extreme amplitude of the decline leads us to favor the tidal disruption interpretation of the flare even if a low-luminosity AGN was present previously. It is likely that a prior AGN accretion disk survived the tidal disruption event; deeper exposures with \\chandra\\ should be able to detect its continuing X-ray luminosity as an excess above the extrapolated $t^{-5/3}$ decay.  Occasional reobservation of these objects is needed, if only to allow the tidal disruption hypothesis to be falsified by detecting renewed activity. While some of the tidal debris should itself spread into a thin accretion disk on the viscous time-scale, weak emission from this eventual AGN fuel source is not expected to dominate for several thousand years \\citep{li02}.  In the likely event that all three X-ray sources studied here continue to decline, deeper \\chandra\\ observations are required to track their luminosities and provide further observational constraints on the decay curve that the theory of the fall-back phase predicts."
    },
    "0402/astro-ph0402174_arXiv.txt": {
        "abstract": "A sample of close major-merger pairs (projected separation ${\\rm 5 \\leq r \\leq 20 h^{-1}}$ kpc, ${\\rm K_s}$ band magnitude difference $\\delta {\\rm K_s} \\leq 1$ mag) is selected from the matched 2MASS-2dFGRS catalog of Cole et al. (2001). The pair primaries are brighter than ${\\rm K_s} = 12.5$ mag. After corrections for various biases, the comparison between counts in the paired galaxy sample and counts in the parent sample shows that for the local `M$_*$ galaxies' sampled by flux limited surveys, the fraction of galaxies in the close major-merger pairs is 1.70$\\pm 0.32\\%$.  Using 38 paired galaxies in the sample, a ${\\rm K_s}$ band luminosity function (LF) is calculated. This is the first unbiased LF for a sample of objectively defined interacting/merging galaxies in the local universe, while all previously determined LFs of paired galaxies are biased by mistreating paired galaxies as singles. A stellar mass function (MF) is translated from the LF. Compared to the LF/MF of 2MASS galaxies, a differential pair fraction function is derived. The results suggest a trend in the sense that less massive galaxies may have lower chance to be involved in close major-merger pairs than more massive galaxies. The algorithm presented in this paper can be easily applied to much larger samples of 2MASS galaxies with redshifts in near future. ",
        "introduction": "Galaxy-galaxy interactions/mergers are very important in the cosmic evolution of galaxies. In the hierarchical galaxy formation paradigm, galaxies and galactic structures are formed through merging of smaller galaxies/structures (Kauffmann et al. 1993; Cole et al. 2000).  There is strong evidence that galaxy-galaxy interactions/mergers can significantly enhance the star formation rate (SFR) in galaxies involved (see Kennicutt 1998 for a review).  There is also evidence that the strong evolution in the cosmic SFR is due to a population of peculiar/interacting starburst galaxies which are closely related to galaxy mergers (Brinchmann et al. 1998). Theoretical simulations (Barnes 1990) and observations (Schweizer 1978; Kormendy \\& Sanders 1992) show that gas-rich late-type galaxies transform to gas-poor early-type E/S0 galaxies through galaxy mergers. This process may also be responsible for the formation of bulges of disk galaxies. The AGN activity, another physical process linking the mass to the light, may also be closely related to galaxy-galaxy interactions/mergers if the central black holes in AGNs are built up mostly in galaxy mergers, given the tight correlation between the black hole mass and the bulge mass (Franceschini et al. 1999). During the merging, the tidal torque can send large amount of gas and stars deep into the galactic nuclear region, and therefore feed a pre-existing black hole very efficiently, leading to enhanced AGN activity. In summary, galaxy-galaxy interactions/mergers play a central role in four of the most important processes in galaxy evolution, including mass assembly, star formation, morphological transformation and AGN activity.  There are two classical methods to select interacting/merging galaxies. One is to identify binary galaxies. And the other is to find galaxies with peculiar morphology (e.g. with tidal tails, double nuclei, or distorted disc). The latter has the advantage of including both pre- and post-mergers, and the identifications need not to be confirmed by redshift data. Studies based on this method found strong evolution in the fraction of major-mergers (galaxy pairs with mass ratio $\\lsim 3$), particularly for massive galaxies (Brinchmann et al. 1998; Le F\\'evre et al. 2000; Conselice et al. 2003). However, these results may suffer uncertainties due to difficulties in quantifying the morphological peculiarity and in detecting the interaction signatures in high z galaxies because of poor spatial resolutions and the cosmic dimming. In contrast, it is easy to define binary galaxies quantitatively and objectively. This makes an objectively defined comparison between local merger events and high z merger events possible. However, earlier studies of pair fraction and its cosmic evolution have suffered seriously from the contamination of unphysical pairs because of the lack of redshifts or highly incomplete redshift data (Zepf \\& Koo 1989; Burkey et al 1994; Carlberg et al. 1994; Yee \\& Ellington 1995; Woods et al. 1995; Patton et al. 1997; Wu \\& Keel 1998). In two recent studies using samples of galaxies with measured redshifts, Le F\\'evre et al (2000) and Patton et al. (2002) found $m=2.7\\pm 0.6$ and $m=2.3\\pm 0.7$, respectively, where m is the evolution index in the evolution function of the merger rate: $R_{mg} \\propto (1+z)^m$.  In a series of papers, Patton et al. (1997; 2000; 2002) pointed out that in studies of merger rate and evolution, it is very important to correct various systematic biases, and results from comparisons between mismatched samples of low-z and high-z galaxies are not very meaningful.  The best way to constrain the merger rate and its cosmic evolution is to compare differential pair fraction functions at different redshifts. A differential pair fraction function (DPFF) is defined by ratios between number of paired galaxies and the total number of galaxies in luminosity (or stellar mass) bins. Such functions are not sensitive to sample selection (flux limited or volume limited) and therefore can be compared without bias between different studies. DPFF can be determined by comparing the luminosity (mass) function of paired galaxies with that of total galaxies.  In this paper, we estimate the local ${\\rm K_s}$ (2.16$\\mu m$) band luminosity function of close major-merger pairs, and derive from it the DPFF in the z=0 universe. The very close relation between the ${\\rm K_s}$ band luminosity and the stellar mass means that for the first time we can have the mass function of the paired galaxies and the mass dependence of the merger rate. Since this can be compared directly to the predictions of the hierarchical galaxy formation simulations (e.g. Benson et al. 2002), it will provide an important test for these simulations. Through out this paper, we adopt the $\\Lambda$-cosmology with $\\Omega_m=0.3$ and $\\Omega_\\Lambda = 0.7$, $h = {\\rm H}_0/(100\\; {\\rm km~sec}^{-1} {\\rm Mpc}^{-1})$. ",
        "conclusions": ""
    },
    "0402/astro-ph0402342_arXiv.txt": {
        "abstract": "Recent observations of the interactions between radio sources and the X-ray-emitting gas in cooling flows in the cores of clusters of galaxies are reviewed. The radio sources inflate bubbles in the X-ray gas, which then rise buoyantly outward in the clusters transporting energy to the intracluster medium (ICM). The bright rims of gas around the radio bubbles are cool, rather than hot, and do not show signs of being strongly shocked. Energy deposited into the ICM over the lifetime of a cluster through several outbursts of a radio source helps to account for at least some of the gas that is missing in cooling flows at low temperatures. ",
        "introduction": "\\label{blanton:intro}  The vast majority of cooling flow clusters contain powerful radio sources associated with central cD galaxies. Initial evidence of radio sources displacing, and evacuating cavities in, the X-ray-emitting intracluster medium (ICM) was found with {\\it ROSAT} observations of a few sources including Perseus (B\\\"ohringer et al.\\ 1993), Abell 4059 (Huang \\& Sarazin 1998), and Abell 2052 (Rizza et al.\\ 2000). Models predicted that the radio sources would shock the ICM, and that the X-ray emission surrounding the lobes would appear hot (Heinz, Reynolds, \\& Begelman 1998). High-resolution images from {\\it Chandra} have revealed many more cases of radio sources profoundly effecting the ICM by displacing it and creating X-ray deficient ``holes'' or ``bubbles.'' The {\\it Chandra} data allow us to study the physics of the interaction in much more detail (i.e. Hydra A, McNamara et al.\\ 2000; Perseus, Fabian et al.\\ 2000; Abell 2052, Blanton et al.\\ 2001; Abell 2597, McNamara et al.\\ 2001; Abell 496, Dupke \\& White, 2001; MKW 3s, Mazzotta et al.\\ 2002; RBS797, Schindler et al.\\ 2001; Abell 2199, Johnstone et al.\\ 2002; Abell 4059, Heinz et al.\\ 2002; Virgo, Young et al.\\ 2002; Centaurus, Sanders \\& Fabian 2002; Cygnus A, Smith et al.\\ 2002; Abell 478, Sun et al.\\ 2003).  A long-standing problem with cooling flow models has been that the mass of gas measured to be cooling from X-ray temperatures, based on surface-brightness and spectral studies with {\\it Einstein} and {\\it ROSAT} has not been detected in sufficient quantities at cooler temperatures. High-resolution spectroscopy with {\\it XMM-Newton} provided direct evidence that gas was cooling in these clusters, but very large masses of gas (hundreds of solar masses) were seemingly cooling over only a limited range of temperatures.  Emission lines such as Fe XVII expected from gas cooling below approximately 2 keV were not detected and in general, most of the ICM seems to cool to about one-half to one-third of the cluster ambient temperature (Peterson et al.\\ 2003). Several possible solutions have been proposed for the lack of cool X-ray gas seen in the new observations (Fabian et al.\\ 2001; Peterson et al.\\ 2001).  These include mixing, heating by central active-galactic nuclei (AGN), inhomogeneous abundances, and differential absorption. Heating of the gas by a central radio source has also been discussed recently by B\\\"ohringer et al.\\ (2002), Churazov et al.\\ (2002), Ruszkowski \\& Begelman (2002), Kaiser \\& Binney (2003), and Br\\\"uggen (2003), among others.  A common finding with the {\\it Chandra} data is X-ray deficient holes that correspond with radio emission from the lobes of a central AGN. These holes are typically surrounded by bright shells of dense, X-ray-emitting gas. One of the surprises of the {\\it Chandra} observations is that the X-ray-bright rims surrounding the radio sources observed in cooling flow clusters were found to be cooler, rather than hotter, than the neighboring cluster gas (i.e. Perseus, Schmidt et al.\\ 2002; Hydra A, Nulsen et al.\\ 2002; Abell 2052, Blanton et al.\\ 2003). The bright shells show no evidence of current strong shocks. Soker, Blanton, \\& Sarazin (2002) found that for the case of Abell 2052, the morphology was well-explained by weak shocks occurring in the past, and strong shocks were generally ruled out. To date, no case of current strong-shock heating of the ICM in a cooling flow cluster by an AGN has been observed.  Although strong-shock heating from the radio source as proposed by Heinz, Reynolds, \\& Begelman (1998) and Rizza et al.\\ (2000) may not be the explanation for the fate of the missing cool gas, energy input from a radio source in the form of weak shocks (e.g. Reynolds, Heinz, \\& Begelman 2002) and buoyantly rising bubbles of relativistic plasma (e.g. Churazov et al.\\ 2002) can still contribute to heating. ",
        "conclusions": "\\label{blanton:conclusions}  Central radio-emitting AGN strongly affect the X-ray-emitting gas in cooling flows.  The radio sources create cavities or ``bubbles'' in the X-ray gas, which, in turn, confines the radio sources. In all clusters observed so far, there is no evidence that the radio sources are strongly shocking the ICM.  The X-ray-bright shells surrounding the bubbles are cool, not hot.  Weak shocks may have occurred in the past, creating the dense shells.  The only evidence for strong-shock heating of a similar nature has been seen in radio-ISM interactions in galaxies, and there are only very few cases of this, so far.  The X-ray pressures derived from the shells surrounding the bubbles are approximately ten times higher than the radio equipartition pressures. There may be problems with some of the equipartition assumptions, or additional contributions to the pressure within the radio bubbles.  One possibility for this additional pressure source is very hot, diffuse, thermal gas.  The bubble interiors are less dense than their surroundings and therefore will rise buoyantly outward into the clusters.  Ghost cavities provide evidence that this process has occurred.  The buoyantly-rising bubbles transport energy and magnetic fields into the clusters and they can also entrain cool gas from the cluster centers outward.  The shell pressures and bubble volumes measured in the X-ray can be used to determine the total energy output of the radio sources.  A radio source repetition rate of $\\approx 10^8$ yr is derived from the buoyancy rise time of the ghost cavities. A rough comparison of the average energy output of radio sources and the luminosity of cooling gas shows that the radio sources can supply enough energy to offset the cooling in cooling flows, at least in some cases."
    },
    "0402/astro-ph0402668_arXiv.txt": {
        "abstract": "We review different models used for reactions involved in nuclear astrophysics. The reaction rate is defined for resonant as well as for non-resonant processes. For low-density nuclei, we describe the DWBA method, the potential model, the $R$-matrix method, and microscopic cluster models. The statistical model is developed for high-level densities. Details of calculations in the low- and high-density regimes are illustrated with new results concerning transfer reactions and level densities. ",
        "introduction": "\\label{sec:intro} Nuclear reactions are important for studying energy generation and nucleosynthesis in astrophysical processes \\cite{Cl83,RR88}. Especially in explosive burning scenarios, large reaction networks include mainly nuclei far from stability which have not been studied experimentally so far, and perhaps will never be accessible in the laboratory. This poses a challenge to both experimentalists and especially theorists, which have to predict reaction cross sections of more or less unknown nuclei. Another challenge is found in the comparatively low interaction energies which pose problems for predictions as well as for measurements already close to and at stability. Here we provide and overview of the current status of theory and details of the remaining challenges, illustrated by several examples of new calculations.  Theoretical models can be roughly classified in three categories \\cite{De03}: \\\\ $(i)$ Models involving adjustable parameters, such as the $R$-matrix \\cite{LT58} or the $K$-matrix \\cite{Hu72} methods; parameters are fitted to the available experimental data and the cross sections are extrapolated down to astrophysical energies. These fitting procedures of course require the knowledge of data, which are sometimes too scarce for a reliable extrapolation. \\\\ $(ii)$ ``Ab initio\" models, where the cross sections are determined from the wave functions of the system. The potential model \\cite{BD85}, the Distorted Wave Born Approximation (DWBA) \\cite{OS91}, and microscopic models \\cite{WT77,La94} are, in principle, independent of experimental data. More realistically, these models depend on some physical parameters, such as a nucleus-nucleus or a nucleon-nucleon interaction which can be reasonably determined from experiment only. The microscopic Generator Coordinate Method (GCM) provides a ``basic\" description of a $A$-nucleon system, since the whole information is obtained from a nucleon-nucleon interaction. Since this interaction is nearly the same for all light nuclei, the predictive power of the GCM is important. \\\\ $(iii)$ Models $(i)$ and $(ii)$ can be used for low level-density nuclei only. This condition is fulfilled in most of the reactions involving light nuclei ($A \\leq 20$). However when the level density near threshold is large (i.e. more than a few levels per MeV), statistical models, using averaged optical transmission coefficients, are more suitable (see Sec.\\ \\ref{sec:hi}).  In Sec.\\ \\ref{sec:rates}, we start with useful definitions commonly used in nuclear astrophysics, such as of the reaction rates which are the actual quantities used in reaction networks, and their relation to cross sections and the relevant energy range. We then discuss different approaches and reaction mechanisms for systems with low (Sec.\\ \\ref{sec:low}) and high level densities (Sec.\\ \\ref{sec:hi}). In nuclear astrophysics, the relevant projectile energy range often covers both the low and high density regimes. The difficulties concerning the prediction of reaction rates bridging the two regimes are briefly addressed in Sec.\\ \\ref{sec:trans}. Concluding remarks are presented in Sec.\\ 5. ",
        "conclusions": "The methods to predict nuclear reaction rates are well established. Approaches based on nucleon-nucleon interactions have the advantage that they are first-principle methods, automatically including all relevant effects such as antisymmetrization, correlations, Pauli blocking, etc. On the other hand, depending on the model, effective NN interactions might have to be adjusted for each case. Common to all microscopic approaches is that the involved numerics limits the application to few-nucleon systems. The other class of reaction theories, based on optical potentials derived either from microscopic structure considerations or from phenomenology, average over certain properties and can be viewed as approximations to the exact calculation. Their strength lies in the wide applicability and the fact that the neglection of certain effects can be compensated by slightly adjusted optical potentials (which is automatically ascertained when deriving phenomenological potentials from experiment).  However, any model can only be as good as the underlying nuclear structure model required to provide the necessary input for the reaction calculation. While the combination of experimental data and predicted nuclear structure leads to good accuracy in the theoretical rates at and close to stability, larger errors should be expected further off stability. Future radioactive ion beam facilities (RIA, GSI upgrade) will provide the possibility for improvement. The situation is also eased by the fact that nuclei far off stability are usually produced in reaction equilibria (such as the (n,$\\gamma$)--($\\gamma$,n) equilibrium of the $r$ process) where individual reactions are not important. Thus, reactions have to be explicitly included only closer to stability (hydrostatic and explosive burning, freeze-out from high-temperature scenarios). Due to this fact it is actually possible to produce meaningful astrophysical models. We can also expect the future facilities to directly study this relevant range of nuclei.  \\begin{ack} A part of this text presents research results of the Belgian program P5/07 on interuniversity attraction poles initiated by the Belgian-state Federal Services for Scientific, Technical and Cultural Affairs. This work was partially supported by the Swiss National Science Foundation (grant 2000-061031.02). T. R. acknowledges a PROFIL professorship from the Swiss NSF (2024-067428.01). \\end{ack}"
    },
    "0402/astro-ph0402032_arXiv.txt": {
        "abstract": "s{We introduce the concepts of adiabatic (curvature) and isocurvature (entropy) cosmological perturbations and present their relevance for parameter estimation from cosmic microwave background anisotropies data. We emphasize that, while present-day data are in excellent agreement with pure adiabaticity, subdominant isocurvature contributions cannot be ruled out. We discuss model independent constraints on the isocurvature contribution. Finally, we argue that the Planck satellite will be able to do precision cosmology even if the assumption of adiabaticity is relaxed.}  \\newcommand{\\ETAL}{{\\it et al.}} \\newcommand{\\EG}{e.g.~} \\newcommand{\\IE}{i.e.~} \\newcommand{\\CF}{cf}  \\newcommand{\\CRIT}{{\\rm crit}} \\newcommand{\\TOT}{{\\rm tot}} \\newcommand{\\MAT}{{\\rm m}} \\newcommand{\\CDM}{{\\rm c}} \\newcommand{\\BAR}{{\\rm b}} \\newcommand{\\FG}[1]{Fig.~#1}  \\newcommand{\\Om}{\\Omega} \\newcommand{\\La}{\\Lambda} \\newcommand{\\pert}[1]{\\frac{\\delta \\rho_{#1}}{\\rho_{#1}}}   \\newcommand{\\UUNIT}[2]{ {\\;\\mathrm{#1}^{#2}} }  \\newcommand{\\MIX}{{\\rm MIX}} \\newcommand{\\AD}{{\\rm AD}} \\newcommand{\\BI}{{\\rm BI}} \\newcommand{\\CI}{{\\rm CI}} \\newcommand{\\NID}{{\\rm NID}} \\newcommand{\\NIV}{{\\rm NIV}}   \\newcommand{\\be}{\\begin{equation}} \\newcommand{\\ee}{\\end{equation}} \\def\\gsim{\\raise 0.4ex\\hbox{$>$}\\kern -0.7em\\lower 0.62 ex\\hbox{$\\sim$}} \\def\\lsim{\\raise 0.4ex\\hbox{$<$}\\kern -0.8em\\lower 0.62 ex\\hbox{$\\sim$}} ",
        "introduction": "The cosmic microwave background (CMB) constitutes one of the pillars of modern cosmology. Since the first measurement of primary anisotropies by the COBE satellite\\cite{Bennett:96} in 1992, the steady increasing precision has culminated with the WMAP results\\cite{WMAP}.  The dependence of the CMB power spectrum on cosmological parameters is well understood, and fast numerical codes produce a theoretical power spectrum in a few seconds with better then percent accuracy. Therefore the extraction of cosmological parameter by grid or Monte Carlo techniques is now a well established practice. By ``cosmological parameters'' we mean a set of 5 dimensionless numbers which describe the matter content of the universe today ($\\Om_\\La, \\Om_c,\\Om_b, \\Om_r, \\Om_\\kappa$ parameterizing the energy density in terms of the cosmological constant, cold dark matter (CDM), baryons, radiation and curvature, respectively), supplemented with the value of the Hubble parameter today, $H_0$, and the optical depth to reionization, $\\tau$. In the  simplest scenario which contains scalar perturbations only, we need two other parameters to describe the amplitude ($A_s$) and spectral dependence ($n_s$) of the initial perturbation spectrum. Certain combination of those parameters constitute ``orthogonal sets'' with respect to the CMB data, hence can be determined with very high accuracy\\cite{Kosowsky:02}.  Combination of CMB data with other cosmological data sets allows to constrain the above 9 standard parameters within a few percent. This is a spectacular achievement, even more so since many totally independent measurements seem to be converging toward the same values\\cite{Tegmark:03}. The accuracy of parameter extraction relies on the assumption that the initial conditions (IC) for the perturbations are of the simplest possible type, namely purely adiabatic. It is much more difficult to use the CMB to constrain cosmological parameters and at the same time to learn more about possible deviations from adiabaticity in the IC. Nevertheless the CMB represents the most promising data set to learn  about the type of initial conditions realized in the observed Universe: it is our window to the very early universe. ",
        "conclusions": ""
    },
    "0402/astro-ph0402518_arXiv.txt": {
        "abstract": "We report on a 10\\,ksec observation of the double pulsar system J0737--3039 obtained with the {\\em Chandra\\/} X-ray Observatory's Advanced CCD Imaging Spectrometer. We detect the system as a point source with coordinates (J2000) $\\mbox{RA} = 07^{\\rm h}37^{\\rm m}51\\fs23$, $\\mbox{Decl.} = -30\\arcdeg39'40\\farcs9$, making this the first double neutron star system to be detected at X-ray energies. Only 77 source counts are detected.  The data are represented satisfactorily by a power-law spectrum with photon index $\\Gamma = 2.9 \\pm 0.4$.  The inferred 0.2--10\\,keV luminosity is $\\approx 2\\times10^{30}$\\,erg~s$^{-1}$ for a distance of 0.5\\,kpc. We do not detect any significant variability with orbital phase in these data, which span one full orbit. These results are consistent with the X-ray emission originating solely from the magnetosphere of the energetic A pulsar, with an efficiency for conversion of rotational energy to X-rays in this band of $\\sim$ 0.04\\%, although we cannot exclude other possibilities. ",
        "introduction": "\\label{sec:intro}  Thirty years following the discovery of the first double neutron star \\cite{ht75}, only eight such systems are known.  A few of these rare binaries are wonderful laboratories for the study of relativistic gravity and gravitational radiation.  One of the more recently discovered of these systems, J0737--3039 \\cite{burgay03}, is fundamentally different from all of the others observationally in that it contains {\\em two\\/} detected pulsars \\cite{lyne04}.  The first pulsar discovered in the system (A) spins every 22\\,ms, has a low inferred surface magnetic field strength, a characteristic age of 210\\,Myr, and a relatively high rate of rotational energy loss, $\\dot E = 5.8\\times10^{33}$\\,erg~s$^{-1}$. Pulsar B has a period of 2.7\\,s, a characteristic age of 50\\,Myr, and an $\\dot E$ that is 3600 times smaller than that of A.  Its mass, 1.25\\,M$_{\\odot}$, is the lowest known for any neutron star and is 0.09\\,M$_{\\odot}$ less than that of A.  Both pulsars orbit their common center of mass every 2.4\\,hr with eccentricity of 0.09. We observe the system nearly edge-on, with an  orbital inclination of $87\\arcdeg$.  A unique aspect of the system is that radio emission from pulsar B is only detected strongly in two repeatable orbital phase ranges, each of duration $\\sim 10$\\,min.  A $\\sim 30$\\,s-long eclipse of pulsar A is also observed at its superior conjunction.  Both of these effects most likely result from the interaction of the fast-spinning pulsar's relativistic wind with the magnetosphere of its much less energetic companion \\cite{lyne04}.  This system will be an exceptional laboratory for relativistic astrophysics and should provide unique insights into the magnetospheres and the local environment of both neutron stars.  High-energy observations, particularly at X-ray energies, could be especially valuable for understanding the energetics of the system.  Whether resulting from magnetospheric or thermal processes, most pulsars emit of order $10^{-3}\\dot{E}$ in soft X-rays \\cite{bt99}. This translates to an expected luminosity of order $10^{31}$\\,erg~s$^{-1}$ for pulsar A, and means that B should be undetectable due to its very low $\\dot E$ and relatively large age.  But in this particular system, while still unlikely, it is conceivable that a portion of A's relativistic Poynting/particle flux might heat the surface of B, perhaps causing it to emit detectable radiation.  Such emission could prove invaluable for probing the atmosphere of B, which is likely to deviate significantly from standard models.  A more likely source of ``unusual'' X-ray emission might be the shock that undoubtedly forms near pulsar B owing to the collision between A's relativistic wind and B's magnetosphere. Some of this emission could be time variable, for instance due to Doppler boosting.  Many of these emission modes are clearly speculative and hard to quantify with the present understanding of the system.  Nevertheless, given the potential for unusual insights into the physics of this unique interacting binary, we have obtained a short X-ray observation of this system, awarded through the {\\em Chandra\\/} Director's Discretionary Time program, in order to better plan future multi-wavelength studies. In this Letter we report the detection of a point source at the position of the J0737--3039 system and describe its properties.  \\clearpage ",
        "conclusions": "\\label{sec:discussion}  We have detected the J0737--3039 system as a faint X-ray source that is best described by a steep power-law spectrum with $\\Gamma \\approx 2.9$ and $L_x \\approx 2\\times10^{30}(d/0.5\\,\\mbox{kpc})^2$\\,erg~s$^{-1}$ (see Table~\\ref{tab:xrayprops}).  This result is compatible with the emission originating solely from the magnetosphere of pulsar A:  Becker \\& Tr{\\\" u}mper (1999) find that for recycled radio pulsars that are detected at X-ray energies, $L_x \\sim 10^{-3} \\dot E$, and that when their spectra can be fit by a power law, $2 \\lapp \\Gamma \\lapp 2.4$.  While by comparison our spectral index is a bit steep and our efficiency for conversion of rotational to X-ray energies is a little small ($4\\times10^{-4}[d/0.5\\,\\mbox{kpc}]^2$), uncertainties in spectral parameters and distance make this interpretation perfectly plausible.  We note in passing that the X-ray luminosity detected from the system is roughly equal to the entire spin-down luminosity of pulsar B.  Another intriguing possibility is that these X-rays are produced when the relativistic winds of A and B collide.  As shown by Arons \\& Tavani (1993), the synchrotron emission produced in such interactions would also be expected to have a power law X-ray spectrum with $\\Gamma \\sim 2$. Unfortunately, the data from this short observation are not sufficient to discriminate between these two possibilities.  More sensitive observations with better time resolution are necessary in order to: obtain a better spectral fit; determine what portion of the detected X-ray flux is pulsed at the rotation period of pulsar A; quantify better any possible time variability (the observed radio flux density of B varies dramatically on $\\sim 5$--10\\,min timescales that are impossible to investigate with the present data);  measure any thermal component that may be contributing to the total X-ray flux of this system."
    },
    "0402/astro-ph0402204_arXiv.txt": {
        "abstract": "We compare a very deep X--ray image of M87, at the center of the Virgo Cluster, to high-quality optical images of the low excitation emission-line gas in the same region. There are striking coincidences of detail between the two. We explore the possiblity that this represents a thermal interaction between hot gas at $10^7$~K and warm gas at $10^4$~K. We find two temperatures are present in the X--ray gas, with the lower more prevelant in the vicinity of the optical filaments. Electron conduction from the hot phase to the cooler one provides a quantitatively acceptable energy source for the optical filaments, and we show additionally that it can do so for the brightest X--ray cluster, Perseus. If operative, conduction in the presence of gas-rich galaxy mergers, may explain the presence of ``cool cores'' in clusters of galaxies. ",
        "introduction": "Recent observations with the new generation of X--ray satellites, XMM-Newton and Chandra, have shown that the conventional cooling-flow theory (Fabian 1994) is no longer tenable. Data from XMM-Newton has revealed that while there is relatively cool X--ray emitting gas in the centers of galaxy clusters, there is an absence of predicted cooler gas below $\\sim 10^7$~K (Molendi \\& Pizzolato 2001, B\\\"ohringer et~al.\\ 2002, Matsushita et~al.\\ 2002, Molendi 2002, Donahue \\& Voit 2003). Two mechanisms in particular have been invoked to compensate for energy lost through X--ray radiation, namely heating by an intermittently active nucleus (e.g.\\ Tucker \\& Rosner 1983, Reynolds et~al.\\ 2001, Ruszkowski \\& Begelman 2002, Kaiser \\& Binney 2003) and thermal electron conduction of heat into the inner regions from the vast reservoir of thermal energy in the cluster e.g.\\ Tucker \\& Rosner 1983, Rosner \\& Tucker 1989, Sparks, Macchetto, \\& Golombek 1989, Sparks 1992, Voigt et~al.\\ 2002, Fabian, Voigt, \\& Morris 2002, Zakamska \\& Narayan 2003).  The Virgo cluster of galaxies is the nearest cluster with substantial X--ray emission, hosting a small but typical cool core and the well-studied active radio galaxy M87. XMM-Newton spectroscopy of the core of the Virgo cluster was amongst the pivotal work that showed large amounts of cooling gas are not present. The lower limit to the temperature found by XMM-Newton was 0.8~keV, or $0.9\\times 10^7$~K (Molendi 2002). The proximity of M87, 16.1~Mpc (Tonry 2001), makes it an attractive target for detailed physical studies, with spatial resolution of 1~arcsec corresponding to 77~pc. Here, we describe a very deep Chandra X--ray image and show that it displays a wealth of fine structure on the arcsec scale. Comparison with optical images reveals that these filaments are closely related to the $H\\alpha+[NII]$ emitting warm gas at $T\\sim 10^4$~K. This in turn suggests interaction between two very different gas phases, and hence provides clues to the physical transport processes operating in the cluster center.  From an optical perspective, the emission filament system was studied by Sparks, Ford, \\& Kinney (1993), and earlier, by Ford \\& Butcher (1979). Sparks et~al.\\ (1993) found that the emission filaments are dusty, that the innermost ones are likely to be outflowing due to interaction with the jet, but that the outer filaments are infalling. The overall excitation state is low, typical of filament systems in elliptical galaxies and galaxy clusters. The kinematics and geometry inferred from the dustiness and correlation with radio source structures led Sparks et~al.\\ (1993) to propose that the filaments have their origin in a small merger event. Weil, Bland-Hawthorn \\& Malin (1997) found evidence of diffuse low surface brightness light in M87 at large radii to the Southeast, which they interpreted as tidal debris of a recent encounter.  Young, Wilson, \\& Mundell (2002) presented an analysis of a 40~ks Chandra X--ray image of the Virgo cluster. They showed that the central regions are cooler than the outer parts of the cluster, and that there are localized cool regions associated with the complex radio source of M87, with its jet, inner lobes, intermediate lobes, and outer, apparently buoyant, relic radio structure. Young et~al.\\ (2002) further showed that the X--ray structures to the North and East of the nucleus correlate with optical $H\\alpha+[NII]$ filaments, and that the X--ray gas was coolest in the vicinity of the optical filaments.  Here we present a $\\sim 150$~ks Chandra X--ray image, $\\sim 4\\times$ longer integration than the Young et~al.\\ (2002) data. With this improved $S/N$, we show a much clearer view of the X--ray filaments and the detail within them, and hence reveal with greater clarity the nature of the correlation between the optical and X--ray emission. We also utilize the spectral information of the new data to probe the physical conditions within the filaments compared to their surroundings. Motivated by an earlier suggestion that electron conduction from the hot gas into the cold gas of an accretion or merger event can explain the peculiar symptoms of a ``cooling flow'' (Sparks et~al.\\ 1989), we assess whether this mechanism could provide the energy to ionize the filament system. ",
        "conclusions": "Our primary goal has been to present a very deep X--ray image of the center of the Virgo cluster, and to stress an important empirical observation, namely that warm gas with $T \\sim 10^4$~K emitting strong optical line emission is strikingly similar to a phase which is a thousand times hotter and radiates at X--ray wavelengths. This correlation may be offering important clues to the nature of transport processes in the cluster interstellar medium, since it implies that two dramatically distinct phases, close in space, ``know about one another'' and are in some form of thermal communication. X--ray spectral analysis shows that the hot phase itself has two temperature components, and that the lower of these is more prevelant in the vicinity of the optical filaments.  We build on these observations to show, quantitatively, that electron conduction (possibly saturated) can provide the energy seen in the form of emission line radiation in these and other optical filaments. We do not claim that this is the {\\it only\\/} viable model. There are important uncertainties particularly with regard to magnetic field structures and strengths, and conduction would have to proceed at or close to Spitzer or saturated levels at least in one direction. Nevertheless, the model appears promising and offers an acceptable {\\it quantitative\\/} relationship between $10^7$~K coronal and $10^4$~K plasma properties. If the material of the filaments has an external merger origin, then other ancillary properties such as their dustiness have a natural explanation.  Given the dynamic and violent nature of galaxy cluster centers, influenced by mergers, active nuclei and energy transport processes, we suggest that the cool-core phenomenon is transient and merger induced."
    },
    "0402/astro-ph0402048_arXiv.txt": {
        "abstract": "{The emission line survey within the \\textbf{C}alar \\textbf{A}lto \\textbf{D}eep \\textbf{I}maging \\textbf{S}urvey (CADIS) detects galaxies with very low continuum brightness by using an imaging Fabry-Perot interferometer. With spectroscopic follow-up observations of $\\rm{M}_{\\rm{B}} \\ga -19$ CADIS galaxies using FORS2 at the VLT and DOLORES at TNG we obtained oxygen abundances of 5  galaxies at $z \\sim 0.4$ and 10  galaxies at $z \\sim 0.64$. Combining these measurements  with  published oxygen abundances of galaxies with $\\rm{M}_{\\rm{B}} \\la -19$ we find evidence that a metallicity-luminosity relation   exists at medium redshift, but it is displaced to lower abundances and higher luminosities compared to the  metallicity-luminosity relation in the local universe. Comparing the observed metallicities and luminosities of galaxies at  $z \\la  3$ with P\\'egase2 chemical evolution models we have found a favoured scenario in which the metallicity of  galaxies  increases  by a factor of $\\sim 2$ between  $z \\sim 0.7$ and today, and their luminosity decreases  by $\\sim 0.5-0.9$\\,mag. ",
        "introduction": "In the local universe, metallicity is well correlated with the absolute luminosity (stellar mass) of galaxies (Skillman et al. \\cite{skilm}, Zaritsky et al. \\cite{zarit}, Richer \\& McCall \\cite{richer}, Garnett et al. \\cite{garnet}, Hunter \\& Hoffman \\cite{hunter}, Melbourne \\& Salzer \\cite{melb}, among others) in the sense that more luminous galaxies tend to be more metal rich. This metallicity-luminosity relationship is often attributed to the action of galactic superwinds: massive (more luminous) galaxies reach higher metallicities because they have deeper gravitational potentials which are better able to retain their gas against the building thermal pressures from supernovae, whereas low-mass systems eject  large amounts of metal-enriched gas  by supernovae-driven winds before high metallicities are attained (e.g., MacLow \\& Ferrara \\cite{maclow}). However,  a greater degree of gas consumption in luminous galaxies and/or suppressed infall of low-metallicity or pristine gas may also play a major role (Pagel \\cite{pagel97}).   Optical emission lines from \\hii\\, regions have long been the primary mean of gas-phase chemical diagnosis in galaxies (see, e.g., the review by Shields  \\cite{shields}). The advent of large  telescopes and sensitive spectrographs has enabled direct measurement of the chemical properties in the ionized gas of cosmologically--distant galaxies with the same nebular analysis techniques used in local \\hii\\ regions. Distant galaxies subtend small angles on the sky, comparable to typical slit widths in a spectrograph. A typical ground-based resolution element of 1\\farcs0 corresponding to a linear size of $\\sim6$\\,kpc at $z=0.5$  encompasses the entire galaxy at medium redshift; so galaxy spectra tend to be integrated spectra. The use of {\\it spatially-integrated} emission line spectroscopy for studying the chemical properties of star-forming galaxies at earlier epochs has been explored by Kobulnicky et al. (\\cite{kob99}). They concluded that spatially-integrated  emission line spectra can reliably indicate the chemical properties of distant star-forming galaxies. They also found  that, given spectra with sufficient signal-to-noise, the oxygen abundance of these galaxies can be measured to within $\\pm 0.2$\\,dex using the \\R23\\ method, first suggested by Pagel et al. (\\cite{pagel}). The  \\R23\\ method allows to derive metallicities also for faint galaxies where only prominent emission lines are observable (see more details in Sect. 4).    At medium redshift, all previous attempts to study the evolution of metallicity with cosmic time (Kobulnicky \\& Zaritzky \\cite{kobzar}, Hammer et al. \\cite{hammer},  Carollo \\& Lilly \\cite{carol}, Contini et al. \\cite{contini}) have been based on continuum selected samples, the emission lines of which were identified using spectroscopy. Owing to magnitude-limited selection effects in these surveys, the samples of medium redshift galaxies   for which metallicities were determined are biased towards luminous galaxies. Thus, for instance, the CFRS sample used by Carollo \\& Lilly (2001),  selected by  $I <22.2$, contains only bright galaxies ($M_{B}<-20$) at redshift $0.6 \\la z \\la 1$, the continuum of which is dominated by an evolved stellar population.  Kobulnicky \\& Koo (\\cite{kobko}) found that star-forming Lyman-break galaxies at $z \\sim 3$ are 2$-$4\\,mag more luminous than local spiral galaxies of similar metallicity, and thus are offset from the local luminosity-metallicity relation. Also Pettini et al. (\\cite{pettini})  found that Lyman break galaxies at $z\\sim 3$ are significantly overluminous for their metallicities.  It seems therefore possible that the whole metallicity-luminosity relation is displaced to lower abundances at high or medium redshifts, ilustrating how galaxies participate in the chemical evolution process.         The \\textbf{C}alar \\textbf{A}lto \\textbf{D}eep \\textbf{I}maging \\textbf{S}urvey (CADIS, Meisenheimer et al. \\cite{meise98}, \\cite{meise03}) allows the selection of galaxies   by their emission line fluxes detected via narrow-band (Fabry-Perot) imaging, regardless of their continuum brightness. Thus, CADIS can reach galaxies with very low continuum brightness (and absolute luminosity), i.e.  galaxies which have  such faint continuum that they are only detectable by their emission lines. Combining metallicities for such faint CADIS galaxies in several narrow redshift bins with published abundances for more luminous galaxies at medium $z$  we aim to establish the metallicity-luminosity relation  for several look-back time bins in the range $0.3<z<1$. The metallicity-luminosity relation at different reshifts is a valuable  probe of chemical galaxy evolution, and  a consistency check for chemical evolution models.    This paper is structured  in the following way: In Sect. 2 we describe how we select  galaxies at $0.39 \\la z \\la 0.65$ from the CADIS emission line sample in order to determine their metallicities by follow-up spectroscopy. In Sect. 3 we present the spectroscopic follow-up of such emission line galaxies. In Sect. 4 we compute oxygen abundances of the spectroscopic sample. Finally, in Sect. 5  we discuss the evolution of the metallicity-luminosity relation with redshift. The ``concordance'' cosmology with  $\\rm{H}_{0}=70$ \\kmsmpc, $\\Omega_{0}=0.3$, $\\Omega_{\\Lambda}=0.7$ is used throughout this paper. For the solar oxygen abundance we use the value $12+\\rm{log}(\\rm{O/H})=8.87$ from Grevesse et al. (\\cite{grevesse}). Note that ``metallicity'' and ``abundance''  normally denotes ``oxygen abundance'' throughout this paper. ",
        "conclusions": ""
    },
    "0402/astro-ph0402081_arXiv.txt": {
        "abstract": "New X-ray observations from the {\\it Chandra} and XMM-{\\it Newton} observatories have shown that cooling of the intracluster medium is occurring at rates that are now approaching the star formation rates measured in cD galaxies at the bases of cooling flows.  Star formation proceeds in repeated episodes, possibly indicating an intermittent fuel supply. Coupled with new evidence for heating by AGN, a new paradigm of self-regulated cooling and star formation in cluster cores is emerging. ",
        "introduction": "\\label{McNamara:intro}  Little has changed in our understanding of star formation in cooling flows since the the topic was last reviewed at the meeting on cooling flows at Haifa in 1996. However, the nature of the so-called ``cooling flow problem'' has changed dramatically since then. This controversial problem concerns the deficit between the large cooling rates of the keV gas in the centers of clusters and the much smaller star formation rates observed in central cluster galaxies located at the bases of cooling flows (Fabian 1994).  While the cooling rates based on {\\it Einstein} and {\\it ROSAT} X-ray observations (but not ASCA, see Makishima this conference) were claimed to be tens to hundreds of solar masses per year, the star formation rates are generally between a few to several tens of solar masses per year.  In spite of recent detections of molecular gas in cooling flows (Edge 2001), this apparent violation of mass continuity cannot be reconciled by a repository of molecular gas clouds or star formation with the Local initial mass function in central cluster galaxies.  Adherents of the cooling flow paradigm dodged the problem by appealing to a repository of elusive matter, such as low mass stars or a mist of cold clouds.  Others proposed heating mechanisms (e.g., active nuclei, or AGN, heat conduction, cosmic rays, mergers, magnetic reconnection, etc.) with the potential to inject enough energy into the keV gas to balance radiative losses.   Nevertheless, these proposals generally suffered various problems, such as the need for fine tuning, but mostly they lacked observational support. ",
        "conclusions": "I have shown that the star formation rates estimated with optical, ultraviolet, and infrared observations are within factors of several of the new cooling limits, and in some cases, agree to within their errors.  What we know about the star formation histories in cooling flows points to repeated bursts of star formation lasting $10^7 -10^9$ yr. In some cases, the starbursts are triggered by the central radio sources. Constant star formation over the ages of clusters, the history of star formation predicted by steady cooling flow models, is inconsistent with most data.  The overall picture of star formation in cooling flows is consistent with, and indeed has helped to shape, the emerging paradigm of self-regulated cooling in clusters. The X-ray, optical, and radio data taken together point to episodes of cooling and periodic reheating by radio outbursts and their associated bubbles.  This scenario (e.g., Churazov et al. 2002) has received observational support from the newly-discovered trend between central X-ray luminosity and the instantaneous kinetic luminosity of bubbles in clusters (B\\^{\\i}rzan et al. 2004). Thermal conduction may be an additional element of the feed-back loop that regulates cooling (Ruszkowski \\& Begelman 2002, Nulsen 2003). This emerging picture of cooling flows has broad implications for theories of structure formation and evolution.  Feed-back processes may have been important during the early development of galactic bulges and their central black holes, and they may regulate the thermal balance of the hot gas in giant elliptical galaxies today (Nulsen, this conference).  The possibility that such processes are active in large, bright, relatively nearby clusters provides a unique opportunity to test these models in detail."
    },
    "0402/astro-ph0402612_arXiv.txt": {
        "abstract": "We present an evaluative study of some current methods utilized in the analysis of infrared (IR) observations of star-forming regions.  A series of self-consistent radiative transfer models are constructed, with the outputs analysed using these methods to infer source properties such as dust tempertaure, mass, opacity function, and density distribution.  Any discrepancies between the inferred and model quantities can be attributed to the analysis methods. The range of validity of most methods is smaller than expected, due to two effects:  (1) limited applicability of the Rayleigh-Jeans limit except to very long wavelengths, and (2) significant errors in the isothermal approximation, even when $\\Delta T(r) < 2\\mathrm{K}$ over 90 per cent of a region.  Still, an accurate mean $T_{\\mathrm{dust}}$ can be found using a modified Wiens law.  This temperature can yield dust masses to within 10-25 per cent -- much better than masses inferred from the integrated luminosity. Using long wavelengths ($> 1000-2000 \\mu \\mathrm{m}$), the opacity index can be determined from the far-IR spectrum to within 20 per cent. Fitting the spectrum yields better results.  The density distribution can be somewhat constrained by fitting the surface brightness, for well-resolved sources. Better results are found by fitting the flux spectrum with detailed models. ",
        "introduction": "Infrared (IR) radiation, due to thermal emission from dust grains, is a key feature of star-forming regions. These grains play an important role in the chemistry, thermal balance, and evolution of these regions. This role, coupled with the ability for IR radiation to escape from areas of high optical depth, makes the determination of source parameters through the study of IR observations of great importance.  The problem of obtaining source parameters from the IR observations is a difficult one.  As a result, it is important to study and evaluate the analysis methods used to judge their value and regimes of validity.  In general, the analysis of IR observations can be broadly broken into two groups.  The first approach relies upon the use of relatively simple, semi-analytic expressions which are derived from idealized and simplified approximations and assumptions (see e.g. Hildebrand 1983). It has been shown previously that the strict use of this approach may lead to interpretations which are ambiguous or erroneous (see e.g., Schmid-Bergk \\& Scholz 1976; Butner et al. 1991; Doty \\& Leung 1994; Men'shcikov \\& Henning 1997; Shirley, et al. 2000). The second approach involves the construction (usually a grid) of detailed, self-consistent radiative transfer models for the computation of model spectra and comparison with observations.  While such codes are available for use by the general community (e.g. CSDUST3 by Egan, Leung, \\& Spagna 1983; DUSTY by Ivezic, Nenkova, \\& Elitzur 1999), constructing such models requires substantial time and effort.  As a result, the first approach remains a common choice.  In this paper we critically evaluate semi-analytic methods of analysis of IR observations of star-forming regions. Here we consider internal and externally heated dust clouds as analogues of early-type star-forming regions. In particular, we examine the usual assumptions of homogeneities in density and temperature, as well as the neglect of opacity effects when applied to externally heated IR sources.  We do this by constructing a series of realistic models of these regions, treating the model output as simulated observations.  We then apply common semi-analytic methods to infer the source properties, and compare them to those used in constructing the models. The discrepancies between the input and inferred source properties yield a measure of the limitations of the semi-analytic methods used.  The outline of this paper is as follows.  In Section 2 we summarize the problem and current methods of analysis. We describe the models run and testing procedure in Section 3. In Section 4, we discuss the role and determination of the dust temperature.  We consider the dust mass in Section 5. In Section 6 we discuss the determination of the dust opacity function.  We consider the determination of the dust density distribution in Section 7. Finally, in Section 8, we summarize and draw conclusions from this study. ",
        "conclusions": "We have critically evaluated current semi-analytical methods in analysing regions of low-mass star-formation. In particular, we have investigated the validity of the underlying assumptions for the methods and determined the conditions under which these methods are reliable.  Based upon this work, we find that:  1.  The temperature distribution plays a significant role in determining the spectrum of a source.  As there are no current semi-analytic methods for determining the temperature distribution explicitly, a common assumption is that these sources are isothermal. Although an isothermal temperature with some physical basis can be found using a modified Wiens Law (Fig. 3 and equation [18]), using this temperature does not accurately determine the other source parameters (See Sections 4, 5, and 6).  2.  It is difficult to reliably determine the total dust mass without an accurate estimate of grain properties, density distribution, and the distance to the source (Section 5).  However, a ratio of mass from two similar sources can be determined using the observed flux at a given frequency. The dust masses determined by this method are accurate to within $50$ per cent for all cases where the estimated temperature is in reasonable agreement (5 K) with the average grain temperature -- namely for regions with no or weak internal sources. It is also possible to determine a ratio of dust masses using the integrated luminosity. This method is only applicable for low optical depths, as errors approach $100$ per cent for $\\tau_{0} \\ge 100$.  3.  For an optically thin source and in a region of frequencies where the Rayleigh Jeans approximation holds, the opacity index, $\\beta$, can be determined using the ratio of two different frequencies as in equation (12). In order to determine the opacity index to within $20$ per cent for all optical depths, the pair of wavelengths used need to be longer than about $2500\\mu \\mathrm{m}$, due to the failure of the RJA and the lack of isothermality of the source. This may be problematic due to the existence of free-free emission which may dominate the thermal dust continuum for some sources. Another method for determining the opacity index is by fitting the spectrum. With a single blackbody fit, $\\beta$ can be determined to within $60$ per cent for all cases. If, instead, a two temperature fit is used, the results for $\\beta$ are improved considerably, with a maximum error of only $7$ per cent.  4.  The density distribution index, $m$, can be determined for a resolved source that is isothermal along any line of sight by using the ratio of intensities at two different wavelengths.  For a theoretically small pencil beam, the results have an error of up to $100$ per cent. Errors increase as the beam size gets larger, growing dramatically for pointings in which the source does not fill the beam. The errors also increase somewhat as the wavelengths used decrease. Finally, due to the isothermal assumption, it is important to probe impact parameters that are neither too close to the edge of the source nor too close to the core. We find $0.03 r_{\\mathrm{out}} < p < 0.3 r_{\\mathrm{out}}$ works well in the cases studied here.  5.  It should be possible to infer the dust density distribution, $m$, from the surface brightness distribution across the source. This, however, requires good signal-to-noise data toward the source edge.  The range in surface brightness distributions is due to the differences in the actual temperature distributions in the sources, highlighting the need for self-consistent modeling to ensure proper interpretation of observational data."
    },
    "0402/hep-th0402225_arXiv.txt": {
        "abstract": "{Motivated in part by string theory, we consider the idea that the standard $\\Lambda$CDM cosmological model might be modified by the effect of a long-range scalar dark matter interaction. The variant of this widely-discussed notion considered here is suggested by the Brandenberger-Vafa \\cite{bv} picture for why we perceive three spatial dimensions. In this picture there may be at least two species of dark matter particles, with scalar ``charges'' such that the scalar interaction attracts particles with like sign and repels unlike signs. The net charge vanishes. Under this condition the evolution of the mass distribution in linear perturbation theory is the same as in the $\\Lambda$CDM cosmology, and both models therefore can equally well pass the available cosmological tests. The physics can be very different on small scales, however: if the scalar interaction has the strength suggested by simple versions of the string scenario, nonlinear mass concentrations are unstable against separation into charged halos with properties unlike the standard model prediction and possibly of observational interest.}  \\date{February, 2004} ",
        "introduction": "\\label{INTRODUCTION}  In this paper and its companion \\cite{gpTwo} we investigate models for a  long-range scalar-mediated interaction among dark matter particles.  Our starting assumptions are that the scalar in question is massless, or very nearly so, and that the scalar has negligible effect on physics in the visible sector---apart from what results from the gravitational interaction with the dark sector.  The fundamental theory that motivates our models offers reasons to doubt both assumptions, but the arguments are not convincing enough to discourage investigation of the possible observational significance of a string-inspired scalar field modification of the cold dark matter (CDM) cosmological model.  In addition to ideas from superstring theory, we are motivated by the thought that there surely is considerable room in the dark sector of cosmology for richer physics than the standard CDM model. Many current notions about viable examples of richer physics trace back to the long history of ideas about scalar-tensor gravity physics, elements of which are reviewed in a paper by G. Farrar and one of us \\cite{fp}.\\footnote{We take this opportunity to remedy an omission: the survey should have included the work on this subject by Frieman and Gradwohl \\cite{Frieman:1993fv}.} The convergence of these long-standing ideas about gravity physics with what superstring theory might be suggesting is appealing and just might signify that we are being led to a better approximation to cosmology.  The inspiration from string theory is the Brandenberger-Vafa \\cite{bv} picture for why we perceive three spatial dimensions. It suggests that the long-range interaction of dark matter is the sum of the gravitational force and the force from the exchange of massless scalars and gauge bosons. The dark matter particles may be assigned ``charges'' ---to be defined more precisely in section~\\ref{BV}---as in electrostatics, but particles with like signs of the scalar charge are attracted and unlike scalar signs are repelled. The model requires zero net scalar charge, which implies that in linear perturbation theory the evolution of the mass disttribution is unaffected by the scalar interaction. This means that initial conditions can be chosen so our model and the $\\Lambda$CDM cosmology equally well pass the cosmological tests, including the new tests based on the measurements  of the large-scale distributions of mass and the 3~K thermal background radiation \\cite{Bennettetal}. The scalar interaction introduces a new mode of instability, however, in which particles with different signs of charge tend to concentrate in different places. If there were a significant primeval amplitude in this charge separation mode it would tend to produce observable mass concentrations that are not closely correlated with the large-scale mass distribution, an observationally unacceptable situation. This is avoided if the initial conditions are very close to adiabatic, as we will assume. If the scalar force were comparable to the prediction of the simplest of string theory setups, then even with purely adiabatic initial conditions a strongly nonlinear dark matter mass concentration would be unstable against separating into massive charged halos, and halos of different signs would tend to avoid each other. We will suggest that such curious behavior may have something to do with the confusing situation in the theory and phenomenology of structure formation on the scale of galaxies. Other viable models for a modified CDM cosmology would have different implications for small-scale structure, of course. An example based on a supersymmetric picture for the dark matter, which allows a non-zero net scalar charge, is discussed in our companion  paper \\cite{gpTwo}.  We organize this paper as follows.  In section~\\ref{BV} we introduce the Brandenberger-Vafa scenario and describe its late-time dynamics in terms of an effective lagrangian for the dark matter.\\footnote{It should perhaps be emphasized that, for us, late time means the matter- and radiation-dominated epochs: this makes our investigation somewhat different in spirit from others in the brane gas literature.}  In section~\\ref{PERTURBATIONS}, we characterize the long-range forces in the dark sector by a matrix of dimensionless quantities which encode the physics of subsequent interest at scales comparable to the sizes of galaxies.  Using this matrix, we demonstrate in examples how the scalar force leads to a charge separation instability in the mass distribution. In section~\\ref{OBSERVATIONS}, we discuss possible observational consequences of this effect for structure formation on the scale of galaxies, and we comment on some issues of small-scale extragalactic structure that might point to physics along the lines of this model. These considerations are continued in our companion paper \\cite{gpTwo}.  In the last section, we return to the fundamental issues mentioned in our introductory paragraph: we consider why the scalar force does not manifest itself in the visible sector, in the E\\\"otv\\\"os experiment and measurements of orbits in the Solar System; we discuss the naturalness of assuming such a small mass for the scalar; and we outline some other issues of the fundamental physics of the model. ",
        "conclusions": ""
    },
    "0402/gr-qc0402059_arXiv.txt": {
        "abstract": "We examine in the context of general relativity the dynamics of a spatially flat Robertson-Walker universe filled with a classical minimally coupled scalar field $\\phi$ of exponential potential $V(\\phi)\\sim exp(\\!-\\mu\\phi)$ plus pressureless baryonic matter. This system is reduced to a first-order ordinary differential equation for $\\Omega_{\\phi}(w_{\\phi})$ or $q(w_{\\phi})$, providing direct evidence on the acceleration/deceleration properties of the system. As a consequence, for positive potentials, passage into acceleration not at late times is generically a feature of the system for any value of $\\mu$, even when the late-times attractors are decelerating. Furthermore, the structure formation bound, together with the constraints $\\Omega_{m0}\\approx 0.25-0.3$, $-1\\leq w_{\\phi 0}\\leq -0.6$ provide, independently of initial conditions and other parameters, the necessary condition $0<\\mu\\lesssim 1.6\\sqrt{8\\pi G_{N}}$, while the less conservative constraint $-1\\leq w_{\\phi}\\leq -0.93$ gives $0<\\mu\\lesssim 0.7\\sqrt{8\\pi G_{N}}$. Special solutions are found to possess intervals of acceleration. For the almost cosmological constant case $w_{\\phi}\\approx -1$, the general relation $\\Omega_{\\phi}(w_{\\phi})$ is obtained. The generic (non-linearized) late-times solution of the system in the plane $(w_{\\phi}, \\Omega_{\\phi})$ or $(w_{\\phi}, q)$ is also derived. ",
        "introduction": "Scalar-field cosmologies continuously receive through time special attention. The interest on scalar fields has been stimulated when it has been realized that a scalar field might be responsible for inflation of the very early universe. If its self-interaction potential energy density is sufficiently flat, it can violate the strong energy condition and drive the universe into an accelerated expansion. Even if the potential is too steep to drive inflation, such models still have important cosmological consequences. Recent observational evidence shows that the energy density of the universe today may be dominated by a homogeneous component with negative pressure (quintessence)~\\cite{perl}. Slowly rolling self-interacting scalar fields may provide in the late-times evolution of the universe a dynamical mechanism to achieve a small effective cosmological constant. \\par Models with a variety of self-interaction potentials have been studied and the cosmological conclusions are naturally strongly model dependent. There are attempts to base the choice of the potential on some fundamental dynamical considerations. A class of potentials that is commonly investigated and which arises in a number of physical situations has an exponential dependence on the scalar field. Such potentials are motivated by higher-dimensional supergravity, superstring theory, or M-theory  that have been compactified to an effective four-dimensional theory~\\cite{sugra}. Also, higher-order gravity theories are conformally equivalent to general relativity plus a scalar field with potential which asymptotically tends to an exponential form~\\cite{conformal}. \\par Since the first inflationary cosmological models were presented, most calculations have been performed using one of a variety of approximations. While these are usually sufficient and often unavoidable, there is considerable interest in scalar field cosmologies that can be solved exactly. Exponential potentials have been extensively considered in cosmology by various authors, basically in Robertson-Walker models with no other source besides the scalar field~\\cite{cosmology}. Solutions with anisotropic/inhomogeneous geometries were obtained in~\\cite{anisotropic}. When a perfect fluid is added, only a few solution are known (see references~\\cite{matter}-\\cite{heard} for relevant works). In the context of ekpyrotic or pre-big bang collapse, exponential potentials were analyzed in~\\cite{ekpyrotic}, while non-minimal couplings of scalar fields with exponential potentials have been considered in~\\cite{nonminimal}. Attention has also been directed to the possible effect of such a scalar field on the growth of large-scale structure in the universe~\\cite{structure}. Other recent studies on exponential potentials have appeared in~\\cite{recent}. \\par In the present paper, we examine the dynamics of a spatially flat Robertson-Walker universe filled with a classical minimally coupled scalar field of exponential potential plus pressureless baryonic matter (luminous and nonluminous) non-interacting with the field. Clearly, this model cannot be used from the beginning of the universe evolution, but only after decoupling of radiation and dust. Thus, we do not take into account inflation, creation of matter, nucleosynthesis, e.t.c. By performing changes of phase-space variables, we reduce the system of Einstein equations into a first order nonlinear differential equation, which can be expressed as an equation for $\\Omega_{\\phi}(w_{\\phi})$ or $q(w_{\\phi})$. Solving this equation leads to integration of the whole system. Furthermore, studying this with respect to its acceleration properties, as well as with respect to standard observational constraints, restrictions on the parameter determining the steepness of the potential are obtained. For the almost cosmological constant case $w_{\\phi}\\approx -1$, the solution $\\Omega_{\\phi}(w_{\\phi})$ is obtained. The late-times behaviour $\\Omega_{\\phi}(w_{\\phi})$ of the system is also found. ",
        "conclusions": "In conclusion, we have studied how long a scalar field of exponential potential can play the role of quintessence in a realistic universe filled with pressureless baryonic matter. By deriving and studying directly the first order differential equation of the deceleration parameter $q$ with respect to the state parameter of the scalar field $w_{\\phi}$, we found that there exist whole classes of solutions which enter into accelerating eras not at late times, independently of the parameter $\\mu$ determining the steepness of the potential. Moreover, the condition $w_{\\phi}<-1/3$ does not necessarily imply acceleration. Observational constraints of $\\Omega_{m0}, w_{\\phi 0}$ provide immediate bounds on $\\mu$ for the viable solutions, irrespectively of initial conditions or other parameters, in accordance -- as it has been claimed \\cite{kolda} -- with all the current observations. Using the same first order equation, we have obtained the solution $\\Omega_{\\phi}(w_{\\phi})$ for the almost cosmological constant case $w_{\\phi}\\approx -1$, and the generic late-times solution of the system in the plane of the variables $(\\dot{\\phi}/H,\\sqrt{V}/H)$, $(w_{\\phi},\\Omega_{\\phi})$, $(w_{\\phi},q)$, and $(w_{\\phi},(1+z)^{3}/H^{2})$.             \\[ \\] {\\bf Acknowlegements} We wish to thank J. Garriga, E. Kiritsis, and R. Maartens for useful discussions. This work is partially supported by the European RTN networks HPRN-CT-2000-00122 and HPRN-CT-2000-00131.  G.K. acknowledges partial support from FONDECYT grant 3020031 and from Empresas CMPC. Centro de Estudios Cient\\'{\\i}ficos is a Millennium Science Institute and is funded in part by grants from Fundacion Andes and the Tinker Foundation."
    },
    "0402/astro-ph0402534_arXiv.txt": {
        "abstract": "Detailed modeling of the millisecond brightness oscillations during thermonuclear bursts from low mass X-ray binaries can provide important information about neutron star structure. Until now the implementation of this idea has not been entirely successful, largely because of the negligible harmonic content in burst oscillation lightcurves. However, the recent discovery of non-sinusoidal burst oscillation lightcurves from the accreting millisecond pulsar XTE J1814-338 has changed this situation. We, therefore, for the first time, make use of this opportunity to constrain neutron star parameters. In our detailed study of the lightcurves of 22 bursts, we fit the burst oscillation lightcurves with fully general relativistic models that include light-bending and frame-dragging for lightcurve calculation, and compute numerically the structure of neutron stars using realistic equations of state. We find that for our model and parameter grid values, at the 90\\% confidence level, $Rc^2/GM > 4.2$ for the neutron star in XTE J1814-338. We also find that the photons from the thermonuclear flash come out through the layers of accreted matter under conditions consistent with Thomson scattering, and show that the secondary companion is a hydrogen burning main sequence star, with possible bloating (probably due to X-ray heating). ",
        "introduction": "\\label{sec: 1}  Nearly coherent brightness oscillations have been discovered with the {\\it Rossi} X-ray Timing Explorer ({\\it RXTE}) during thermonuclear X-ray bursts from more than a dozen neutron stars in low-mass X-ray binaries (Strohmayer \\& Bildsten 2003). The large modulation amplitudes at the onset of bursts, the time evolution of the pulsed amplitude during the rise of bursts (Strohmayer, Zhang, \\& Swank 1997), the coherence of the oscillations (Smith, Morgan, \\& Bradt 1997; Strohmayer \\& Markwardt 1999; Muno et al. 2000), and the long term stability of the oscillation frequencies (Strohmayer et al. 1998) led quickly to an interpretation in terms of a ``hot spot\" on the surface of the star, so that the observed flux is modulated by rotation (Strohmayer et al. 1996). This interpretation has been strongly supported by observations of the transient accretion-powered pulsars SAX~J1808--3658 (Chakrabarty et al. 2003) and XTE~J1814--338 (Markwardt, Strohmayer, \\& Swank 2003), in which the frequency of burst oscillations is indeed extremely close to the spin frequency as inferred from persistent oscillations.  Early on it was suggested that the intensity profiles of the pulse lightcurves contain valuable information about the stellar mass and radius (and hence about the equation of state of the dense matter in the stellar core), as well as the surface emission properties and system geometry (Strohmayer, Zhang, \\& Swank 1997; Miller \\& Lamb 1998; Weinberg, Miller, \\& Lamb 2001). The reason is that the lightcurves are affected by general relativistic light deflection as well as special relativistic beaming and aberration, which can be significant at the $\\sim 0.1-0.2$c linear rotational speeds implied for spin frequencies of a few hundred Hertz. However, much of the required information is encoded in the ratios of pulse amplitudes at different harmonics of the spin frequency (although some constraints can be obtained if the amplitude at the fundamental is particularly high; see Nath, Strohmayer, \\& Swank 2002). Until recently, no source had harmonics detected definitively (for upper limits see Muno, \\\"{O}zel, \\& Chakrabarty 2002).  This situation has now changed. Analysis of XTE~J1814--338, the fifth known accreting millisecond pulsar, shows a clear overtone in the pulse profiles of many individual bursts, with an amplitude that can be more than 0.25 times that of the fundamental (Strohmayer et al. 2003). In addition, the 314~Hz frequency of the fundamental oscillation is consistent with the spin frequency seen from the persistent pulsations.  The existence of harmonics and the confirmation of the basic hot spot picture make this source very promising for detailed analysis and constraints on stellar, geometrical, and emission properties.  We note that even more precise information is in principle available from the extremely accurately characterized lightcurves of the persistent pulsations of millisecond pulsars, which have several harmonics (e.g., see Cui, Morgan, \\& Titarchuk 1998; Poutanen \\& Gierli{\\'n}ski 2003). However, the spectrum and profiles of pulses during accretion are affected by complicated details of shock physics and Comptonization. In contrast, thermonuclear burst spectra are well fit by blackbodies, without any trace of a high-energy tail related to Comptonization. They are also consistent with coming from very near the stellar surface. For these reasons, the pulses of XTE~J1814--338 are most promising for an initial investigation.  Combined with analyses of surface spectral line profiles (e.g., from EXO 0748-676; Cottam, Paerels, \\& Mendez 2002), thermonuclear X-ray bursts are key sources for constraining the state of matter inside neutron stars.  Here we report our analysis of the {\\it RXTE} PCA data on 22 thermonuclear bursts from XTE~J1814--338.  In \\S~2 we briefly describe our data extraction techniques.  In \\S~3, we discuss our theoretical models, the fitting procedure and the method of calculation of confidence intervals. We discuss our results in \\S~4. In particular, we show that for the two high-density equations of state (EOS) that we analyze, there is a preference for comparatively less compact neutron stars. We also show that the emission pattern from points on the surface is constrained moderately strongly and is close to that expected from the law of darkening for passive radiative transport in a Thomson scattering atmosphere (Chandrasekhar 1960). In this paper we use geometric units, i.e., $G=c\\equiv 1$. ",
        "conclusions": "\\label{sec:4}  As discussed in \\S~3, we perform comparisons using two EOS models. Figure 6 shows the goodness of fit for a certain burst group and channel range, and for a representative combination of parameter values. The $\\chi^2$ value for this example comes out to be $\\sim 25.5$, for the number of degrees of freedom equal to 24.  This goodness of fit suggests that the overall model framework is a reasonable representation of the data.  In Figure 7, we plot the likelihoods with the parameter $R/M$, for two EOS models, using the data from all 22 bursts. While we can not determine an upper limit of $R/M$, lower limits $(4.7$ for A18$+\\delta v+$UIX and $4.2$ for A18) of $90$\\% confidence regions (given in Table 2) can be computed. We do not consider $R/M$ values greater than $6.0$, as for these values, the stellar mass is too small $(< 1.07 M_\\odot$ for A18 and $< 1.29 M_\\odot$ for A18$+\\delta v+$UIX) to be realistic for an accreting neutron star. A simple extrapolation of the likelihood curves in Figure 7 shows that the probability of $R/M < 3.6$ is very small. However, for $R/M < 3.52$ and for Schwarzschild spacetime, photons from a single point on the stellar surface may reach the observer by multiple paths. As a result, the corresponding lightcurves may be qualitatively different from the ones we calculate. Therefore, we can say that one interpretation of the data is that $R/M$ is greater than 4.2 with $90$\\% confidence, but we can not completely exclude the possibility of $R/M < 3.52$.  Figure 8 displays the likelihood distribution with the observer's inclination angle $i$. This figure shows that $i > 22^{\\rm o}$ with very high probability. Here, as well as for the likelihood calculation of other parameters, we focus our fitting procedure on the range $20^{\\rm o} - 50^{\\rm o}$ for $i$. This is because, with a smaller number of grid points for other parameters, our fitting for inclination angle values in the range $5^{\\rm o} - 90^{\\rm o}$ shows that the likelihood values for $i < 20^{\\rm o}$ are very small, and the inclusion of inclination angles greater than $50^{\\rm o}$ produces only minor changes to our results. However, constraints on inclination angle by observations in other wavelengths would help us constrain other parameters more effectively. Presently, M. Krauss et al. (2004, in preparation) are working on this using optical data.  We keep the beaming parameter $n$ (defined in \\S~3) fixed in all three burst groups in our statistical procedure. This is because the thermonuclear flash is expected to occur at the bottom of a pile of accreted matter. Therefore, although the original emission may be isotropic, what comes out through the layers of accreted matter should be beamed. As the optical depth of these layers is expected to be very high, we expect the law of darkening (and hence the value of $n$) to be the same for all the bursts.  The likelihood distributions for $n$ for two EOS models are shown in Figure 9. For both EOS models, the beaming parameter $n$ peaks at $n=0.75$. This figure shows that $n$ is well constrained on lower side, and the lower limit of the $90$\\% confidence regions (for both the EOS model) is $0.55$ (given in Table 2). The law of darkening for semi-infinite plane-parallel layers with a constant net flux and Thomson scattering is given in Chandrasekhar (1960). In Figure 10, we compare this law of darkening $I(\\psi)$ with our emission function $\\cos^n(\\psi)$. We find that, for $n=0.55$ (where likelihood value is $\\sim 15-20$\\% of the peak likelihood value), our emission function matches with $I(\\psi)$ quite well, except for nearly tangential emission (for which the emitted flux is small). The matching is also reasonably good for the most probable value $n=0.75$. Therefore, the data are consistent with Thomson scattering through an optically thick layer (probably in the accumulated accreted matter) of thermal electrons.  In Table 3, we show the 90\\% confidence intervals for the polar angle of the center of the hot spot $(\\theta_{\\rm c})$ for both the EOS models. Here we give the results for the three burst groups separately, as the values of this parameter may change from one group to another. The union of 90\\% confidence intervals for all burst groups and EOS models is $60^{\\rm o}-139^{\\rm o}$. Assuming that the hot spot is at the magnetic pole, this implies an angle of $40^{\\rm o}-90^{\\rm o}$ between the spin axis and the magnetic axis. The peak of the likelihood distribution for $\\theta_{\\rm c}$ always occurs around the interval $90^{\\rm o}-110^{\\rm o}$ (see Figures 11, 12 \\& 13), hence the magnetic pole may be close to the rotational equator.  Table 4 displays the likelihood values for the angular radius $(\\Delta\\theta)$ of the hot spot for different burst groups and EOS models. For the burst group $1-8$, $\\Delta\\theta$ is comparatively better constrained, and a smaller spot $(\\Delta\\theta \\sim 5^{\\rm o}-25^{\\rm o})$ is likely. For the other two burst groups, comparatively larger spots are probable.  A direct comparison of the likelihoods of two EOS models shows that the difference is not enough to constrain EOS models by this direct method.  The pulsar mass function and orbital period for this source are $0.002016\\,M_\\odot$ and 4.27 hours respectively (Markwardt, Strohmayer, \\& Swank 2003). Hence, a reasonable range of neutron star mass $(1.4\\,M_\\odot - 2.0\\,M_\\odot)$ and our $i > 20^{\\rm o}$ inclination angle constraint imply a companion mass of $0.17\\,M_\\odot - 0.72\\,M_\\odot$. For this range of masses and for the observed orbital period, the companion can be neither a helium main sequence star nor a degenerate star (Bhattacharya \\& van den Heuvel 1991). It is also too large to be a brown dwarf (as was predicted for another accreting millisecond pulsar system SAX J1808.4-3658; Bildsten \\& Chakrabarty 2001). The most probable option is that it is either a hydrogen burning main sequence star, or an evolved (sub-)giant star. However, the companion mass is too small for it to have evolved off the main sequence (B\\\"ohm-Vitense 1992). Therefore, the most likely option is a hydrogen main sequence star. This is also seen in Figure 14. If the companion is a normal hydrogen burning main sequence star (as shown in Figure 14), its maximum possible mass is $\\sim 0.48 M_\\odot$. This is because, the radius of the companion can not be smaller than that of a hydrogen main sequence star (for a given mass), but it can be bigger, because it can be bloated due to X-ray heating by the neutron star and/or the accretion disk. This upper limit of companion mass corresponds to a lower limit of observer's inclination angle $24^{\\rm o}$ (using the pulsar mass function value and a lower limit of neutron star mass $1.4 M_\\odot$). Interestingly, this lower limit of inclination angle is close to that found by our lightcurve fitting. However, if the mass fraction of hydrogen in the hydrogen burning main sequence companion star is very low, then it may have a higher mass than $0.48 M_\\odot$ (see Rappaport \\& Joss 1984). If we consider the lower limit $(0.17\\,M_\\odot)$ of the companion mass, then its radius is about $98$\\% bigger than the normal radius (Bhattacharya \\& van den Heuvel 1991).  In conclusion, the significant overtones observed in the pulse profiles of burst brightness oscillations from XTE~J1814--338 open up new possibilities for constraints on neutron star structure as well as on source geometry and emission properties. In a subsequent paper, we will study the analysis possible with future large-area detectors, in particular whether EOS models may be constrained strongly based on their likelihoods."
    },
    "0402/astro-ph0402228_arXiv.txt": {
        "abstract": "A phenomenological approach is proposed to the problem of universe accelerated expansion and of the dark energy nature. A general class of models is introduced whose energy density depends on the redshift $z$ in such a way that a smooth transition among the three main phases of the universe evolution (radiation era, matter domination, asymptotical de Sitter state) is naturally achieved. We use the estimated age of the universe, the Hubble diagram of Type Ia Supernovae and the angular size\\,-\\,redshift relation for compact and ultracompact radio structures to test whether the model is in agreement with astrophysical observation and to constrain its main parameters. Although phenomenologically motivated, the model may be straightforwardly interpreted as a two fluids scenario in which the quintessence is generated by a suitably chosen scalar field potential. On the other hand, the same model may also be read in the context of unified dark energy models or in the framework of modified Friedmann equation theories. ",
        "introduction": "In the last few years, an increasing bulk of data has been accumulated leading to the emergence of a new cosmological scenario. The Hubble diagram of type Ia Supernovae (SNeIa) first indicated that the universe expansion is today accelerating \\cite{Riess98,Perlm99}. The precise determination of first and second peaks in the anisotropy spectrum of cosmic microwave background radiation (CMBR) by the BOOMERanG and MAXIMA collaborations \\cite{CMBR} strongly suggested that the geometry of the universe is spatially flat. When combined with the data on the matter density parameter $\\Omega_M$, these results lead to the conclusion that the contribution $\\Omega_X$ of dark energy is the dominant one, being $\\Omega_M \\simeq 0.3, \\Omega_X \\simeq 0.7$. This picture has been strenghtened by the recent determination of CMBR spectrum measured by the WMAP team \\cite{WMAP}.  According to the standard recipe, pressureless cold dark matter and a homogenously distributed cosmic fluid with negative pressure, referred to as {\\it dark energy}, fill the universe making up of order $95\\%$ of its energy budget. What is the nature of this dark energy still remains an open and fascinating problem. The simplest explanation claims for the cosmological constant $\\Lambda$ thus leading to the so called $\\Lambda$CDM model \\cite{Lambda}. Although being the best fit to most of the available astrophysical data \\cite{WMAP}, the $\\Lambda$CDM model is also plagued by many problems on different scales. If interpreted as vacuum energy, $\\Lambda$ is up to 120 orders of magnitudes smaller than the predicted value. Furthermore, one should also solve the {\\it coincidece problem}, i.e. the nearly equivalence of the matter and $\\Lambda$ contribution to the total energy density.  As a response to these problems, much interest has been devoted to models with dynamical vacuum energy, dubbed {\\it quintessence} \\cite{QuintFirst}. These models typically involve scalar fields with a particular class of potentials, allowing the vacuum energy to become dominant only recently (see \\cite{PR02,Pad02} for comprehensive reviews). Altough quintessence by a scalar field is the most studied candidate for dark energy, it generally does not avoid {\\it ad hoc} fine tuning to solve the coincidence problem. On the other hand, a quintessential behaviour may also be recovered without the need of scalar fields, but simply by taking into account the effective contribution to cosmology of some (usually neglected aspects) of fundamental physics \\cite{review}. A first tentative were undertaken showing that a universe with a non vanishing torsion field is consistent with SNeIa Hubble diagram and Sunyaev\\,-\\,Zel'dovich data on clusters of galaxies \\cite{torsion}. The same quintessential framework can be obtained with the extension of Einstein gravity to higher order curvature invariants leading to a model which is in good agreement with the SNeIa Hubble diagram and the estimated age of the universe \\cite{curv}. It is worth noting that these alternative schemes provide naturally a cosmological component with negative pressure whose origin is simply related to the geometry of the universe itself thus overcoming the problems linked to the physical significance of scalar fields.  Despite the broad interest in dark matter and dark energy, their physical properties are still poorly understood at a fundamental level and, indeed, it has never been shown that the two are in fact two different ingredients. This observation motivated the great interest recently devoted to a completely different approach to quintessence. Rather than the fine tuning of a scalar field potential, it is also possible to explain the acceleration of the universe by introducing a cosmic fluid with an exotic equation of state causing it to act like dark matter at high density and dark energy at low density. An attractive feature of these models is that they can explain both dark energy and dark matter with a single component (thus automatically solving the coincidence problem) and have therefore been referred to as {\\it unified dark energy} (UDE) or {\\it unified dark matter} (UDM).  Some interesting examples are the generalized Chaplygin gas \\cite{Chaplygin}, the tachyonic field \\cite{tachyons} and the condensate cosmology \\cite{Bruce}.  It is worth noting that all the dark energy models (both with scalar fields or UDE) proposed up to now predict that the expansion of the universe is a two phase process\\,: it is first determined by a matter\\,-\\,like term and it is then driven by the quintessence\\,-\\,like component towards an asymptotically de Sitter state. However, it is well established that there is also a third phase preceding these two, i.e. the radiation dominated era. It is thus interesting to look for a model which is able to predict a smooth transition from one phase to the following one in a natural way.  A quite confused picture emerges from the previous discussion about dark energy and its nature. Many models have been proposed to succesfully reproduce the astrophysical observations available up to date, but they are so different each other that the mistery of the dark energy is far to be solved. Given the state of the art, a different approach to the problem is welcome. We think that a first step toward understanding the nature of dark side (dark energy and dark matter) of the universe is to explore phenomenological models which are able to reproduce what we observe. To this aim, we consider a single fluid whose energy density scales with the redshift in such a way that the radiation dominated era, the matter domination and the accelerating phase have been naturally achieved. We constrain the model parameters using the estimated age of the universe, the SNeIa Hubble diagram and the angular size\\,-\\,redshift relation of radio structures. Although phenomenologically motivated, nonetheless the model we propose may be physically interpreted in terms of an effective two fluids scenario with the dark energy component represented by a scalar field with a suitably chosen interaction potential. On the other hand, the models may also be considered as a new member of the UDE class.  The plan of the paper is as follows. In Sect.\\,II we present the main feature of the class of phenomenological models we propose. Some considerations on the allowed range for the model parameters lead us to restrict our attention (at least in this first analysis) to a particular class of models for the reasons we explain in Sect.\\,III. The estimated age of the universe, the SNeIa Hubble diagram and the angular size\\,-\\,redshift relation for compact radio structures are used in Sects.\\,IV, V and VI to test whether the model is a viable one and to constrain its parameters. Sect.\\,VII is devoted to the interpretation of the model in terms of an effective two fluids model with a scalar field whose interaction potential is reconstructed. Finally, we summarize and conclude in Sect.\\,VIII. ",
        "conclusions": "The concordance cosmological model assumes that pressureless cold dark matter and the dominant quintessence field drive the evolution of the universe leading to the observed accelerated expansion. As an alternative to this picture, unified dark energy models have been proposed where a single cosmic fluid acts both as dark matter and dark energy depending on the value of the energy density.  Here we have presented a general class of models and tested it against the astrophysical observations available up to date. The starting point is the assumption that the fluid energy density $\\rho$ depends on the scale factor $R$ as shown in Eq.(\\ref{eq: rhor}). The model is characterized by five parameters, but we have limited our attention to those models where the energy density smoothly interpolates among the three main phases of the universe evolution, i.e. a radiation dominated era followed by matter domination and an asymptotic de Sitter state. This ansatz is motivated by the phenomenology we observe since every consistent picture of the universe evolution predicts the onset of the three different phases we have quoted above. The model has been tested against some astrophysical observations, namely the age of the universe, the SNeIa Hubble diagram and the angular size\\,-\\,redshift relation for compact radio structures. The succesful results of these tests is a strong evidence of the reliability of this phenomenological approach. It is worth noting that our results may be interpreted both in the frame of UDE models and in the standard picture of a two fluid universe made of matter and dark energy. In this latter case, if a scalar field $\\phi$ is assumed to be the origin of this component, the interaction potential $V(\\phi)$ may be directly reconstructed from the observations without any a priori hypothesis on its form.  Although the model has been shown to successfully fit the available observations, further analysis is still needed.  First, there are other tests that can be used to better constrain the model parameters. In particular, we have not discussed the growth of perturbations which determines the CMBR anisotropy spectrum. Since the energy density $\\rho$ scales as the matter term in $\\Lambda$CDM model during the structure formation era, perturbations should grow in the same manner so that the fit to the observed CMBR spectrum is likely to be successful. However, this naive expectation has to be carefully checked since it depends on how we interpret the model. Actually, this qualitative picture should indeed be valid if Eq.(\\ref{eq: rhor}) is meant as a phenomenological description of a two fluids scenario since, in this case, the quintessence\\,-\\,like fluid almost vanishes during the epoch of structure formation and therefore everyting works as in the usual scenario. On the other hand, if the UDE interpretation is preferred, one has to explicitely solved the perturbation equations by using explicitely Eq.(\\ref{eq: rhor}) for the matter enery density. We are, however, confident that the main results are not changed for the model with $(\\alpha, \\beta) = (3, 4)$ since, in this case, Eq.(\\ref{eq: rhor}) reduces to the usual expression $(\\rho \\propto R^{-3})$ in the structure formation epoch. Finally, a different approach has to be considered if the interpretation of the model in the framework of modified Friedmann equations. In this case, one should first obtain the corresponding modified Newtonian potential and then work out the consequences on the equations describing the growth of perturbation to finally obtain a coherent description of the structure formation process and of its imprint on the CMBR anisotropy spectrum.  Secondly, in this analysis, we have held fixed the two slopes parameters $(\\alpha, \\beta)$ to the values they must have to perfectly mimic the scaling of the energy density with the scale factor $R$ during the radiation and matter dominated era. However, it is still possible that other values of $(\\alpha, \\beta)$ fit the astrophysical data we have considered. It is thus interesting to repeat the same analysis performed here also varying these parameters. To this aim, however, more data are welcome since adding more parameters may introduce strong degeneracies among some of them. In order to solve these problems, one should also consider the observations of large scale distribution of matter as measured by the two major ongoing galaxy surveys (the 2dFGRS \\cite{2dFGRS} and the SDSS \\cite{SDSS}) or the data on the Lyman $\\alpha$ forest \\cite{Lyalpha}.  As a final remark, we want to stress again the underlying philosophy of this work. Contrary to the usual approach of proposing a theory and then test it against observations, we have preferred to start from the phenomenology we observe to investigate what are the main features a theory should have to give a consistent and realistic picture of the universe."
    },
    "0402/astro-ph0402158_arXiv.txt": {
        "abstract": "The idea that the radio jets in AGN contain material in relativistic motion is supported by many lines of observational evidence, including morphology, brightness temperature estimated with interferometers and with intrinsic variations, interstellar scintillations, X-rays, and superluminal motion. These are largely independent, and taken together make an irrefutable case for relativistic motion. ",
        "introduction": "Scientists were surprised in the early 1950s, when it became clear that the ``radio stars\" in fact were extragalactic and hence very energetic, and the general idea of a calm orderly Universe had to be questioned. In later years the discovery of black holes, quasars, pulsars, and other exotic objects changed the view completely, and brought us to the {\\it violent} Universe as we understand it today.  In a prescient article, Rees (1966) suggested bulk relativistic motion as an explanation for the rapid variations in flux density seen in quasars; and he also predicted the phenomenon of superluminal apparent transverse velocity, which was discovered five years later. At that point most people believed that cosmic jets contained relativistic jets of plasma, but what process could collimate and accelerate whole solar masses of material to relativistic speeds? The current paradigm is that a giant black hole lies at the center of a galaxy, and the process of accretion combined with a magnetic field generates the bulk relativistic flow. The details of this scenario are not well understood, but gravitational accretion is the only known means for generating the required energy per unit volume.  In this paper we review the radio observations that provide evidence for relativistic motion.  The plan of the paper is to successively discuss frequency shift, morphology and Doppler boosting, brightness temperature measured with interferometers and estimated with intrinsic variability, interstellar scintillations, inverse-Compton X-rays, and superluminal motion.  Individually, and especially taken together, these make a convincing case for relativistic motion in the radio jets. This conclusion is independent of any specific cosmology.  Relativistic motion implies that the observed radiation patterns are narrow, that Doppler boosting increases the observed flux density, and that the observed time scale is shrunk below that in the frame of the AGN. These affect the observable quantities; namely, frequency, luminosity, apparent transverse speed, and variability; and there should be statistical relations among them.  We briefly touch on this in \\S 4.  A number of earlier papers have discussed the different types of observational data bearing on relativistic motion. We mention only a few of them here: \\cite{Gh93} compared X-ray and VLBI data to derive Doppler factors and compared them with other beaming indicators; Vermeulen and Cohen (1994) assembled all the existing data on superluminal sources and analyzed several models involving the distributions of Lorentz factor and of the ratio of pattern to beam velocity; Lister and Marscher (1997) made Monte Carlo calculations of distributions of the measureable parameters, and of their correlations; and \\cite{LV99} estimated Lorentz factors and also the intrinsic temperature in the synchrotron sources.  These past surveys mainly used inhomogeneous data culled from the literature. However, two large surveys on the VLBA, at 5 MHz \\cite{VB03} and 15 GHz (Kellermann \\etal~2004, hereafter K04) are now bearing fruit, and analyses with these superior data sets should soon give much better results. ",
        "conclusions": "We have described five independent sets of radio phenomena which show that there is relativistic motion in the compact jets of AGN. These are the observations of morphology, interferometry, intrinsic variability, interstellar scintillations, and superluminal motion. X-ray measurements provide further confirmation, although these results are dependent on the VLBI measurements from which the brightness temperatures are determined. Each of these phenomena has a theoretical difficulty which is readily alleviated by assuming that there is bulk relativistic motion towards the observer.  The morphology, interferometry, interstellar scintillations, and the X-ray method are independent of redshift and hence of cosmology. Interpretations based on intrinsic variability and on internal proper motion, however, require a knowledge of the distance to the source. The variability methods, and the interferometry method, require an assumption of an intrinsic temperature in the synchrotron cloud.  The distribution of $\\delta$ peaks at low values and decreases to $\\delta \\sim 30$ The Lorentz factor $\\gamma$ is a more fundamental quantity than $\\delta$, but is harder to estimate.  Values estimated for $\\gamma$ range up to 34.  A comparison of $\\rm \\beta_{app}$ and $\\rm \\delta_{var}$ shows that the superluminal speeds and the variability Doppler factors are roughly consistent with the standard relativistic motion paradigm, provided the intrinsic temperature in the synchrotron clouds is $\\sim 2\\times 10^{10}$ K. This value is below the various theoretical upper limits, and is close to Readhead's estimate (1994), based on equipartition between the energy in the particles and that in the magnetic field."
    },
    "0402/astro-ph0402472_arXiv.txt": {
        "abstract": "{\\small{ From a survey of edge-on disks, we find that disk galaxies show a sharp, mass-dependent transition in the structure of their dusty ISM.  In more massive, rapidly rotating disks with $V_c>120\\kms$, we see the well-defined dust lanes traditionally associated with edge-on galaxies.  However, in more slowly rotating, lower mass galaxies with $V_c<120\\kms$, we find {\\emph{no}} dust lanes.  Instead, the distribution of dust in these galaxies has a much larger scale height and thus appears more diffuse.  Evidence suggests that the change in scale height is due primarily to changes in the turbulent velocities supporting the gas layer, rather than to sharp changes in the gas surface density.  A detailed analysis of our sample reveals that the decrease in the dust scale height is associated with the onset of disk instabilities, evaluated for a mixed star$+$gas disk.  Specifically we find that all of the high-mass galaxies with dust lanes are gravitationally unstable, and thus are prone to fragmentation and gravitational collapse along spiral arms. Empirically, our data implies that turbulence has lower characteristic velocities in the presence of disk instabilities, leading to smaller gas scale heights and the appearance of narrow dust lanes.  The drop in velocity dispersion may be due either to a switch in the driving mechanism for turbulence from supernovae to gravitational instabilities or to a change in the response of the ISM to supernovae after the ISM has collapsed to a dense layer.  We hypothesize that the drop in gas scale height may lead to significant increases in the star formation rate when disk instabilities are present.  First, the collapse of the gas layer increases the typical gas density, reducing the star formation timescale.  Second, the star formation efficiency increases due to lower turbulent velocities.  These two effects can combine to produce a sharp increase in the star formation rate with little change in the gas surface density, and may therefore provide an explanation for the Kennicutt surface density threshold for star formation.  Our data also suggest that star formation will be systematically less efficient in low mass disks with $V_c<120\\kms$, since these galaxies are stable and lie entirely below the Kennicutt surface density threshold.  In these stable systems the effective nucleosynthetic yield is reduced because the star formation timescale becomes longer than the gas accretion timescale, supressing the metallicity.  This effect can possibly produce the observed fall-off in metallicity at rotation speeds less than $V_c<120\\kms$.  Thus, infall provides an equally plausible explanation of the mass-metallicity relation in disks as global outflows driven by supernova winds.  The transitions in disk stability, dust structure, and/or star formation efficiency may also be responsible for the observed changes in the slope of the Tully-Fisher relation, in the sharp increase in the thickness of low mass galaxy disks, and in the onset of bulges in galaxies with $V_c\\gtrsim120\\kms$.  The latter observation lends support to theories in which bulges in late-type galaxies grow through secular evolution in response to disk instabilities.  We include in this paper relationships between the surface density and the vertical stellar velocity dispersion as a function of galaxy rotation speed, which may be useful constraints on galaxy formation models.  }} ",
        "introduction": "Dust lanes have long been recognized as a characteristic signature of edge-on disk galaxies.  Optical images show that edge-on disks are typically bisected by a dark narrow layer of dust which has a typical scale height half as large as that of the stellar disk (Xilouris et al.\\ 1999).  However, not all galaxies show this conspicuous feature. In many late-type, low mass galaxies the dust lane is absent.  It has been assumed that the lack of a dust lane reflects an overall lack of dust, possibly due to low metallicity (van den Bergh \\& Pierce 1990). However, low mass galaxies {\\emph{do}} have dust (see compilation in Lisenfeld \\& Ferrara 1998).  The difference lies in the morphology of the dust distribution, which is less concentrated within the disk and significantly more porous (see e.g.\\ Matthews \\& Woods 2001).  We are engaged in an extensive multicolor imaging program of edge-on bulgeless disks (Dalcanton \\& Bernstein 2000) that is ideal for probing variations in dust lane morphology.  Our study of this sample indicates that the transition from a thicker, more porous distribution of dust to a narrow lane is remarkably sharp.  The dramatic change in the morphology of the dusty ISM with increasing galaxy mass cannot be due solely to increasing metallicity or dust content, but rather must result from the onset of a distinct physical mechanism.  As discussed in more detail in \\S\\ref{theorysec}, dust is an excellent tracer of gas in galaxies, with the regions of highest extinction tracing regions of the highest gas density, namely the cool neutral and the cold molecular medium.  This correlation is a clue to the observed transition in the dust distribution, and suggests that the dust morphologies indicate a change in the structure of the cool and cold ISM.  This connection allows us to use the transition in dust morphology to diagnose the processes which shape the global distribution of the cold ISM.  The outline of the paper is as follows.  In \\S\\ref{datasec} we present the evidence in our sample for a sharp change in the dust lane morphology with decreasing mass.  We suggest that the observed disappearance of dust lanes results from a sharp increase in the thickness of the cold ISM in low mass galaxies.  In \\S\\ref{quantsec} we discuss the physics which might drive this morphological dichotomy, and find that disk stability is the only intrinsic property that changes sharply with dust morphology.  In \\S\\ref{theorysec}, we discuss the relationship between gravitational instabilities and interstellar turbulence, and describe how it might lead to a thinner layer of cold ISM in unstable high mass galaxies.  In \\S\\ref{SFsec} we discuss the implications of our observations for the characteristic scales and velocities of turbulence driven by gravitational instabilities, and suggest that these may provide an explanation for the Kennicutt star formation threshold.  In \\S\\ref{metallicitysec}, we discuss how the resulting drop in star formation efficiency in low mass disk galaxies leads to lower metallicities and nucleosynthetic yields at the same characteristic velocity where dust lanes disappear, without resorting to global galactic outflows. In \\S\\ref{observationsec} we discuss how connections between dust morphology, disk stability, and star formation efficiency may be evident in other observations, such as the Tully-Fisher relation, the bulge-disk ratio as a function of galaxy mass, the thickness of dwarf galaxies, and the properties of galaxies at high redshift.  We summarize our conclusions in \\S\\ref{conclusionsec}. ",
        "conclusions": "\\label{conclusionsec}  Our analysis of the distribution of dust in a large sample of edge-on late-type disks shows that the dust lane morphology appears only within galaxies rotating faster than $V_c\\approx120\\kms$. Based on ground-based and archival HST images, we suggest that dust lanes appear in high mass galaxies because the thickness of the dust layer drops substantially, increasing its line-of-sight opacity and its contrast against the stellar disk.  Dust is closely coupled to the cold ISM, and thus these data indicate that the scale height of the cold ISM is likely to have dropped as well.  The scale height is set by the balance between the surface density of the disk and the velocity dispersion of the gas.  However, the surface density does not increase sharply above the $V_c=120\\kms$ transition, and thus we suggest that the smaller scale height is due to lower velocity dispersions in high mass galaxies.  We have investigated possible origins for the observed transition at $V_c=120\\kms$.  We find that this transition point corresponds to the rotation speed above which disks become gravitationally unstable.  We speculate that the disk instabilities lead to a thin dust lane through collapse and fragmentation of high density gas within spiral arms.  We speculate on two possible effects which might reduce the {\\emph{rms}} velocity dispersion within fragmenting spiral arms and thus lead the new equilibrium scale height for the gas to be lower.  The first possibility is that the fragmentation process changes the dominant driving mechanism for turbulence from supernovae to fragmentation. The second possibility is that supernovae continue to be the dominant mechanism driving the turbulent velocities, but that the resulting velocities are lower in the dense collapsed gas layer.  Either of these two scenarios would affect the velocity dispersion of only the gas within the spiral arms.  We suggest that the change in dust lane morphology may provide a physical explanation for the Kennicutt (1989) star formation threshold.  Our data indicate that when disk instabilities are present, the cold ISM collapses to a denser layer with smaller turbulent velocities.  The increased density should increase the star formation rate by roughly a factor of three, assuming a Schmidt law, with no change in gas surface density.  In addition, numerical simulations suggest that star formation is more efficient when turbulent velocities are smaller, further increasing the star formation rate.  The combination of these two effects can lead to sharp increases in the star formation rate in the gravitationally unstable regions of disks, reproducing Kennicutt star formation threshold.  Our data suggest that galaxies with $V_c<120\\kms$ will tend to lie entirely below the Kennicutt threshold.  They will therefore have systematically low star formation efficiencies and long star formation timescales.  The long star formation timescales can potentially explain why the metallicity and effective yield of galaxies decrease systematically below $V_c\\sim120\\kms$, but are constant above.  For the low mass stable disks, the star formation timescale is likely to be sufficiently long that accretion can dilute the gas in between significant star formation events.  In this regime, the metallicity and the effective yield are systematically depressed, so that the galaxies with the longest star formation timescales show the largest deviations from theoretical closed box yields.  For unstable disks, however, star formation is extremely efficient and always shorter than the infall timescale.  In this regime, chemical evolution models suggest that the metallicity and effective yield approach a constant value.  This interpretation may explain the observed trends between metallicity and galaxy mass without invoking mass loss through outflows.  We give evidence that the onset of disk instabilities at $V_c\\sim120\\kms$ may be associated with the onset of bulges at a comparable velocity.  This gives support to theories of secular evolution for bulge formation in late-type galaxies.  We also demonstrate that there is a significant jump in the thickness of stellar disks for galaxies with $V_c<120\\kms$.  We argue that this results naturally from the increase in the thickness of the cold ISM in these stable galaxies.  In a subsequent paper we will explore the implication of our results for the molecular gas content of galaxies."
    },
    "0402/astro-ph0402644_arXiv.txt": {
        "abstract": "{ The X--ray afterglow of the Gamma-Ray Burst \\gr,  associated to SN2003dh at $z$=0.1685, has been observed with \\xmm 258 days after the burst explosion. A source with flux of (6.2$\\pm$2.3)$\\times10^{-16}$ erg cm$^{-2}$ s$^{-1}$ (0.5-2 keV) has been detected at the GRB position. This measurement, together with a re--analysis of the previous X--ray observations, indicates a flattening of the  X--ray light curve  $\\sim$40 days after the burst. This  is in remarkable agreement with the scenario invoking the presence of two jets with different opening angles. The wider jet should be responsible for the observed flattening due to its transition into the non--relativistic Sedov--Taylor phase. ",
        "introduction": "Multifrequency studies of  Gamma--Ray Burst (GRB) afterglows at late times ($\\gsim$10 days)  can be used to test some predictions of the standard, as well as alternative, models. This has been done so far for a handful of bursts (e.g., Frail, Waxman \\& Kulkarni 2000; Frail et al. 2004). Thanks to its high flux and low redshift ($z$=0.1685, Greiner et al. 2003; Caldwell et al. 2003), GRB 030329 is  an ideal target for such kind of investigations.  In fact \\gr had a very large fluence of $\\sim$10$^{-4}$ erg cm$^{-2}$ (30--400 keV, Ricker et al. 2003), in the top 1\\% of all  detected GRBs, and its optical transient had magnitude 13 one hour after the explosion (Peterson \\& Price 2003; Torii 2003). \\gr is also the first GRB unambiguously associated with a supernova (SN2003dh, Stanek et al. 2003; Hjorth et al. 2003). Systematic multiwavelength monitoring of its afterglow (e.g., Lipkin et al. 2004; Matheson et al. 2003; Tiengo et al. 2003; Sheth et al. 2003) is yielding important information for the understanding of the jet structure, GRB energetics and circumburst environment.    Here we report on a new \\xmm observation  performed 258 days after the burst in order to study the late evolution of the X--ray afterglow. We also re--analysed the previous observations to derive in a consistent way all the X--ray fluxes  in the 0.5-2 keV energy range. This is the range where the effective area of the \\xmm EPIC instrument is highest and the corresponding fluxes are less affected by uncertainties in the afterglow spectral shape. ",
        "conclusions": "The combination of the exceptional brightness of \\gr and the high sensitivity of \\xmm has allowed us to study a GRB X--ray afterglow up to very late time, leading to a $\\lsim$3$\\sigma$ detection of the afterglow 258 days after the burst.   Although other interpretations are possible, the \\xmm data at $t=$37, 61 and 258 days are consistent with the X--ray light curve being dominated by the wide jet component envisaged by Berger et al. (2003), Lipkin et al. (2004) and Sheth et al. (2003). In particular, and more importantly, we have evidence that this component became non--relativistic at $\\sim$40 days after the GRB explosion, in remarkable agreement with what predicted assuming the parameters (kinetic energy and density of the external medium) derived by the above authors by fitting the radio and optical light curves.  The immediate prediction of this interpretation is that also the radio light curve should show a flattening from $\\delta=2$ to $\\delta_{NR}\\sim 1.5$ after $t\\sim$40 days. The publicly available radio data at these late times have large uncertainties and cannot exclude such a behavior. Only further analysis of more radio data could settle this issue. The lack of a  flattening in the radio light curve would imply that the fireball has not yet reached the non--relativistic phase, due to a larger total kinetic energy of the wide jet and/or a smaller density for the interstellar medium. This would then require a  different reason for the observed flattening in the X--ray light curve, e.g. a larger supernova contribution in the X--ray band. This would imply an X--ray luminosity at least one order of magnitude higher than that observed in other type  Ic and  II supernovae (see Fig. 3 of \\cite{kouveliotou})."
    },
    "0402/astro-ph0402122_arXiv.txt": {
        "abstract": "We discuss the Skyrme model for strong interactions and the concept of a Skyrmion fluid. The corresponding compact objects, namely Skyrmion stars, constructed from the equation of state describing such a fluid are presented and compared to models of neutron stars based on modern equation of states. We suggest plausible Skyrmion star candidates in the {\\bf 4U 1636$-$53} and {\\bf 4U 1820$-$30} low mass X-ray binary systems where the suggested masses of the accreting compact companion ($\\sim 2.0M_{\\odot}$) remain a challenge for neutron star models. ",
        "introduction": "\\label{one}  We are closing in on neutron stars both observationally and theoretically. Observationally more and better masses  and radii are determined by a number of different methods. The ``large\" masses implied in few cases (as in the Vela X-1 pulsar where a mass of $\\sim 1.78 M_{\\odot}$ has been suggested; \\cite{barziv01}) poses a real challenge to models of neutron stars built using the so-called modern equation of state (EOS) where the uncertainties are reduced by improved two- and three-body forces, relativistic effects and many-body calculations \\cite{heiselbergb00}. Even the stiffest EOS so far developed seem to be facing difficulty in accounting for the extreme values (up to $2.4M_{\\odot}$; Sect. 4). In \\cite{ouyed99}, hereafter OB, as an alternative, we constructed an EOS of dense matter based on the Skyrme model for strong interactions which represents baryons as solitons of classic pionic fields. The resulting compact objects we named {\\it Skyrmion stars} (SSs) \\cite{ouyed02} are intrinsically heavier (due to the stiffness of the Skyrmion fluid; hereafter SF) than any other type of compact stars computed using modern EOS. SSs can be as massive as $\\sim 2.8M_{\\odot}$ leading us to speculate (given the above mentioned observations) that these might exist in nature\\footnote{Skyrmion stars, like neutron stars, are likely to be born with masses around $1.5 M_{\\odot}$. We expect only older SSs that have accreted enough mass to reach these extreme masses.}. The paper is presented as follows: Sect. 2 is devoted to the discussion of the Skyrme model and its link to Quantum-ChromoDynamics (QCD). The problem of the {\\it missing attraction} in the Skyrme model is described. In Sect. 3, we discuss the role the dilaton (the glueball potential in QCD) could play in curing such a problem by binding Skyrmions together to form the SF. We end Sect. 3 by  reminding the reader of the basic properties of the SF and the resulting EOS used to construct models of SSs. The astrophysical implications follow in Sect. 4 where we compare SSs to stars constructed using modern EOS of dense matter. We conclude in Sect. 5. ",
        "conclusions": "We gave a brief historical overview of the Skyrme model, its predictions of hadronic interactions and its interesting connection to QCD. Here we showed how the repulsive term in the Skyrme model can be removed by coupling the  Skyrmion to the dilaton field. This lead to the concept of the Skyrmion fluid and the related hypothetical stars we called Skyrmion stars. The stiffness of the Skyrmion fluid allows for SSs to be as massive as $2.8M_{\\odot}$. The SSs show unique features; for a given mass their radii are in general larger than those of neutron stars constructed using modern EOSs. We discussed examples in astrophysics where SSs might constitute plausible candidates. Future observations constraining the mass-radius plane of compact stars would most likely prove or rule out the existence of SSs in nature.  \\vskip\\baselineskip {\\bf Acknowledgements.} I am grateful to S. Morsink, G. K\\\"albermann and R. Bhaduri for encouraging help and valuable discussions. The research of R.O. is supported by grants from the Natural Science and Engineering Council of Canada (NSERC)."
    },
    "0402/astro-ph0402408_arXiv.txt": {
        "abstract": "The distance to the black hole binary \\gx\\ remains a topic of debate. We examine high-resolution optical spectra of the \\nad\\ lines resolving the velocity structure along the line of sight.  We find this to be complex, with at least nine components, mostly blue-shifted, spanning a velocity range of nearly 200\\,km\\,s$^{-1}$. The presence of components with a large blue-shift rules out a nearby location and requires that the binary be located at or beyond the tangent point, implying a lower limit to the distance of $\\sim6$\\,kpc. The presence of a significant red-shifted component at $+30$\\,km\\,s$^{-1}$ is even more intriguing as \\gx\\ also has a slightly positive systemic velocity, suggesting that the source, and this cloud, could be on the far side of the Galaxy, where the radial velocities due to Galactic rotation become positive again.  If this is the case, we require a distance of $\\sim15$\\,kpc.  This is less secure than the 6\\,kpc lower limit however.  We discuss the implications of these possible distances for the outburst and quiescent luminosities, and the nature of the companion star, and argue that a large distance is consistent with these characteristics.  In particular, it would explain the non-detection of the companion star during the faintest states. ",
        "introduction": "The distance of X-ray binaries is challenging to determine precisely. They are usually too distant for current parallax measurements, and there are no really reliable standard candles, other than type-I X-ray bursts showing radius expansion, and these only occur in a subset of neutron star systems (see \\citealt{Kuulkers:2003a} and references therein).  For black holes, such as \\gx, the problem is much less tractable.  Fortunately the majority of such systems undergo periods of deep quiescence when it is possible to detect the companion star, and determine the binary parameters.  Hence both the radius and temperature (i.e.\\ surface brightness) can be estimated and the distance derived from the apparent brightness (see e.g.\\ \\citealt{Gelino:2001a}).  For \\gx\\ this is not possible, as even in a very low luminosity state, no spectral features of the companion star were seen in VLT/FORS observations \\citep{Shahbaz:2001a}.  This is frustrating for distance and parameter estimation, and very puzzling in its own right.  One explanation for the non-detection of the companion is that the distance has been significantly underestimated.  Current distance estimates are very uncertain; \\citet{Buxton:2003a} has provided a convenient summary.  Early estimates were rather low, mostly falling in the 1--4\\,kpc range.  More recently, larger estimates have been proposed.  The non-detection of the companion star \\citep{Shahbaz:2001a} was used as evidence for a distance $\\ga 5.6$\\,kpc.  A more novel approach was taken by \\citet{Maccarone:2003a} who has shown that in most black hole binaries, transitions between the high/soft and low/hard states occur at approximately a fixed fraction of the Eddington limit.  It was suggested that the transition luminosity can then be used as a mass-dependent standard candle. Assuming a minimum black hole mass for \\gx\\ of $\\ga5.8$\\,M$_{\\odot}$ \\citep{Hynes:2003a}, \\citet{Maccarone:2003a} estimate a distance $\\ga 6$\\,kpc, allowing for scatter about the relation.  The actual distance of \\gx\\ thus remains very uncertain, but could be rather large.  This introduces considerable uncertainty into interpreting its quiescent and outburst luminosities, and the nature of the companion.  The distance also has important implications for interpretation of the recently discovered resolved jets; \\citet{Gallo:2004a} estimated that the motion of the jet head requires an apparent velocity of 0.9\\,c even at 4\\,kpc.  A significantly larger distance would therefore imply apparent superluminal motion.  We have therefore analyzed interstellar features in high resolution spectra described in Section~\\ref{ObservationSection} to elucidate the properties of the line-of-sight in more detail.  In Section~\\ref{ModelSection} we perform a detailed study of the properties of line-of-sight \\NaI\\ and \\CaII\\ absorption, resolving most individual clouds in \\NaI\\ and hence determining cloud-by-cloud velocities, velocity dispersions, and \\NaI\\ column densities.  We also have some information about the relative abundances of \\CaII\\ to \\NaI\\ as a function of velocity.  We discuss the implications for the distance of \\gx\\ in Section~\\ref{RotationSection}, and for the reddening in Section~\\ref{ReddeningSection}.  Finally in Section~\\ref{DiscussionSection} we consider how the range of distances proposed affect estimates of the outburst and quiescent X-ray luminosities and the nature of the companion star, and in Section~\\ref{ConclusionSection} we summarize our conclusions. \\newpage ",
        "conclusions": "\\label{ConclusionSection} We have analyzed \\NaI\\ and \\CaII\\ absorption along the line of sight to \\gx.  We find a rich absorption spectrum with a large velocity dispersion.  Velocities are predominantly blue-shifted and span the range of velocities expected for Galactic rotation along this line-of-sight. The presence of multiple features with a significant column density near the tangent point velocity suggests that \\gx\\ is either at or beyond the tangent point, i.e.\\ $d\\ga 6$\\,kpc.  The presence of absorption at red-shifted velocities, together with the systemic velocity of \\gx\\ itself, suggest the possibility that the object could be at an even greater distance, on the far-side of the Galaxy, $d\\ga15$\\,kpc.  Neither of these pieces of evidence are as secure as that for a location beyond the tangent point, however.  Our analysis also indicates a total \\NaI\\ column density along the line of sight of at least $3.5\\times10^{13}$\\,cm$^{-2}$, and this suggests a reddening of at least $\\ebv=0.5$, depending on assumed ratios of \\NaI\\ to dust densities.  This is consistent with more straightforward calibrations of the equivalent widths, which imply $\\ebv \\ga 0.85$, but it could be considerably larger dependent on the degree of saturation.  Hence even with well resolved observations of the \\nad\\ lines it is impossible to derive a precise reddening from these lines and published estimates should be viewed with caution.  We have examined the implied outburst and quiescent luminosities, and the nature of the companion, in the context of these proposed scenarios.  Neither a distance of 6\\,kpc, nor 15\\,kpc, presents difficulties in these respects.  In fact the larger distance simplifies matters somewhat, for example in explaining the non-detection of the companion (likely a late-type sub-giant) in quiescence.  Consequently we propose that a large distance, $\\sim15$\\,kpc, should be given consideration."
    },
    "0402/astro-ph0402314_arXiv.txt": {
        "abstract": "Clusters of galaxies are powerful cosmological probes, particularly if their masses can be determined. One possibility for mass determination is to study the cosmic microwave background (CMB) on small angular scales and observe deviations from a pure gradient due to lensing of massive clusters. I show that, neglecting contamination, this technique has the power to determine cluster masses very accurately, in agreement with estimates by Seljak and Zaldarriaga (1999). However, the intrinsic small scale structure of the CMB significantly degrades this power. The resulting mass constraints are useless unless one imposes a prior on the concentration parameter $c$. With even a modest prior on $c$, an ambitious CMB experiment ($0.5'$ resolution and $1\\mu K$ per pixel) could determine masses of high redshift ($z>0.5$) clusters with $\\sim 30\\% $ accuracy. ",
        "introduction": "Clusters of galaxies are powerful cosmological probes. The abundance of clusters is very sensitive to the amplitude of fluctuations in the mass density field. For example, the present-day mass distribution of clusters offers perhaps the cleanest way to measure the normalization of the matter power spectrum~\\cite{Eke1996,Henry1997, ViaLid1999,ReiBoh2002,Ikebe2002,Schuecker:2002ti,Pierpaoli2003,Allen2003,Viana2003,Dodelson,Rozo2004}. As we gain the ability to probe the cluster abundance at high redshifts, we can hope to measure the evolution of this normalization, an evolution which is sensitive to the dark energy and its equation of state~\\cite{HMH2001,Wang:1998gt,Schuecker:2002yj}.   Because the mass function varies so rapidly with mass, accurate mass determination is critical if cluster constraints are to realize their potential. Finding clusters is not enough: even if a method finds all clusters and returns no false positives, it still does not necessarily have constraining power. It is quite possible that one technique will be optimal for finding clusters while another is a more powerful probe of cluster masses. For example, radio surveys which can measure the Sunyaev-Zel'dovich effect are very promising ways of detecting clusters even at high redshift~\\cite{White:2003af}. But they do not necessarily yield an accurate mass estimator. In contrast, weak gravitational lensing has drawbacks as a cluster finder~\\cite{MetzWhi2001,Whitevan2002} but is potentially a powerful way of measuring mass~\\cite{KaiSqu1993,SquKai1996,Hoekstra2001,Hoekstra2003, Schneider2003,Dod2003}.  While lensing of background galaxies has been studied in depth over the last decade, lensing of the cosmic microwave background (CMB) has not been explored as carefully, mainly because the angular resolution needed to see the predicted signal is only now being attained or planned~\\cite{ACT,ALMA,Stark1998}. Here I discuss how accurately cluster masses can be determined with observations of small angle ($\\sim 1'$) observations of the CMB.  On these very small scales, to a first approximation the background primordial CMB is a gradient. The photons passing close to the center of the cluster get deflected so that they appear to be originating further away from the cluster center. On the cool side of the gradient, this means the temperature is slightly hotter than it would appear without lensing, and on the hot side slightly cooler. So, subtracting off the dipole, the lensing pattern is quite distinctive~\\cite{Seljak:1999zn,Zal2000}, as shown in Fig.~\\rf{pattern}.  \\Sfig{pattern}{Idealized temperature pattern due to lensing of a cluster. Here, the background CMB is assumed to be a pure gradient with strength $T_y=13\\mu$K arcmin$^{-1}$ and this gradient has been removed. The cluster, at redshift 1 with a mass $\\mtwo=10^{15} h^{-1} M_\\odot$, has an NFW profile. The area shown is within the virial radius of $3.3'$. Jaggedness at the edge of the map is an indication of finite size $(0.5')$ of the (noiseless) pixels.}  Although there are many possible contaminants to this signal, its distinctive morphology allows one to hope that the signal can be extracted from any sources of noise~\\cite{Seljak:1999zn,HolKos,Vale2004}. Here I examine how accurately CMB lensing can determine the masses of clusters, accounting for the small scale structure present in the CMB, i.e., the fact that the background CMB is {\\it not} a pure gradient. There are two sources of this small scale structure: one is the power which remains in the damping tail of the primordial ($z\\sim1100$) spectrum and the other is lensing by structures along the line of sight. Figure~\\rf{cldt} gives an indication of the strength of the signal and the contributions from the CMB, which I will call {\\it CMB noise}. The CMB noise dwarfs the signal on ``large'' scales. We will see that this significantly impairs the constraining power of CMB-cluster lensing.   \\Sfig{cldt}{Anticipated signal from the lensing of a massive cluster and noise in $l$-space. The cluster has mass $10^{15} h^{-1} M_\\odot$ and is at redshift $1$. The light CMB curve is without lensing along the line of sight; the heavy CMB line includes this effect. The instrumental noise corresponds to a $0.5'$ resolution experiment with $1\\mu$K noise per pixel.}  Section II describes the model for the mass density I will use throughout and the ensuing constraints if CMB noise were absent. Section III shows how CMB noise pollutes the signal and which modes retain their power. Figures~\\rf{cont_cmb1d.5}-\\rf{dncontcijp_m10z1} summarize the main results of the paper and can be understood without reading the text. Finally, in the conclusion, I emphasize some of the simplifications made throughout. ",
        "conclusions": "Since gravitational lensing is directly sensitive to the mass of a cluster, it seduces us into hoping that we can reject more traditional phenomenological mass indicators such as X-ray temperature and galaxy counts. In the case considered here, lensing of the CMB, unavoidable realities spoil some of this allure. Since the primordial CMB does have some small scale structure, this {\\it CMB noise} significantly degrades the constraining potential of the lensed field. The constraints may still prove useful in constraining cosmological parameters such as the dark energy equation of state, but we will have to include this degradation when evaluating different experimental proposals.  While I believe this work is an important step on the road to assessing the utility of CMB-cluster lensing, it is incomplete in several ways. First, I have neglected a number of other contaminants such as point sources, the thermal Sunyaev-Zel'dovich effect, and the kinetic Zel'dovich effect. It is at least conceivable that these other effects differ sufficiently in frequency dependence and spatial morphology from the cluster lensing signal that they will not further degrade the projected constraints much. For example, if the morphology of the kinetic Sunyaev-Zel'dovich effect is much different from that depicted in the right panel of Fig.~\\rf{u17}, then it would have no impact on the mass determination, for we have seen that this is the mode which most tightly constraints the mass. Second, I have assumed a particularly simple parametrized form for the cluster mass distribution. Accounting for asphericity and substructure requires numerical simulations~\\cite{HolKos,Vale2004}, and these effects might add further noise to the mass estimators. Note that all of these effects would serve only to degrade the mass constraints further.  To close on a positive note, I point out that CMB experiments have consistently exceeded our expectations. There is thus every reason to believe that the high resolution, low noise experiments anticipated here will take place. If they do, then -- along with input from numerical simulations such as the correlation between mass and concentration -- we can hope to obtain interesting limits on the masses of thousands of clusters from CMB maps.  {\\it Note:} After this work was completed (but before it was circulated), Holder and Kosowsky~\\cite{HolKos} and Vale, Amblard, and White~\\cite{Vale2004} put out preprints studying very similar issues. Although we reached the same general conclusion -- realistic effects degrade the ability of CMB-cluster lensing to determine masses -- our papers are by and large complementary. While I have used analytic techniques and therefore simplified cluster models, the other two papers use more realistic numerical simulations. The simplicity has allowed me to provide quantitative measures of the efficiency of CMB-cluster lensing for a wide range of cluster masses, redshifts, noise levels, and pixel sizes. Ultimately, these estimates will need to be reinforced with the realistic simulations performed in \\cite{HolKos,Vale2004}."
    },
    "0402/astro-ph0402587_arXiv.txt": {
        "abstract": "{ The BATSE mission aboard CGRO monitored the whole sky in the 20 keV - 1 MeV energy band continuously from April 1991 until June 2000.  Although BATSE had very poor intrinsic angular resolution, the data can be used to survey the entire soft gamma-ray sky with $<1^{\\circ}$ angular resolution and $\\sim$mCrab sensitivity by using the Earth occultation method.  This method determines flux by measuring the step in the count rate profile in each BATSE detector as a source rises above or sets below the Earth's limb.  A maximum likelihood imaging technique can then be used to build up all-sky maps from the images of the Earth's limb produced by occulting sources.  However, since the Earth seen from BATSE has a radius of $\\sim$70$^{\\circ}$, the limb images that intersect at the positions of bright point sources have a significant effect over the area of the all-sky map.  A method for performing image cleaning on likelihood data has also been developed and is used to effectively remove artefacts from the all-sky maps.  This paper describes the 'LIMBO' imaging technique and presents preliminary all-sky maps of 25 - 160 keV emission, the first to be made since the HEAO1-A4 mission of 1978-79. ",
        "introduction": "Conventional imaging techniques, of focusing light with reflecting or refracting surfaces, are not possible in the hard X-ray or gamma-ray regions of the electromagnetic spectrum due to the highly penetrating nature of these photons (however, see \\citet{ballmoos} for a report on the first use of a Bragg diffraction gamma-ray lens in astrophysics).  For photons above $\\sim 10$ keV, where grazing incidence techniques become highly inefficient, it is necessary to make images by transforming the spatial information of the object onto the detector plane with some sort of modulation or coding technique that can be used to form a reconstructed image.  The ability to form an image has many advantages over simple measurements of source fluxes made with pointings of basic collimated instruments.  The ability to locate a source with an image allows comparison with measurements made at different times, with different instruments at different wavelengths.  Crowded fields can be resolved, minimising the possibilty of source confusion and the nature of extended or diffuse emission, such as that which may be seen from Supernova Remnants (SNR) or the Galactic Plane (GP), can be investigated.  Imaging detectors can also have a large gain in sensitivity over simpler instruments by reducing the amount of time needed to make accurate observations;  It is possible to make an estimation of the background in the region of a source at the same time as the observation, negating the need to perform control observations away from the source.  In an instrument with a large Field of View (FOV) it is possible to observe many sources in a single observation, which is often referred to as the \\emph{multiplex advantage}.  \\begin{figure*}[ht] \\centering \\includegraphics[width=17cm]{figs/shaw0026f1.eps} \\caption{The Compton Gamma-Ray Observatory and one of the 8 BATSE detector modules (drawings courtesy of the National Space Science and Technology Centre, Huntsville, AL, USA).} \\label{fig:cgro} \\end{figure*}  The natural progression of the last point above is to image the entire celestial sphere and perform an all-sky survey.  Knowing the locations and properties of a large number of sources is important in all energy bands of astronomy if only for the purposes of design, planning and proposal justification of other instruments and observations.  Even a survey with crude angular or spectral resolution can be highly important; as long as the sources in the survey catalogue have well determined uncertainties then other observations by higher resolution instruments are possible.  Contemporaneous measurements of a large number of sources are important in conducting morphological studies of families of objects (e.g. \\citet{seyferts}).  This is possible for an instrument with a large FOV, which is able to conduct a survey of such sources free from the biases that may be introduced by conducting many observations on different smaller FOV instruments at different times.  Similarly, by observing a large part of the sky, it is possible to maximise the probabilty of detecting unknown or transient objects, such as Gamma-Ray Bursts (GRBs) or X-ray novae, in a fixed time.  The nature of hard X-ray and gamma-ray imaging systems means that they often have much larger FOVs than is possible at optical wavelengths.  In the extreme case it is possible to observe the entire celestial sphere at once.  In high energy astronomy all-sky images have been provided between 0.1 - 2 keV by ROSAT \\citep{rosat} and $>$ 1 MeV by COMPTEL \\citep{comptelal26, comptelcat}, but there is sparse information in the intervening band, provided almost exclusively by the HEAO1-A4 experiment of 1978-79 \\citep{HEAO1A4}.  This experiment consited of two NaI/CsI phoswich detectors arranged so that they observed simultaneously in opposite directions and spun about an axis perpendicular to the optic axis and in the plane of the ecliptic.  This meant that it took 6 months for the whole sky to be observed and so most areas of the sky were only sampled two or three times during the 16 month mission.  The HEAO1-A4 survey catalogue contains 72 sources in four broad energy channels ranging from 13 - 180 keV, at a flux sensitivity of 14 mCrab.    The A4 experiment suffered from some problems with changing efficiencies between the two detectors and the sensitivity in the high energy band was limited; only 14 of the sources were detected at energies greater than $80$ keV.  Although the HEAO1-A4 was an important mission the usefulness of the survey is limited by the low spectral range and resolution and the fact that many transient sources were missed because of the slow sampling of the whole sky.  The aim of this work is to show how the wealth of the BATSE data set, which comprises 9 uninterrupted years of all-sky monitoring, can be exploited to achieve the maximum scientific gain. ",
        "conclusions": "The first true all-sky image of the gamma-ray sky since HEAO1-A4 has been produced, in a single 25 - 160 keV energy band, using the LIMBO maximum likelihood imaging technique with BATSE occultation data.  In order to produce the all-sky map it has been necessary to develop several new, novel techniques, such as: An accurate background model for the BATSE experiment (BAMM; \\citet{bammpaper}); Methods for producing all-sky images from BATSE occultion data; A method for cleaning artefacts from maximum likelihood images.  The all-sky map is based on $\\sim$ 500 days of data and shows the location of about 25 known gamma-ray sources, including several extragalactic objects at $4-5 \\sigma$ level ($\\sim$ 15 mCrab).  In almost all cases, LIMBO has identified emission from within one pixel of a source from the ISDC high energy catalogue.  However, whether all the detections from the LIMBO map are real sources of gamma-ray emission, that have been associated with the correct object from the catalogue, can not be determined at present.  It is clear that the images presented here suffer from the coarse $2^{\\circ}$ resolution of the imaging grid, which makes identification of sources in crowded regions, such as the Galactic Centre, very difficult.  The performance of LIMBO in these cases can be improved significantly by simply increasing the resolution of the imaging grid towards the theoretical imaging limit of $0.5^{\\circ}$, although this would be at the expense of an increase in computing time.  Since a single energy band has been considered, no measurements have been made of the spectra of any of the detected sources and fluxes are given only in comparison to the detected Crab Nebula flux.  The all-sky map was made by summing the images from 7 CONT energy channels.  Tests have shown that there is significant imaging response in the individual energy channels and further information could be gained from summing energy channels above 160 keV.  It is envisaged that further investigation will be made of the spectral performance of LIMBO with the aim of calibrating the response so that the continuum spectra of gamma-ray sources may be examined.  A rough estimate of the statistical sensitivity of the LIMBO technique to persisitent gamma-ray sources can be found by scaling the detection significance of the Crab nebula (990$\\sigma$) in the 500 day image.  This gives a 5$\\sigma$ sensitivity in the 25-160 keV range of $\\sim$5 mCrab.  If the whole BATSE CONT dataset of $\\sim$3300 days is considered then the sensitivity is improved by the square root of the ratio of the observation times, leading to a sensitivity of $\\sim$2 mCrab.  Depending on the brightness of the source, it should also be possible to make time resolved spectra for the duration of the BATSE mission.  Assuming an equal contribution to the integrated flux from each of the 7 CONT energy channels and using the same simple statistical arguments as above, it should be possible to make broad spectra for gamma-ray sources with a constant flux of $\\sim$6 mCrab.  Work continues to produce a publically accessible database of daily all-sky maps at $0.5^{\\circ}$ resolution for several energy channels of the full 9 years of available BATSE data."
    },
    "0402/astro-ph0402252_arXiv.txt": {
        "abstract": "Most Tidal Dwarf Galaxies (TDGs) so-far discussed in the literature may be considered as young ones or even newborns, as they are still physically  linked to their parent galaxies by an umbilical cord: the tidal tail at the tip of which they are usually observed. Old Tidal Dwarf Galaxies, completely detached from their progenitors, are still to be found. Using N--body numerical simulations, we have shown that tidal objects as massive  as $10^{9}$ solar masses may be formed in  interacting systems and survive for more than one Gyr. Old TDGs should hence exist in the Universe. They may be identified looking at a peculiarity  of their  ``genetic identity card\": a relatively high abundance in heavy elements, inherited from their parent galaxies. Finally, using this technique, we revisit the dwarf galaxies in the local Universe trying to find arguments pro and con a tidal origin. \\noindent ",
        "introduction": "The presence of compact star-forming regions in the tidal tails of colliding galaxies is commonly observed. The formation, in that environment,   of super-star clusters with  masses up to those of globular clusters has also often been reported (e.g., Knierman et al., 2003). The existence of even more massive tidal objects, with global properties characteristics of dwarf galaxies, has been claimed for more than a decade (see the  review by Duc \\& Mirabel, 1999). But whether such ``Tidal Dwarf Galaxies\" are genuine galaxies -- i.e.  they are  gravitationally bound entities -- is still strongly debated (see in this volume the contribution by Hibbard \\& Barnes  and  that of Amram et al. and Braine et al. for an alternative view). But  perhaps the more important issue of whether  TDGs survive, escape from their parent galaxies, and significantly  contribute to the overall population of dwarf galaxies, is  a matter of speculation. Indeed, a more  plausible fate is their tidal destruction or their falling back onto their progenitors. There are two ways to tackle the problem: using numerical simulations of  interacting galaxies and following in them the buildup and evolution of tidal objects, or trying to identify in catalogs of  dwarf  galaxies  those which might have a tidal origin. We have explored both approaches. ",
        "conclusions": ""
    },
    "0402/astro-ph0402246_arXiv.txt": {
        "abstract": "We report centimeter continuum observations with the VLA and the VLBA as well as mid-infrared observations with COMICS/SUBARU toward the components of the YLW~15 very young binary system, VLA~1 and VLA~2.  The centimeter emission of the two components traces partially thick free-free emission, likely due to collimated, ionized winds. VLA~1 is an embedded protostar, undetected in the near-IR, and possibly in the Class 0 to Class I transition and powering a Herbig-Haro outflow. Its mid-IR emission appears slightly resolved with a diameter of $\\sim 16$~AU, possibly tracing circumstellar material from both the envelope and the disk. VLA~2 is a typical Class I object, unresolved in the mid-IR, and is the responsible of the strong X-ray emission associated with YLW~15. The expected centimeter ''peri-stellar'' emission associated with the X-ray emission is not detected with the VLBA at 6~cm likely due to the high optical depth of the free-free emission.  Strikingly, the near to mid-IR properties of YLW~15 suggest that VLA~1 is a more embedded YSO, or alternatively, less luminous than VLA 2, whereas orbital proper motions of this binary system indicate that VLA 1 is more massive than VLA 2. This result is apparently against the expected evolutionary scenario, where one expects that the more massive YSO in a binary system is the more evolved and luminous YSO. Finally, the nearby source YLW~16A is detected with the VLA, its position coincides with reported near-IR and X-ray sources. ",
        "introduction": "}  The $\\rho$ Ophiuchi molecular cloud complex is one of the nearest star formation regions ($\\sim$ 120~pc: Knude \\& Hog 1998) and is also one of the best studied low mass star forming regions.  YLW~15 (IRS 43) is a young stellar object (YSO) embedded in the Oph-F dense core, located in the southeastern region of the molecular cloud complex (\\eg\\ Motte, Andr\\'e \\& Neri 1998).  The near-IR to mid-IR (\\eg\\ Wilking, Lada, \\& Young 1989), far-IR (\\eg\\ Bontemps \\et\\ 2001; Pezzuto \\et\\ 2002;  Nisini, Gianni \\& Lorenzetti 2002) and millimeter (\\eg\\ Andr\\'e \\& Montmerle 1994; Bontemps \\et\\ 1996) properties suggest that YLW~15 belongs to the Class I stage of evolution.  The spectroscopic features in the near-IR show that YLW~15 is a YSO  with K5 spectral type and with dwarf like surface gravity, \\ie\\ a luminosity class IV/V, having a stellar mass of 0.5~\\mo\\ (Greene \\& Lada 2002).   It has been proposed that YLW~15 constitutes a wide-binary system with GY~263, located 6$''$ (720~AU) northwest of YLW~15 (\\eg\\ Allen \\et\\ 2002).  YLW~15 powers a compact molecular outflow (Bontemps \\et\\ 1996) and a Herbig-Haro outflow (\\eg\\ Grosso \\et\\ 2001).  YLW~15 is one of the youngest low mass stars associated with X-ray emission. First detected by ROSAT (Casanova \\et\\ 1995), the X-ray emission shows quasi-periodic flares (Tsuboi \\et\\ 2000).  A ``super-flare''  was observed during a few hours (Grosso \\et\\ 1997).   Grosso \\et\\ (1997) suggested that the flare arises from a magnetically confined, low-density plasma bubble with a diameter of $~\\sim0.05$--0.3~AU around the young star.   The quasi-periodic X-ray flares of YLW~15 were explained by Montmerle \\et\\ (2000) as due to ``fast rotation of the central star with respect to the accretion disk, which results in star-disk shearing of the magnetic lines, producing magnetic reconnection and mass loss, and eventually extremely high X-ray luminosities'', which is different from the ``classical'' quasi steady-state situation expected for the T Tauri stars.  Its radio continuum emission has been well studied at a moderate angular resolution ($\\sim 10''$) showing no variability (\\eg\\ Leous \\et\\ 1991).  VLBI milliarcsecond angular resolution observations  of YLW~15 show that the emission is resolved out at ``peristellar'' scales  (\\ie\\  its size is larger than $1.2\\times10^{13}$~cm), suggesting that the  emission is thermal in origin and probably comes from circumstellar ionized  winds (Andr\\'e \\et\\ 1992). Sub-arcsecond angular resolution VLA observations revealed that the radio emission arises from a binary system, YLW~15~VLA~1 and YLW~15~VLA~2 (Girart, \\ro\\ \\& Curiel 2000, Paper I).  In these maps VLA~1 was clearly elongated with a position angle of PA$=25\\arcdeg \\pm 2\\arcdeg$.  This elongation aligns well with HH~224NW1 and a chain of near-IR HH knots (Grosso \\et\\ 2001).  VLA~2 appeared unresolved down to scales of $\\sim 0\\farcs4$.  In Paper I we suggested  that the two radio sources trace a binary young stellar system with each component possibly having a different evolutionary status.  Using a 12 year time baseline of subarcsecond angular resolution VLA observations, Curiel \\et\\ (2003, Paper II) have measured absolute proper motions of 24~mas~yr$^{-1}$ in YLW~15.  In addition, the VLA~1 and VLA~2 relative astrometry analysis done in Paper II reveals orbital proper motions that suggest a total mass of the binary system of $\\ga$1.7~\\mo.  These results show also that VLA~1 is likely more massive than VLA~2.   In this paper we present VLA sub-arcsecond angular resolution observations at 3.6 and 6~cm toward YLW~15, as well as VLBA 6~cm and COMICS/SUBARU 8.8~$\\mu$m observations.  In \\S~2 and \\S~3 we give the details of the observations and results with the VLA, VLBA and SUBARU. In \\S~4 we discuss the nature of the radio emission and compare its properties with those obtained at other wavelengths in order to elucidate the nature of this interesting binary protostellar system. ",
        "conclusions": "We have carried out VLA and VLBA radio as well as COMICS/SUBARU mid-IR (8.8~$\\mu$m)  observations of the YSO binary system in YLW~15 and carried out a study of the properties of YLW~15 from previous results available in the literature. The main conclusions are:  \\begin{enumerate} \\item The properties of the radio continuum emission of VLA~1 and VLA~2 show that these two sources have partially thick free-free emission, likely due to collimated, ionized winds, as it is the case of other Class I and 0 YSOs, with associated molecular and/or HH outflows. \\item The centimeter VLA~1 emission has a jet-like morphology, with the axis coinciding in position angle with the HH outflow found by Grosso et al. (2001) and a jet length of $\\sim 0\\farcs32$ or $\\sim 40$~AU. Undetected in the near-IR, becomes ''visible'' at 8.8~$\\mu$m. The 8.8~$\\mu$m emission is slightly resolved, with a deconvolved size of $\\sim 16$~AU, suggesting that the emission arises from a compact circumstellar structure around the protostar. \\item The centimeter VLA~2 emission is partially resolved at scales of 20~AU. Comparison with data taken at other wavelengths shows that VLA~2 is the responsible of the near-IR emission and it is the source of the strong X-ray emission.  This source appears unresolved at 8.8~$\\mu$m, implying a small circumstellar structure. \\item VLBA observations failed to detect continuum emission at 6~cm. The lack of ''peri-stellar'' emission is likely due to a high optical depth of the free-free emission, that arises at larger scales. Shorter wavelengths should be used to detect the expected centimeter emission associated with the X-ray emission. \\item The near to mid-IR properties of YLW~15 show that VLA~1 is less luminous than VLA~2 in this wavelength range. This could be due to a lower bolumetric luminosity of VLA~1 compared with VLA~2 or alternatively indicate that VLA~1 is a more embedded YSO, and thus, younger or less evolved, than VLA~2. VLA~2 has the SED properties of a typical Class I YSO. Therefore, VLA~1 could be either in the same evolutionary state than VLA~2 or either in the Class 0 to Class I transition. \\item VLA~1 appears to be more massive but more embedded or, alternatively, less luminous than VLA~2. This result is apparently against the expected evolutionary scenario, where one expects that the more massive YSO in a binary system is also the more evolved and luminous YSO. \\item YLW~16A is detected with the VLA. Its position coincides well with the near-IR and X-ray emission observed in the region. \\end{enumerate}"
    },
    "0402/astro-ph0402593_arXiv.txt": {
        "abstract": "We suggest that large \\ion{H}{i} holes observed in the interstellar medium (ISM) of galaxies such as the Large Magellanic Cloud (LMC), the Small Magellanic Cloud (SMC) and \\objectname{Holmberg II} (Ho II, DDO 50, UGC 4305) can form as the combined result of turbulence coupled with thermal and gravitational instabilities. We investigate this problem with three dimensional hydrodynamical simulations, taking into account cooling and heating processes and the action of the self-gravity of the gas. We construct an algorithm for radiative transfer to post-process the simulated data and build emission maps in the 21 cm neutral hydrogen line. With this approach, we are able to reproduce the structure of the shells and holes as observed in regions of the ISM, where no stellar activity is detected. In order to quantify the comparison of our synthetic maps to the observations, we calculate the physical scale-autocorrelation length relation ($L-L_{cr}$ relation) both on the synthetic \\ion{H}{i} maps and the \\ion{H}{i} map of Ho II. The $L-L_{cr}$ relation shows a linear increase of the autocorrelation length with the physical scale up to the scale of energy injection and flattens for larger scales. The comparison of the $L-L_{cr}$ relation between the observations and the synthetic maps suggests that turbulence is driven in the ISM of Ho II on large scales ($\\sim~6$ kpc). The slope of the $L-L_{cr}$ relation in the linear regime in Ho II is better reproduced by models where turbulence is coupled  with a weak efficiency cooling of the gas. These results demonstrate the importance of the interplay between turbulence and the thermodynamics of the gas for structure formation in the ISM. Our analysis can be used to determine the scale on which kinetic energy is injected in the ISM of dwarf irregular galaxies, and to derive, in a first approximation, the cooling rate of the gas. ",
        "introduction": "The question of which physical processes produce the holes observed in the 21 cm \\ion{H}{i} line in the interstellar medium (ISM) of many galaxies such as in the Large Magellanic Cloud (LMC) (Kim et al 1999), Holmberg II (Ho II, DDO 50, UGC 4305) (Puche et al. 1992), the Small Magellanic Cloud (SMC) (Staveley-Smith et al. 1997, Stanimirovi\\'c et al. 1999) and the Western Magellanic Bridge (Muller et al. 2003) has been a controversial and puzzling issue in the past. On the observational side, there are evidences that many of the shell and void-like structures observed in the ISM of different galaxies, especially in the outer parts of their galactic disks are unlikely to be the result of supernova explosions. Kim et al. (1999) came to this conclusion for the LMC in view of the weak correlation that exists between the positions of some of the \\ion{H}{i} shells and the \\ion{H}{ii} regions. Using a multicolor BVR imaging approach, a careful study by Rhode et al. (1999) of the 51 \\ion{H}{i} holes in Ho II, previously selected by Puche et al. (1992), shows no signs of the presence of remnant stellar clusters inside 86 percent of these holes. Additionally, using H$\\alpha$ photometry, they could not confirm the presence of any ongoing star formation process inside the holes. Based on these observations, they concluded that those holes can not be supernova driven expanding shells. Furthermore, Stewart et al. (2000) observed the far-ultraviolet (FUV) emission of Ho II. The FUV emission is sensitive to massive star formation on a timescale which is comparable to the lifetime of the largest holes in Ho II ($\\sim$ 100 Myrs). Only 3 out of the 51 \\ion{H}{i} holes selected by Puche et al. (1992) have associated FUV emission whereas in the rest of the holes the emission is concentrated at their edges. Furthermore, only very little FUV emission is detected outside the 2$\\arcmin$-3$\\arcmin$ radii inner region, indicating that star formation has played only a little role in the outer region. X-ray observations of Ho II using ROSAT PSPC by Zezas et al. (1999) and Kerp et al. (2002) have shown that most of the \\ion{H}{i} holes are devoid from hot gas. However, Kerp et al. (2002) showed that a few X-ray sources (sources 2 and 4 in Fig. 1 in their article), located outside the Stellar Disk but within the \\ion{H}{i} distribution of Ho II have emission that might be associated with thermal plasma radiation which is characteristic of supernova remnants. This suggests that, the supernova scenario can not be entirely ruled out in the case of those few holes.  Alternative mechanisms which could lead to the creation of the \\ion{H}{i} holes in Ho II have been discussed by several authors (Rhode et al. 1999; Bureau \\& Carignan 2002). Some of the proposed scenarios are related to the supernova scenario, like invoking a non standard initial mass function and modifying the holes energetics which helps to reduce the number of supernovae needed to create them. The latter scenario was pointed out by Rhode et al. (1999). It would however require a reduction of the supernova rate by a factor of $\\sim 5$ in order to have a remnant stellar cluster with a luminosity that falls below their detection limit. This in turn would require that the expansion velocities in the case of the largest holes have been overestimated by a factor $\\sim 3$ by Puche et al. (1992), which seems unlikely. The scenario of an extremely top-heavy IMF in Ho II is also not supported by any observations of significant variations in the IMF measured in different environments like the Milky Way or other nearby galaxies (Leitherer 1999) and thus, seems also very unlikely. \\ion{H}{i} holes could also be created by the impact of High-Velocity Clouds (HVCs) on the gaseous disk. In the case of the Milky Way and other galaxies, it has been shown that some holes are associated with high velocity gas (Heiles 1984 for the Milky Way; Van der Hulst \\& Sancisi 1988 for M101). In the case of Ho II, only one object has been identified by Rhode et al. (1999) as a possible HVC. This object has an estimated mass of 1.2$\\times$10$^{7}$ M$_{\\odot}$ and a total kinetic energy of $\\approx 5 \\times 10^{53}$ ergs. This amount of energy is enough to create any of the observed \\ion{H}{i} holes in Ho II. Aside from being statistically not very significant, this cloud has been detected on only one pixel in the whole data set. As already pointed out by Rhode et al. (1999), the signal detected from that pixel could also be assimilated to noise under the assumption that the noise in the data set is Gaussian. Only a more sensitive study of  the \\ion{H}{i} structure around Ho II can shed more light on the existence and role of HVCs. By re-processing the low resolution map of Puche et al. (1992), Bureau \\& Carignan (2002) have unveiled the large scale structure of the \\ion{H}{i} in Ho II to be comet-like. They interpret this as gas in Ho II experiencing ram-pressure stripping under the effect of the intra-group medium (IGM) and that this process might lead to the formation of some holes or to the enlargement of existing ones. This scenario would only work if the IGM density is $n_{IGM}\\gtrsim 4.0\\times 10^{-6}$ cm$^{-3}$. Bureau \\& Carignan (2002) pointed out to the fact that tidal interactions on the scale of the subgroup of galaxies containing Ho II, Kar 52 and UGC 4483 could also affect the \\ion{H}{i} morphology of Ho II. However, a preliminary analysis of the \\ion{H}{i} structure on the subgroup scale by Bureau et al. (2004) does not support the scenario of gas being tidally stripped from the disk of Ho II in the direction of Kar 52 and vice versa. It is also inconclusive with regard to the role played by ram pressure stripping.  In this work, we investigate whether the \\ion{H}{i} morphology could also emerge as the result of internal gas dynamics as earlier suggested by Elmegreen (1997). Our aim is to explore the role played by turbulence and by thermal (Field 1965; Field et al. 1969)  and gravitational instabilities (Kolesnik 1991) (TI and GI, respectively) on the formation of large scale structure in the ISM. We use three dimensional hydrodynamical isothermal simulations of driven turbulence to investigate if holes like those in Ho II are formed as a natural consequence of the turbulence present in the medium. We then study the effect of heating and cooling and take the self-gravity of the gas into account. In a post-process approach, we build two dimensional emission maps in the  21 cm \\ion{H}{i} hyperfine structure line and compare them to observations.  TI has been traditionally considered as an important agent of structure formation in the ISM. Burkert \\& Lin (2000) performed a linear perturbation analysis of TI and studied numerically under which condition the transition from the linear to the non-linear regime occurs. S\\'anchez-Salcedo et al. (2002), used high resolution 1D simulations to study the behavior of TI in the presence of forced flows. Koyama \\& Inutsuka (2000, 2002) have investigated with 1D and 2D simulations which include detailed cooling, heating and chemistry, the propagation of shocks in the warm neutral medium (WNM) and cold neutral medium (CNM) and the development of TI in the post-shock region. Kritsuk \\& Norman (2002a,b) showed, using three dimensional simulations, that a time fluctuating background heating supports turbulence in a multiphase ISM. Piontek \\& Ostriker (2004) investigated the interplay between TI and the magnetorotational instability (Balbus \\& Hawley 2001). V\\'azquez-Semadeni et al. (2000) and Gazol et al. (2001) used local two dimensional numerical simulations to investigate TI in the ISM on a scale similar to ours (1 kpc). They focused on the segregation of the medium into two distinct gaseous phases. A major success of their model is that they were able to reproduce a mass fraction of the unstable gas which is in good agreement with the observations (Heiles 2001), but they did not build emission maps that can be compared with observations. Wada \\& Norman (1999, 2001) and Wada et al. (2000) have shown that the combined effect of TI and GI with turbulence can lead to the formation of \\ion{H}{i} holes. They performed global two dimensional numerical simulations of an LMC type galactic disk and constructed from their hydrodynamical data emission maps in the 21 cm \\ion{H}{i}  and the CO(1-0) lines. They transformed their two dimensional grid into a three dimensional one by assuming a constant scale height for the disk, divided in a number of grid cells. Though they used a rather sophisticated approach to solve the radiation equation in each of the newly created cells, the local densities in all cells along a given line of sight were assumed to be the same, which is a major simplification. Additionally, although a similarity in the complex network of holes and shells is observed between their simulated \\ion{H}{i} map and observations of the structure in the 21 cm line of the LMC, they did not quantify the hole and shell-like structures. This contrasts with our simulations which take into account the complexity of the three dimensional structure, which arises self-consistently. The \\ion{H}{i} holes formation process by TI and GI in the ISM of dwarf irregular galaxies is critically reviewed by S\\'anchez-Salcedo (2001). The paper is organized as follows. In \\S 2 we describe our set of simulations. The radiative transfer approach we use to construct the emission maps is described in \\S 3. In \\S 4 we define the observable quantity which we use to compare our models with observations. In \\S 5 we present the observational data and in \\S 6 we evaluate the autocorrelation length in our simulated \\ion{H}{i} maps and compare it with the observations. In \\S 7, we summarize our results. ",
        "conclusions": "In this work we compare 21 cm \\ion{H}{i} simulated maps to the high resolution \\ion{H}{i} map of Puche et al. (1999) of the dwarf irregular, gas rich galaxy Holmberg II (Ho II). The maps are built by using a post processing of the radiative transfer equations on different models of 1 kpc scale local three dimensional hydrodynamical simulations. After mimicking the same spatial resolution in the simulated maps than the one used in the observations, we perform a quantitative comparison between the models and the observations using the autocorrelation length. The medium in the different models is driven in time and space and turbulence is charaterized by a particular injection scale $L_{driv}$. Simulations are either isothermal (Models ISO in Tab.~\\ref{tab1}) or subject to cooling and heating and might include the effect of self-gravity (models CO in Tab.~\\ref{tab1}). Cooling and heating are parametrized with simple functions of the form $\\Lambda(\\rho,T)=\\Lambda_{0}\\rho^{\\alpha}T^{\\beta}$ and $L(\\rho)=q\\Lambda_{0}\\rho^{\\epsilon}$, respectively. In the set of simulations discussed in this paper, we essentially vary the parameter which affects the most the efficiency of thermal instability (TI), namely the cooling rate $\\Lambda_{0}$. For the isothermal models driven at a Mach number $Ma=1$ which is characteristic of the warm phase of the ISM, the largest over-densities reached in the medium when it reaches dynamical equilibrium are of the order of 2-3 cm$^{-3}$, which are only a few times larger than the average density, $\\bar{n}=0.5$ cm$^{-3}$. This contrast to models which include cooling and heating where  TI leads to the formation of condensations with peak densities of a few 100 cm$^{-3}$. The latter condensations usually sit at the intersection of filament and sheet like structures which are separated by low density cavities. In the absence of gravity, the medium settles into a two-phase medium made of cold and dense filaments and warm low density holes and the probability distribution function evolves into a stationary bi-modal form (see Fig.~\\ref{fig2}). In the presence of self-gravity, the evolution of the medium is similar to the case without self-gravity until condensations of a few 100 cm$^{-3}$ are formed at $t=\\tau_{dyn}$, though these condensations have peak densities that are larger than their counterparts in the non-gravitating case. At a later stage of the evolution, the dense and gravity bound condensations acrrete the surrounding gas to become more massive until they reach the stage of condensation-condensation accretion. The evolution of the autocorrelation length at a given scale is independant from the presence of gravity until this stage of condensation-condensation merger is reached. This is simply due to the fact that, even though the presence of gravity leads to the formation of denser structures, it does not affect the spatial distribution of the condensations in the box in comparison to the non self-gravitating models until the phase of condensation-condensation accretion is reached.  The autocorrelation length-physical scale relation ($L-L_{cr}$ relation) shows, both in the observations and the simulations, two regimes. For scales smaller than the scale of the energy injection $L_{driv}$, $L_{cr}$ shows a linear increase with $L$ and flattens for scales larger than $L_{driv}$. This is confirmed by a comparison of the observations with synthetic maps derived from simulations driven at different length scales. The $L-L_{cr}$ relation of Ho II implies that energy is injected into the ISM of the galaxy on a scale of $\\sim 6$ kpc. This clearly favors a large-scale driving rather than a driving by supernova explosions. The kinetic energy power spectrum drawn by Stanimirovi\\'c \\& Lazarian (2001) for the Small Magellanic Cloud (SMC) speaks for a large scale driver operating in the SMC as well. The spectrum shows no signs of energy injection up to the largest scale they have considered (4 kpc). The slope of the $L-L_{cr}$ relation seems to be related to the gas physics and the mode of energy decay from $L_{driv}$ down to the smaller scales. Our results show that the slope of the $L-L_{cr}$ relation in the linear regime can be better reproduced when turbulence is coupled to a weak cooling of the gas. We have checked, in the case of model CO09G2 which is the best fitting model to the data, the effect of numerical resolution by performing two additional models at the resolutions of $64^{3}$ and $192^{3}$. We observed the time evolution of the autocorrelation length in the models with different resolutions and found that the three models yield relatively similar values, in particluar models with the resolutions of $128^{3}$ and $192^{3}$, whereas the model with the resolution of $64^{3}$ shows larger fluctuations.  Wada et al. (2000) have shown, in the case of the LMC, that thermal and gravitational instabilities lead to the formation of a network of shell and hole-like structures in the ISM, however, they did not perform a quantitative comparison between their models and the observations in terms of the \\ion{H}{i} structure. Our approach can be applied to the ISM in different galaxies in order to 1) constrain the global cooling rate of the galaxy which is connected to the slope of the $L-L_{cr}$ relation in the linear regime and most important, 2) estimate the physical scale on which the energy is injected into the medium by the position of the turning point in the $L-L_{cr}$ curve. In a subsequent paper, using 3D global models of dwarf irregular low star forming galaxies, we plan to investigate the relationship that may exist between the structure function index $\\alpha$ ($S(l)=l^{\\alpha}$, where $l$ is the physical scale) or the slope of the $L-L_{cr}$ relation in the linear regime for the \\ion{H}{i} gas and the global metallicity of the galaxy. Such a relation might exist for dwarf irregular galaxies with small star formation rates and where the gas morphology is likely to be shaped by the interplay of turbulence and the thermodynamical properties of the gas. By building a large library of models including also star formation and feedback, it might also be possible to uncover the relation that exists between the structure function index, the metallicity, and the star formation rate. It is not yet known whether a combination of the metallicity (or a radially dependent metallicity profile) and a radially dependent star formation rate can lead to a unique, non degenerate, large scale morphology of the gas. Establishing such relations between the gas morphology and metallicity, eventually with a radial dependence, would be useful for determining the metallicity of galaxies in a very simple way from their gas morphologies. Interestingly, a large sample of \\ion{H}{i} observations for $\\sim$ 40 galaxies has been recently obtained with the VLA by F. Walter (Heidelberg) and M. Bureau (Oxford), among them are $\\sim$ 10 gas rich dwarf irregular galaxies of different metallicities which can serve as the perfect observational sample to compare our models with.  Regarding the energy requirement of the driver, we have injected into our simulation box $\\sim 1.1\\times10^{50}$ ergs Myrs$^{-1}$ kpc$^{-3}$. Summing the kinetic energy derived by Puche et al. (1999) (based on kinetic energy input estimates using the model of Chevalier 1974), normalizing that energy input to the age of the oldest hole in Ho II (in time units of Myrs) and dividing by the volume of the galaxy yields a value of $\\sim7.3\\times10^{48}$ ergs Myrs$^{-1}$ kpc$^{-3}$. This implies that the supernova scenario applied to Ho II underestimates by a factor of $\\sim 15$ the required energy for the holes formation. The estimates of Puche et al. (1999) might underestimate the number of supernova explosions in each hole. This would lead to larger remnant stellar clusters than those already expected, but not found, in the \\ion{H}{i} holes of Ho II (Rhode et al. 1999). The energy balance argument tends to confirm our conclusions that supernova explosions are not responsible for the formation of all of the observed \\ion{H}{i} holes in Ho II.  The natural continuation of this work is to investigate what are the eventual large scale drivers operating in Ho II. In the absence of large scale supernova driving, spiral density waves could constitute a plausible driving candidate. However, Ho II does not show any developed spiral pattern. Another possibility is that Ho II, by infalling to the center of the M81 group and weakly interacting with its neighboring galaxies, might be undergoing a process of ram pressure stripping in the intra-group medium of this group (Bureau and Carignan 2002) and/or tidal stripping by it's closest neighbors. These scenarios could be confirmed by running combined N-Body/SPH simulations of interacting galaxies while including the intergalactic medium gas. Such simulations should be made possible by using new hybrid N-Body/SPH codes like VINE which is currently under development (Wetzstein et al. in preparation). Another possibility to drive turbulence on large scales could be offered by the magneto-rotational instabilities that could take place in the outer parts of the galactic disk of Ho II (Balbus and Hawley 1991; Selwood and Balbus 1999). However, this scenario has not yet been investigated in detail in a galactic context. We should finally mention that the scenario discussed in this paper does not necessarily rule out that \\ion{H}{i} holes can also be formed, in other dwarf galaxies, as a result of multiple supernova explosions as discussed by S\\'anchez-Salcedo (2002) in the case of the \\ion{H}{i} holes observed in the nearby dwarf galaxy IC 2574 (Walter and Brinks 1999)."
    },
    "0402/astro-ph0402070_arXiv.txt": {
        "abstract": "{SS433 is the prototype {\\it microquasar} in the Galaxy and may even be analogous to the ULX sources if the jets' kinetic energy is taken into account.  However, in spite of 20 years of study, our constraints on the nature of the binary system are extremely limited as a result of the difficulty of locating spectral features that can reveal the nature and motion of the mass donor. Newly acquired, high resolution blue spectra taken when the (precessing) disc is edge-on suggest that the binary is close to a common-envelope phase, and hence providing kinematic constraints is extremely difficult.  Nevertheless, we do find evidence for a massive donor, as expected for the inferred very high mass transfer rate, and we compare SS433's properties with those of Cyg X-3.}    \\addkeyword{accretion} \\addkeyword{accretion discs} \\addkeyword{X-rays: binaries} \\addkeyword{X-ray: stars}   \\begin{document} ",
        "introduction": "\\label{sec:intro} SS433 holds a special place in late 20th century astronomy as the first relativistic jet source discovered in the Galaxy (see Margon 1984), and hence is the prototype {\\it microquasar}.  It is remarkable as the only {\\it continuously} emitting microquasar, with its key feature being the 162.5d precession period of the jets (and associated accretion disc) that is revealed by the ``moving lines'' which have now been observed for more than 2 decades.  The radial velocity curves of the moving lines are very well described by the {\\it Kinematic Model} (Margon 1984), thereby revealing key parameters of the jets ($v$=0.26c; $i$=79$^\\circ$) and, combined with radio observations of the associated SNR W50, a distance of 4.8kpc.  It is observed as a weak ($\\sim$10$^{36}$erg~s$^{-1}$) X-ray source (X1909+048), which combined with the high orbital inclination (it is eclipsing) suggests that it may be an Accretion Disc Corona (ADC) type of X-ray source. However, it is well established that the kinetic energy in the jets is at least 1,000 times greater than the observed luminosity (Safi-Harb 2003), although the origin of the jet energy is still a matter of debate, and this would place SS433 in the recently defined class of ULX sources (e.g. King 2003). If the intrinsic luminosity were indeed $\\geq$10$^{39}$erg~s$^{-1}$, then the implied mass-transfer rate would be extremely high (\\mdot$\\sim$10$^{-4}$\\msun$y^{-1}$, see also Fuchs et al 2003).  This has led to King et al (2000) investigating the evolutionary status of SS433 and they show that, given a massive, radiative companion, it is possible for SS433 to avoid the common-envelope phase, by expelling the majority of the transferred material (part of which constitutes the jets).  They conclude that the observed X-rays are then likely to almost entirely arise from the jets themselves. Indeed, the X-ray output of SS433 appears to be dominated by the precessing jets, with moving X-ray emission lines and even spatially resolved X-ray structure along the jet axis (Marshall et al 2002; Namiki et al 2003). Recently, however, the RXTE All-Sky Monitor has revealed a long-term modulation in the precessional X-ray lightcurve (Gies et al 2001). The relativistic doppler boosting from jet precession alone is insufficient to produce this modulation, which requires an additional, presumed geometrical component (Clarkson et al 2004).  For the established high inclination this can most naturally be provided by structure in the edge-on view of the accretion disk, which may be tilted or warped to produce the apparent asymmetry in the precessional lightcurve.   \\begin{figure}[!t]\\centering \\vspace{0pt} \\includegraphics[angle=0,width=0.94\\columnwidth]{asm.ps} \\caption{7-year X-ray lightcurve of SS 433 obtained with the RXTE/ASM. The dotted and dashed curves represent modulation expected from doppler boosting in the relativistic jets (for the cases of bullets and a continuous stream respectively). An additional geometric modulation is clearly required, modelled here as the modulation of a power-law jet by a precessing accretion disk (dot-dashed line). From Clarkson et al (2004).} \\label{fig:asm} \\end{figure} ",
        "conclusions": ""
    },
    "0402/astro-ph0402300_arXiv.txt": {
        "abstract": "Some general properties of extensive air showers are discussed. The main focus is put on the longitudinal development, in particular the energy flow, and on the lateral distribution of different air shower components. The intention of the paper is to provide a basic introduction to the subject rather than a comprehensive review. ",
        "introduction": "Extensive air showers (EAS) are known for about 70 years to be cascades initiated by primary cosmic rays in the atmosphere. As such, they serve as connection to the highest particle energies nature is offering, especially at energies exceeding $10^{15}$~eV where direct measurements of cosmic rays are hampered by the low primary flux. EAS can be viewed as tools for astroparticle physics. The main task is to reconstruct from the shower observables the parameters of the initial cosmic ray, which is in general straightforward for the particle direction. It is more a challenge for the energy and much more for the particle type of the primary. This is due mostly to shower fluctuations, an important EAS property given by the stochastic nature of the elementary interaction processes. Combined with the steep primary spectrum, the fluctuations impose a particular challenge for interpreting air shower data. Use can be made, however, of the fact that an EAS consists of different shower components.  The main body of this paper deals with illustrating, supported by EAS simulations, some characteristics of these shower components and the information they contain about the primary particle. This might serve as introduction to the subject of extensive air showers and help to better understand what air shower experiments are measuring, which information they try to infer from the data, and where limitations are given~\\cite{proc_exp}. A more comprehensive introduction to EAS is given e.g.~in~\\cite{gaisser}.  \\begin{figure} \\centering \\epsfig{file=fig1.eps,width=0.73\\textwidth} \\caption{Schematic sketch of an EAS.} \\label{fig-scheme} \\end{figure}  A schematic sketch of an EAS is given in Figure~\\ref{fig-scheme}. The incident particle (primary energy $E_0$ and primary type $A_0$) hits an air nucleus. Usually considered as extreme illustrations of parimary particle types are protons and iron nuclei (although an analysis of EAS data is not restricted to this range). The notation $A_0$ anticipates the fact that EAS features of primary hadrons, which are known as dominant cosmic rays from direct measurements at energies below $10^{15}$~eV, show some dependence on the nucleon number $A$ of a nucleus, as explained later. In the first interaction, a number of secondary hadronic particles are generated. Mostly by decay of the neutral pions, electromagnetic sub-cascades are initiated, and by decay of (part of) the charged pions, muons and neutrinos are produced. While neutrinos are hardly detectable in classic air shower experiments, muons might reach the observation level even from large altitudes. It is important to note that a competition of the charged pions between decay and interaction exists which depends on the pion energy and the local air density. Some information on the first interactions will be imprinted in this way in the muon component. The remainder of the hadronic secondaries continues to interact with air in subsequent collisions and to feed the other shower components.  In ground arrays, the surviving particles are measured with particle detectors, and with appropriate optical telescopes Cherenkov and fluorescence emission during the shower propagation can be observed.  To characterize more quantitatively the interaction process, it is helpful to introduce the quantity {\\it inelasticity k} as the energy fraction available for the production of secondary particles (or in other words, the initial collision energy reduced by the energy of the most energetic particle). As a rule of thumb, about $\\frac{1}{3} k$ is ``lost'' to the electromagnetic channel per hadronic interaction. For a mean value $k \\simeq 0.6$ typical for high-energy interactions this corresponds to $\\simeq 20\\%$ per interaction that on average will give rise to subsequent electromagnetic cascading. Equipped with this knowledge, we can turn to the longitudinal shower development. ",
        "conclusions": "Extensive air showers consist of different particle components which have different shower characteristics. This gives a handle to determine primary energy and (to some extent) the primary particle type. Extensions of this approach that were not discussed comprise the exploitation of different time structures of the shower front or comparing inclined events (where mostly muons survive the increased atmospheric path-length) to near-vertical ones. Also in the hadronic shower component, information on the primary particle type is imprinted; but maybe even more important is the possibility to test high-energy hadronic interaction models used in the simulations. Shower simulations have grown to an important tool for EAS data reconstruction. Shower fluctuations motivate the development of Monte Carlo simulation techniques. In turn, these simulations might also be applied to find observables that depend less strongly on the specific primary or on shower fluctuations, such as the signal at distances of 600$-$1000~m from the core for primary energy estimations with ground arrays. At the highest energies, further interaction features have to be considered in shower simulations, e.g.~for primary photons the LPM effect and preshower formation in the geomagnetic field.  In summary, even after many years of extensive exploration, air showers continue to be fascinating tools for astroparticle physics.  {\\it Acknowledgements.} The conference organizers are thanked for their kind hospitality and for creating an inspiring conference atmosphere. In particular, I would like to thank Marek Je$\\dot{\\textnormal{z}}$abek and Henryk Wilczy\\'nski $-$ dzi\\c{e}kuj\\c{e} bardzo! The fruitful collaboration with Ralph Engel and Dieter Heck is gratefully acknowledged. The author is supported by the Alexander von Humboldt Foundation."
    },
    "0402/astro-ph0402650_arXiv.txt": {
        "abstract": "We have conducted a study of optical and \\HI properties of spiral galaxies (size, luminosity, \\Halpha flux distribution, circular velocity, \\HI gas mass) to explore the role of gas stripping as a driver of morphological evolution in clusters. We find a strong correlation between the spiral and S0 fractions within clusters, and the spiral fraction scales tightly with cluster X-ray gas luminosity.  We explore young star formation and identify spirals that are (1) {\\it asymmetric}, with truncated \\Halpha emission and \\HI gas reservoirs on the leading edge of the disk, on a first pass through the dense intracluster medium in the cores of rich clusters; (2) strongly \\HI deficient and {\\it stripped}, with star formation confined to the inner 5 \\hkpc and 3 disk scale lengths; (3) reddened, extremely \\HI deficient and {\\it quenched}, where star formation has been halted across the entire disk.  We propose that these spirals are in successive stages of morphological transformation, between infalling field spirals and cluster S0s, and that the process which acts to remove the \\HI gas reservoir suppresses new star formation on a similarly fast timescale. These data suggest that gas stripping plays a significant role in morphological transformation and rapid truncation of star formation across the disk. ",
        "introduction": "Some of the fundamental questions of galaxy formation are the following: Can we reconcile the observed galaxy populations at high redshifts with those found today in the local universe?  Are the proposed mechanisms of interaction between galaxies and other galaxies or clusters sufficient to explain observations of galaxies within clusters and groups?  Can we similarly explain the evolutionary history of field galaxies without invoking separate formation scenarios?  Studies of distant galaxies have established that there is significant evolution in cluster galaxy populations from redshifts z = 0.5 to the present. Butcher \\& Oemler (1978) first noted the excess of blue galaxies in the cores of rich clusters at redshifts as low as z $\\sim$ 0.2; many studies have since verified and extended this so--called ``Butcher--Oemler effect''.  Parallel studies established that the fraction of spiral and S0s varied inversely within rich clusters, with the S0 fraction dropping by factors of 2-3 from local levels by redshifts z $\\sim$ 0.5 (Couch \\etal 1994; Dressler \\etal 1997; van Dokkum \\etal 1998; Fasano \\etal 2000).  In contrast, the relatively constant numbers and tightly constrained colors of the cluster elliptical population indicated that it was very stable, with star formation concentrated at redshifts z $\\ge$ 3 (Ellis \\etal 1997, though van Dokkum \\etal 1999, 2000 suggest that progenitor bias may elevate the apparent formation redshift). Finally, optical spectra identified numerous poststarburst galaxies within clusters (Dressler \\& Gunn 1983, Couch \\& Sharples 1987), a suppressed star formation rate for spiral types relative to the field (Balogh \\etal 1998), and found evidence of an infalling spiral population out to redshifts z~$\\sim$ 0.4 (Poggianti \\etal 1999; Kodama \\& Bower 2001).  In summary, these intermediate redshift studies support a well-known class of scenarios (\\eg Spitzer \\& Baade 1951; Melnick \\& Sargent 1977 for early discussion) wherein rich clusters are continuously rejuvenated by infalling field spirals (Balogh, Navarro, \\& Morris, 2000); star formation is disrupted and morphology is slowly transformed, some turning into the S0 population of today (Jones, Smail, \\& Couch 2000; Kodama \\& Smail 2001; though see Andreon 1998).  The observational evidence obtained from studying local clusters and groups is more ambiguous.  Dressler (1980) argued strongly against the formation of S0s through transformation of spirals, citing their presence in low density, cool regions, the uniformity of the density-morphology relation in both relaxed and unrelaxed clusters, and the enhanced luminosity of S0 bulges relative to that of the complete spiral population.  Solanes \\& Salvador-Sol\\'e (1992, and references therein) addressed the first two points by proposing that an initial correlation between density and bulge mass occurred during protogalaxy formation, leading to early bulges in pockets of initial overdensities.  This innate bias could explain the continued morphology-density relation in low density regions, but would be erased from high density regions during vigorous stages of cluster formation.  Accreted, gas stripped spirals could gradually build up the core S0 population afterward, thus restoring the morphology-density relation.  Pfenniger (1993) addressed the third point with simulations that suggested that S0 bulges could be supplemented after disk formation by a burst of star formation, fueled by gas funneled to the nucleus via a short-lived bar.  This mechanism would also explain the high metallicity gradients observed in the inner regions of bulges (Fisher, Franx, \\& Illingworth 1996).  The combination of low and intermediate redshift observations indicates that morphological transformation is significant within clusters, but there is much to be determined regarding the details of how it occurs.  Three broad types of physical mechanisms have been proposed to model the transformation of spiral galaxies in dense environments: galaxy-galaxy interactions, tidal forces, and gas stripping.  Galaxy-galaxy interactions are most efficient within groups, where relative velocities are low (Zabludoff \\& Mulchaey 1998; Mulchaey \\& Zabludoff 1998; Ghigna \\etal 1998), and tidal forces (Toomre \\& Toomre 1972; Moore \\etal 1996; Mihos, McGaugh, \\& de Blok 1997) are most effective on small spirals (though see Gnedin 2003a, 2003b).  In spite of these factors, recent large-scale surveys (Lewis \\etal 2002, Nichol \\etal 2002) have found clear evidence of the significance of such moderate physical processes in suppressing star formation, particularly in galaxies located beyond cluster cores, though often at a extreme cost in morphological transformation (\\eg harassment conversion into dwarf spheroidals).  Gas stripping mechanisms, in contrast, are extremely efficient at gas removal, while preserving the large-scale structure of spiral galaxy disks.  They allow for a gradual fading of stellar populations, as \\HII regions diminish in intensity and surface brightness simultaneously decreases {\\it and} becomes more uniform across the disk.  Stripping can take on several forms, including absorption of a hot gas outer envelope (Larson, Tinsley, \\& Caldwell 1980), ram pressure sweeping (Gunn \\& Gott 1972), evaporation via turbulent mixing and heat conduction (Cowie \\& Songaila 1977), and turbulent viscous stripping (Nulson 1982).  Hardware limitations (Steinmetz \\& M\\\"{u}ller 1993; Steinmetz 1996) have precluded incorporating the effects into N-body simulations with gas dynamics at adequate resolution and coverage (time steps and size scales), though recent models of ram pressure stripping of individual Virgo spirals have been quite successful (Abadi, Moore \\& Bower 1999; Vollmer \\etal 1999, 2000, 2001a, 2001b; Vollmer 2003; Vollmer \\& Huchtmeier 2003).  Observational studies of the \\HI gas dynamics (\\cf Giovanelli \\& Haynes 1985; Magri \\etal 1988; Solanes \\etal 2001) have found a strong correlation between \\HI deficiency and clustercentric radius throughout a range of local clusters, strongest in early type spirals (Dressler 1986, though note Koopmann \\& Kenney 1998), and that while \\HI deficiency can extend as far out as 3 \\hMpc most gas stripping occurs well within cluster cores, within galaxies on preferentially radial orbits (Solanes \\etal 2001).  Cayatte \\etal (1990, 1994) obtained aperture synthesis \\HI maps of Virgo cluster galaxies and found signatures of ram pressure and of viscous stripping in separate populations, distinguished by the ratio of \\HI to optical diameters, while Bravo-Alfaro \\etal (1997, 2000, 2001) have used a similar technique to identify first pass spirals falling into and out of the core of Coma and around A262.  An alternate technique is to examine the properties of \\Halpha rotation curves of local cluster spirals.  Early pioneering studies (Rubin, Whitmore, \\& Ford 1988; Whitmore, Forbes, \\& Rubin 1988; Forbes \\& Whitmore 1989) suggested that velocity profiles declined at large optical radii more in cluster spirals than in the field, implying gas stripping or diminished halos, but this result has not been supported by later observations (Distefano \\etal 1990; Amram \\etal 1993; Sperandio \\etal 1995; Vogt 1995).  A recent, large-scale study of 510 rotation curves (Dale \\etal 2001) found instead weak trends in the extent and asymmetry of the \\Halpha flux with clustercentric radius, but the strong selection bias towards late type spirals (Dale \\etal 1999) with strong, extended \\Halpha or \\fnii 6584\\AA\\ emission constrains the application of this result to the general cluster spiral population.  In this paper, we identify and explore a population of infalling spirals in a large sample of spiral galaxies, 296 selected from 18 nearby clusters and 33 isolated field galaxies observed for comparative purposes.  The current program integrates both optical and \\HI observations and is thus sensitive to both gas depletion and star formation suppression.  The targeted galaxies are spread over a wide range of environments, covering three orders of magnitude in cluster X-ray luminosity and containing galaxies located throughout the clusters from rich cores out to sparsely populated outer envelopes.  We have obtained \\Halpha rotation curves to trace the stellar disk kinematics within the potential at high resolution and to explore the strength of current star formation, \\HI line profiles to map the overall distribution and strength of \\HI gas, and \\Ib imaging to study the distribution of light in the underlying, older stellar population.  The sample contains spirals of all types, and is unbiased by the strength of flux from \\HII regions or by \\HI gas detection. This paper is a companion to Vogt, Haynes, Herter \\& Giovanelli (2004a, \\pone), which details the observations and reduction of the data set, and to Vogt, Haynes, Herter \\& Giovanelli (2004b, \\pthree), which explores changes in the fundamental parameters (size, luminosity, and mass) and star formation properties of spiral galaxies as a function of the cluster environment. ",
        "conclusions": "A strong inverse relation is obtained between the fraction of spirals and that of S0s, within the inner 1 \\hMpc or including infalling groups of spirals within up to 2 \\hMpc in local cluster membership calculations.  The elliptical fraction, in contrast, holds fairly constant across the full range of three orders of magnitude in cluster X-ray luminosity. We have explored the strong correlation between \\HI gas stripping and the consequential suppression of young star formation, finding a correlation between the distribution of \\HI flux and of \\HII regions within the galaxy disks.  To this end, we have divided the sample into four groups, on the basis of \\Halpha emission properties.  {\\bf Group I}: {\\it normal} spirals, with no particularly striking properties in the extent and strength of the observed \\Halpha emission flux.  Many, but not all, of these spirals have the expected amount of atomic gas for their optical size and morphological type.  Most of the field spirals fall into this category, and the cluster spirals which do so tend to have similar properties as the field galaxies.  This group also includes a small number of spirals located in the cores of rich clusters with patchy, somewhat decremented \\Halpha emission flux extent.  They appear to be infalling at a face-on orientation into the intracluster medium, and the gas reservoir is being stripped simultaneously across the entire radial extent of the disk.  {\\bf Group II}: {\\it asymmetric} spirals, where the radial extents of \\Halpha emission traced on each side of the disk either differ by more than 5 \\hkpc or form a ratio of less than 1:2.  Galaxies with this degree of asymmetry are not observed beyond 1 \\hMpc in any of the clusters, and are found predominantly in the richest cluster cores.  Half are deficient in \\HI, with an upper limit of 30\\% of the expected \\HI gas, and the distribution of detected \\HI gas correlates with that of the \\Halpha emission.  They are oriented preferentially edge-on to the cores, with truncation of \\Halpha flux {\\it and} \\HI flux along the leading edge, suggesting that ram pressure stripping from a first pass through the intracluster medium plays an important role in generating this effect.  The suppression of star formation along the disk occurs on a timescale similar to that of the \\HI gas stripping, as we find many galaxies where the gas stripping process has begun (low \\HI gas content and asymmetry in the two-horn \\HI line profile) but not yet been completed (a substantial amount of \\HI gas remains) where star formation has already been terminated along the leading edge of the disk.  The {\\it asymmetric} galaxies fall into two groups, with \\BI colors either $<$ 1.5 (blue) or $>$ 2.0 (red).  The blue galaxies have the small $B/T$ fractions of late type spirals, but are even bluer in \\BI, have intrinsically smaller lengths of R$_d$, R$_b$, and \\Halpha extent, and are profoundly \\HI deficient -- suggestive of a central starburst phase stimulated by the gas stripping process.  The red galaxies are larger and brighter and have $B/T$ fractions higher than the mean of the early type (Sa through Sbc) galaxies in the sample.  They have redder \\BI colors, range between normal and deficient in \\HI, and the lengths of R$_b$ and the \\Halpha extent are slightly less {\\it in units of R$_d$}, evidence of mild truncation of star formation in the outer regions of the disk.  Taken together, the evidence supports a less extreme interaction with the cluster, supported perhaps by a low impact parameter or by the greater resistive force of a more massive potential well.  These two groups appear to be on parallel, rather than sequential, tracks in their morphological histories.  {\\bf Group III}: {\\it stripped} spirals, where the extent of \\Halpha emission is less than 5 \\hkpc on both sides of the disk and extends to less than 3 disk scale lengths.  The bulk of these galaxies are strongly deficient in \\HI. They are found in a more relaxed orbital distribution than the {\\it asymmetric} galaxies, out to 2 \\hMpc from the cores of a wider range of clusters.  In combination with their slightly reddened colors, this suggests that they are comprised of less massive systems currently undergoing non-catastrophic stripping (\\ie no starburst phase, and still identifiable as spirals), and of blue {\\it asymmetric} galaxies which have already passed through the cores and are now in a post-starburst phase.  {\\bf Group IV}: {\\it quenched} spirals, for which star formation has been halted across the entire disk and \\Halpha is found only in absorption.  These galaxies range in appearance between possible edge-on S0s, and (primarily early type) spirals with clear spiral arm structure.  They are 1 magnitude fainter in \\Ib than early type \\HI normal field spirals, 0.5 magnitudes redder in \\BI, their disk scale lengths are a factor of 2 smaller, and the extent of the stellar \\Halpha (absorption) flux is quite truncated along the disk (relative to the normal \\Halpha extent).  \\HI observations place an upper limit of one-tenth of the expected \\HI flux on the sample, and 90\\% are undetected.  These galaxies may serve to illustrate the transition stage of a morphological transformation between infalling field spiral and cluster S0s. Their current orbital distribution is far less radially concentrated than that of present day S0s, and on a timescale of a few Gigayears they may slowly blend with that population.  In the richest, hottest clusters we find primarily {\\it asymmetric, stripped,} and {\\it quenched} galaxies, and {\\it normal} edge-on infalling spirals within 1 \\hMpc of the cores.  Less rich clusters contain primarily {\\it normal} spirals, at all radii.  This is consistent with a picture of infalling spirals being significantly altered by ram-pressure stripping in the hot cores, while galaxy-galaxy interactions and tidal forces play a role throughout the entire cluster distribution, and at larger radii.  In summary, we have explored the formation and evolution of spiral galaxies in local clusters through a combination of optical and \\HI properties.  We find a clear relationship between \\HI gas stripping and the consequential suppression of young star formation; both occur quickly within spirals infalling into an intracluster medium of hot gas.  We have traced galaxies through a progression of infall stages, beginning with infalling field spirals and transforming via gas stripping and passive fading into cluster proto-S0s (not with a bang but a whimper)."
    },
    "0402/astro-ph0402185_arXiv.txt": {
        "abstract": "We study cosmological perturbations in generalized Einstein scenarios and show the equivalence of inflationary observables both in the Jordan frame and the Einstein frame. In particular the consistency relation relating the tensor-to-scalar ratio with the spectral index of tensor perturbations coincides with the one in Einstein gravity, which leads to the same likelihood results in terms of inflationary observables. We apply this formalism to nonminimally coupled chaotic inflationary scenarios with potential $V=c\\phi^p$ and place constraints on the strength of nonminimal couplings using a compilation of latest observational data. In the case of the quadratic potential ($p=2$), the nonminimal coupling is constrained to be $\\xi>-7.0 \\times 10^{-3}$ for negative $\\xi$ from the $1\\sigma$ observational contour bound. Although the quartic potential ($p=4$) is under a strong observational pressure for $\\xi=0$, this property is relaxed by taking into account negative nonminimal couplings. We find that inflationary observables are within the $1\\sigma$ contour bound as long as $\\xi<-1.7 \\times 10^{-3}$. We also show that the $p \\ge 6$ cases are disfavoured even in the presence of nonminimal couplings. ",
        "introduction": "The inflationary paradigm has been the backbone of high-energy cosmology over the past 20 years \\cite{review}. The striking feature of the inflationary cosmology is that it predicts nearly scale-invariant, gaussian, adiabatic density perturbations in its simplest form. This prediction shows an excellent agreement with all existing and accumulated data within observational errors. In particular the recent measurement of the Wilkinson Microwave Anisotropy Probe (WMAP) \\cite{Spergel} provided the high-precision dataset from which inflationary models can be seriously constrained \\cite{Bridle,Peiris,Barger,Kinney,Leach,Tegmark}.  The prediction of inflationary observables exhibits some difference depending on the models of inflation. It is important to pick up this slight difference in order to discriminate between a host of inflationary scenarios. Conventionally inflationary models can be classified in three classes \\cite{Kolb}--``large-field\", ``small-field\", and ``hybrid\" models-- depending upon the shape of the inflaton potential. The large field model is characterized by the potential \\begin{eqnarray} V(\\phi)=c \\phi^p\\,, \\label{potential} \\end{eqnarray} which includes only one free parameter for a fixed value of $p$. Therefore it is not generally difficult to place strong constraints on the potential (\\ref{potential}) compared to small-field and hybrid models that involve additional model parameters. In fact it was recently shown in Refs.\\,\\cite{Peiris,Barger,Kinney,Leach,Tegmark} that the quartic potential ($p=4$) is under a strong observational pressure due to the deviation of a scale-invariant spectrum in addition to a high tensor-to-scalar ratio. This situation does not change much in the context of the Randall-Sundrum II braneworld scenario \\cite{LS,TL}. Note that the quadratic potential ($p=2$) is within observational contour bounds due to the flatness of the potential relative to the quartic case.  When we try to construct models of inflation based on particle physics, we do not need to restrict ourselves to the standard Einstein gravity. For example, low-energy effective string theory gives rise to a coupling between the scalar curvature and the dilaton field, which leads to an inflationary solution in the string frame \\cite{PBB}. The Jordan-Brans-Dicke (JBD)-like theories can be viewed as the low-energy limit of superstring theory if the Brans-Dicke scalar plays the role of the dilaton \\cite{FT85,C85}. In this sense the proposal of the extended inflation scenario \\cite{extended} stimulated a further study of more generic classes of inflation models in non-Einstein theories \\cite{soft,Lin90}, in spite of the fact that the first version of the extended inflation resulted in failure due to the graceful exit problem \\cite{presc}.  From the viewpoint of quantum field theory in curved spacetime, nonminimal couplings naturally arise due to their own nontrivial renormalization group flows. The ultraviolet fixed point of these flows is often divergent, implying that nonminimal couplings can be important in the early universe. In this respect Futamase and Maeda \\cite{FM} studied the effect of nonminimal couplings on the dynamics of chaotic inflation and estimated the strength of the coupling $\\xi$ from the requirement of a sufficient inflation. While $|\\xi|$ is required to be much smaller than unity in the quadratic potential, such a constraint is absent in the quadratic potential for negative $\\xi$. Fakir and Unruh \\cite{FU} showed that the fine-tuning problem of the self-coupling $c$ in Eq.~(\\ref{potential}) is relaxed for large negative values of $\\xi$ by evaluating the amplitude of scalar metric perturbations. A number of authors further investigated the spectra of scalar and tensor perturbations generated in inflation \\cite{Salopek,Makino,Salopek2,Kaiser,KF0,KF,SY,massive} and particle productions in reheating \\cite{preheating}\\footnote{The presence of negative nonminimal couplings also leads to the strong variation of curvature perturbations in the context of multi-field inflation. See Refs.~\\cite{multi} for details.}. See e.g., Refs.~\\cite{Gun} for other interesting aspects of nonminimally coupled scalar fields.  In this work we do not restrict ourselves to the Fakir and Unruh scenario, but will place constraints on the strength of nonminimal coupling for the general potential (\\ref{potential}) by using a compilation of recent observational datasets. We shall provide a general formalism for scalar and tensor perturbations in generalized Einstein theories including dilaton gravity, JBD theory and a nonminimally coupled scalar field. This analysis explicitly shows the equivalence of inflationary observables in both the Jordan frame and the Einstein frame, which results in the fact that a separate likelihood analysis in terms of observational quantities is not needed compared to the Einstein gravity. Making use of the 2-dimensional observational constraints of the scalar spectral index $n_{\\rm S}$ and the tensor-to-scalar ratio $R$, we shall carry out a detailed analysis about the constraint on the nonminimally coupled inflaton field for the potential (\\ref{potential}). This provides us an interesting possibility to place strong constraints on the strength of nonminimal couplings from the recent high-precision observations. In fact we  will show that even small nonminimal couplings with $|\\xi| \\ll 1$ can alter the standard prediction of the Einstein gravity. ",
        "conclusions": "In this work we studied cosmological perturbations in generalized gravity theories based on the action (\\ref{lag}). We showed that curvature perturbations in the Jordan frame coincide with those in the Einstein frame. Since tensor perturbations are also invariant under a conformal transformation, the inflationary observables ($n_{\\rm S}$, $n_{\\rm T}$, $R$ and $A_{\\rm S}$) are the same in both frames. This property indicates that the same likelihood can be employed as for the standard Einstein gravity. This is similar to what happens for the Randall-Sundrum II braneworld scenario in which the degeneracy of the consistency relation does not explicitly give rise to the signature of the braneworld \\cite{Huey}, although the constraints of model parameters in terms of underlying potentials are different \\cite{TL}. Remarkably the consistency relation (\\ref{slow}) holds even for the generalized action (\\ref{lag}) that includes dilaton gravity, JBD theory, and a nonminimally coupled scalar field.  We then apply our general formalism to the nonminimally coupled inflaton field with potential (\\ref{potential}). Our main aim is to place strong constraints on the strength of nonminimal couplings using the latest observational data including WMAP, the 2dF and SDSS galaxy redshift surveys. We focused on the case of the negative nonminimal couplings, since the positive coupling $\\xi$ was already severely constrained from the requirement of a sufficient inflation \\cite{FM} ($\\xi$ is at least smaller than $10^{-3}$ and is even much smaller depending on the initial condition of the inflaton).  For the quadratic potential ($p=2$), inclusion of negative nonminimal couplings leads to the decrease of the spectral index $n_{\\rm S}$ and the tensor-to-scalar ratio $R$ (see Figs.\\,\\ref{ns} and \\ref{R}). While the minimally coupled case ($\\xi=0$) is within the $1\\sigma$ contour bound, the theoretical points of larger $|\\xi|$ tend to be away from the observational bounds due to the departure from the scale-invariance of the spectral index (see Fig.\\,\\ref{nsR}). We found the constraints $\\xi>-7.0 \\times 10^{-3}$ at the $1\\sigma$ level and $\\xi>-1.1 \\times 10^{-2}$ at the $2\\sigma$ level.  The quartic potential ($p=4$) suffers from a strong observational pressure for $\\xi=0$, since the $\\xi=0$ case is outside of the $2\\sigma$ contour bound. However this situation is easily improved by taking into account negative nonminimal couplings, as seen in Fig.\\,\\ref{nsR}. The strength of the coupling is constrained to be $\\xi<-1.7 \\times 10^{-3}$ at the $1\\sigma$ level and $\\xi<-3.0 \\times 10^{-4}$ at the $2\\sigma$ level. Note that the Fakir and Unruh scenario with large negative $\\xi$ ($|\\xi| \\gg 1$) is deep inside the $1\\sigma$ bound, thus preferred observationally. We also found that the $p \\ge 6$ cases are outside of the $2\\sigma$ bound even in the presence of negative nonminimal couplings (see Fig.\\,\\ref{nsR}).  While we mainly concentrated on slow-roll inflation, the formula (\\ref{tilns}), (\\ref{nTein}) and (\\ref{ratioep}) can be used in more general theories if the terms $\\dot{\\epsilon}$, $\\dot{\\delta}$ and $\\dot{\\beta}$ vanish. Actually this happens for the dilaton gravity ($F=e^{-\\phi}$ and $\\omega=-e^{-\\phi}$) with an exponential potential. Let us consider a negative exponential potential, $\\hV=-V_0 \\exp(-\\sqrt{2/\\alpha}\\,\\hat{\\phi})$, which appears in the Ekpyrotic scenario \\cite{ekpyr}. Note that this potential is the one in the Einstein frame and the dilaton $\\phi$ is related with the separation of two parallel branes $\\hat{\\phi}$ through the relation $\\phi=-\\sqrt{2}\\hat{\\phi}$ \\cite{TBF}. In this case the background evolution is characterized by \\begin{eqnarray} \\hat{H}=\\frac{\\alpha}{\\hat{t}}\\,,~~~~ \\frac{{\\rm d}\\hat{\\phi}}{{\\rm d}\\hat{t}} =-\\frac{\\sqrt{2\\alpha}}{\\hat{t}}\\,, \\label{HEEin} \\end{eqnarray} in the Einstein frame and \\begin{eqnarray} H=-\\frac{\\sqrt{\\alpha}}{t}\\,,~~~~ \\phi=-\\frac{2\\sqrt{\\alpha}}{1-\\sqrt{\\alpha}} {\\rm ln}\\,\\left[-(1-\\sqrt{\\alpha})t\\right]\\,, \\label{Hstring} \\end{eqnarray} in the string frame \\cite{TBF}. Then we have $\\epsilon=1/\\sqrt{\\alpha}$, $\\hat{\\epsilon}=-1/\\alpha$, $\\delta_{\\rm S}=\\delta_{\\rm T}=\\beta=-1/(1-\\sqrt{\\alpha})$ and $\\hat{\\delta}_{\\rm S}=\\hat{\\delta}_{\\rm T}=0$, which are all constant. Therefore (\\ref{tilns}), (\\ref{nTein}) and (\\ref{ratioep}) can be {\\em exactly} employed in spite of the fact that the background evolution is not slow-roll. The spectral indices $n_{\\rm S}$, $n_{\\rm T}$ and the tensor-to-scalar ratio $R$ are invariant under a conformal transformation, and simply given by \\begin{eqnarray} \\hat{n}_{\\rm S}=n_{\\rm S}=1+\\frac{2}{1-\\alpha}\\,,~~~ \\hat{n}_{\\rm T}=n_{\\rm T}=\\frac{2}{1-\\alpha}\\,,~~~ \\hat{R}=R=\\frac{16}{\\alpha}\\,, \\label{nsntR} \\end{eqnarray} which are highly blue-tilted spectra for $0<\\alpha \\ll 1$ \\cite{bounceper}. Note that these are the spectra generated during the contracting phase and may be affected by the physics around the bounce \\cite{CDC}. We simply presented this example in order to show the validity of the formula (\\ref{tilns}), (\\ref{nTein}) and (\\ref{ratioep}) rather than working on the detailed evolution of perturbations in the bouncing cosmology. We note that the formula (\\ref{tilns}), (\\ref{nTein}) and (\\ref{ratioep}) are automatically valid in slow-roll inflation, since the variation of the terms $\\epsilon$, $\\delta$ and $\\beta$ are negligibly small.  There exist other generalized gravity theories where the function $F$ in the action (\\ref{lag}) depends upon not only $\\phi$ but also the Ricci scalar $R$. The $R^2$ inflationary scenario proposed by Starobinsky \\cite{R2} belongs to this class, in which the spectrum of density perturbations was derived in Refs.\\,\\cite{Sta,Kofman,Hwang01}. In particular Hwang and Noh \\cite{Hwang01} showed that the spectra of both scalar and tensor perturbations are invariant under a conformal transformation as in our action (\\ref{lag}). These facts imply that the degeneracy of the consistency relation persists in a wide variety of gravity theories including the higher-curvature gravity theory and the Randall-Sundrum II braneworld."
    },
    "0402/astro-ph0402466_arXiv.txt": {
        "abstract": "We estimate cosmological parameters using data obtained by the Very Small Array (VSA) in its extended configuration, in conjunction with a variety of other CMB data and external priors. Within the flat $\\Lambda$CDM model, we find that the inclusion of high resolution data from the VSA modifies the limits on the cosmological parameters as compared to those suggested by WMAP alone, while still remaining compatible with their estimates. We find that $\\Omega_{\\rm b}h^2=0.0234^{+0.0012}_{-0.0014}$, $\\Omega_{\\rm dm}h^2=0.111^{+0.014}_{-0.016}$, $h=0.73^{+0.09}_{-0.05}$, $n_{\\rm S}=0.97^{+0.06}_{-0.03}$, $10^{10}A_{\\rm S}=23^{+7}_{-3}$ and $\\tau=0.14^{+0.14}_{-0.07}$ for WMAP and VSA when no external prior is included.On extending the model to include a running spectral index of density fluctuations, we find that the inclusion of VSA data leads to a negative running at a level of more than $95\\%$ confidence ($n_{\\rm run}=-0.069\\pm 0.032$), something which is not significantly changed by the inclusion of a stringent prior on the Hubble constant. Inclusion of prior information from the 2dF galaxy redshift survey reduces the significance of the result by constraining the value of $\\Omega_{\\rm m}$. We discuss the veracity of this result in the context of various systematic effects and also a broken spectral index model. We also constrain the fraction of neutrinos and find that $f_{\\nu}< 0.087$ at $95\\%$ confidence which corresponds to $m_\\nu<0.32{\\rm eV}$ when all neutrino masses are the equal. Finally, we consider the global best fit within a general cosmological model with 12 parameters and find consistency with other analyses available in the literature. The evidence for $n_{\\rm run}<0$ is only marginal within this model. ",
        "introduction": "Recent measurements of the cosmic microwave background (CMB) anisotropies have allowed the determination of a large number of cosmological parameters with unprecedented accuracy. The Wilkinson Microwave Anisotropy Probe (WMAP) and pre-WMAP data sets can be fitted by a six-parameter $\\Lambda$CDM model (see, for example, Bennett et al. 2003, Slosar et al. 2003). In order to break the degeneracies inherent in the CMB power spectrum (Efstathiou \\& Bond 1999), various authors have augmented measurements of the CMB with observations of large-scale structure (LSS), for example, the 2dF galaxy power spectrum (Percival et al. 2001, Percival et al. 2002), Lyman-$\\alpha$ forest (Mandelbaum et al. 2003; Seljak et al. 2003), SDSS three-dimensional power spectrum (Tegmark et al. 2003), measurements of cosmic shear (Hoekstra et al. 2002) and the galaxy cluster luminosity function (Allen et al. 2003a), and/or information on the expansion rate of the Universe from measurements of the Hubble constant (Freedman et al. 2001) and high redshift supernovae (Perlmutter et al. 1999; Reiss et al. 1998).  High-resolution ($\\ell\\ge 700$) observations of CMB anisotropies provided by previously released data obtained by the Very Small Array (VSA, Grainge et al. 2003), the Arcminute Bolometer Array (ACBAR, Kuo et al. 2004) and the Cosmic Background Imager (CBI, Pearson et al. 2003), can also be important in reducing the impact of degeneracies and provide information on the parameters relating to the power spectrum of initial density fluctuations over a much wider range of scales. In particular, the WMAP team made use of these data in their analyses in order to improve the significance of their results (Spergel et al. 2003).  In this paper, we study the cosmological implications of the new CMB power spectrum measured by the VSA which has a good signal-to-noise ratio out to a multipole of $\\ell=1500$ (Dickinson et al. 2004). These observations cover 33 fields, as opposed to 9 in Grainge et al. (2003), representing an improvement of $\\sim 2$ in signal-to-noise over the previous data. By virtue of the accurately measured temperature of Jupiter by WMAP, the absolute calibration uncertainty for these data is reduced to $3\\%$ on the power spectrum; something which will be significant in our subsequent discussion. The power spectrum is measured between $\\ell=300$ and $\\ell=1500$ with a resolution in $\\ell$-space of $\\Delta\\ell\\approx60$. Previous measurements of the power spectrum between $\\ell=130$ and $\\ell=900$ using the VSA compact configuration can be found in Scott et al. (2003).  We will first consider the standard six-parameter flat $\\Lambda$CDM model, and then include extra parameters broadly in keeping with the approach taken in the papers published by the WMAP team (Spergel et al. 2003, Verde et al. 2003, Peiris et al. 2003). Our main focus will be on the initial spectrum of fluctuations, quantified by the running of the spectral index, which appears to be particularly sensitive to high-resolution data such as ours. In the case where we do not impose external priors on the CMB data (WMAP+VSA), we find that there is significant evidence ($>2\\sigma$) for negative running; something which is not implied by the WMAP data alone. The significance of this result is sensitive to the inclusion of external priors, the relative calibration of WMAP and VSA, and possible source/cluster contamination of the measured power spectrum, illustrating issues which are of great relevance in the era of precision cosmology. The result, if true, would be a significant challenge to models of slow-roll inflation, and so we also consider a broken spectral index model. As a final point, we consider a 12-parameter model fit to WMAP, WMAP+VSA and all available CMB data beyond $\\ell>1000$, illustrating the effects of external priors on the estimated parameters. Our results within this model are compatible with previous determinations, both by the WMAP team and others. ",
        "conclusions": "\\label{sec:disc}  We have used recent data from the Very Small Array, together with other CMB datasets and external priors, to set constraints on cosmological parameters. We have considered both flat and non-flat $\\Lambda$CDM models and the results are consistent.  Within the flat $\\Lambda$CDM model, we find that the inclusion of VSA data suggests that the initial fluctuation spectrum that is not described by a single power-law. As we have pointed out already, the value of $n_{\\rm run}$ preferred by the data is incompatible with the basic premises of slow-roll inflation. Moreover, the negative running, which reduces the amount of power on small scales and hence the amount of structure at early times, leads to predictions for the epoch of reionization at odds with the best fit to the CMB data. This comes almost directly from the temperature-polarization cross-correlation power spectrum observed by WMAP (Kogut et al. 2003). Given the implications of this result it is important to consider the possible systematic effects that might weaken it.  The absolute calibration uncertainty of the VSA power spectrum is an important contributory factor to this result. The 3\\% uncertainty quoted in Dickinson et al. (2004) relies heavily on the the measurement of the temperature of Jupiter, $T_{\\rm jup}$, by WMAP and this requires an overall factor of 0.92 modification in the power spectrum estimates from the previous VSA results (Grainge et al 2003; Scott et al 2002) which were reliant on on earlier measurements of $T_{\\rm jup}$ given by Mason et al. (1999). It was pointed out in Dickinson et al. that, in fact, using an absolute calibration based on this measurement of $T_{\\rm jup}$ gives the most consistent normalization of the power spectrum when compared to that of WMAP.  We have investigated the effects of possible uncertainties in the calibration in two ways. First, we consider the possibility of using the Mason et al (1999) central value for $T_{\\rm jup}$ while maintaining an overall uncertainty of $3\\%$. As an alternative we just increase the overall uncertainty in the calibration to $10\\%$ while keeping the central value for $T_{\\rm jup}$ from WMAP. The derived limits on $n_{\\rm S}$ and $n_{\\rm run}$ are presented in Table~\\ref{tab:calib} for these two possibilities using the HST and 2dF priors. We see in each case that the preference for $n_{\\rm run}<0$ is weakened to below $2\\sigma$ compared to the calibration based on WMAP's measurement of $T_{\\rm jup}$.  It is clear that refinement of the absolute calibration of the VSA in the light of the WMAP measurements is something requiring further attention.  \\begin{table*} \\caption{Limits on $n_{\\rm S}$ and $n_{\\rm run}$ in the flat $\\Lambda$CDM model with a running spectral index for different absolute calibration schemes. The uncertainty refers to that in the power. See text for discussion.} \\begin{tabular}{lcccc} \\hline $T_{\\rm jup}$ & Uncertainty  & External & $n_{\\rm S}$ & $n_{\\rm run}$ \\\\ \\hline \\\\ Mason et al. & 3\\% & HST & $0.93^{+0.05}_{-0.05}$ & $-0.058^{+0.038}_{-0.038}$ \\\\ \\\\ Mason et al. & 3\\% & 2dF & $0.93^{+0.05}_{-0.05}$ & $-0.028^{+0.047}_{-0.047}$ \\\\ \\\\ \\\\ WMAP & 10\\% & HST & $0.93^{+0.05}_{-0.05}$ & $-0.055^{+0.035}_{-0.035}$ \\\\ \\\\ WMAP & 10\\% & 2dF & $0.95^{+0.06}_{-0.06}$ & $-0.040^{+0.033}_{-0.033}$ \\\\ \\\\ \\hline\\\\ \\label{tab:calib} \\end{tabular} \\end{table*}  Another possible systematic effect is the residual point source correction due to sources below out subtraction limit of $20{\\rm mJy}$. This was computed by normalizing the point source model of Toffolatti et al. (1998) to the observed VSA source counts which can then be extrapolated to lower flux densities. There are clearly some uncertainties in this procedure. It is possible that an imperfect subtraction, either an over-estimate or an under-estimate, could lead to inaccuracies in the derived limits on the cosmological parameters, in particular on $n_{\\rm S}$ and $n_{\\rm run}$. In order to investigate possible effects of such uncertainties we have performed our likelihood analysis with the inclusion of the parameter $A_{\\rm X}$ which was discussed in section~\\ref{sec:cosmomodel}. We note that it is also possible that for Galactic foregrounds might contribute to this. However, it was shown in Dickinson et al. (2004) that the level of foreground contamination of the VSA fields was negligible.  We find that the derived limits on $n_{\\rm S}$ and $n_{\\rm run}$ in this case are less stringent than without including $A_{\\rm X}$ for WMAP+VSA. The marginalized distributions for $n_{\\rm S}$, $n_{\\rm run}$ and $A_{\\rm X}$ are presented using CMB data from WMAP and WMAP+VSA for the external priors from HST and 2dF in Fig.~\\ref{fig:ax} and the derived limits are presented in Table.~\\ref{tab:ax}. In fact the likelihood curves and derived limits for WMAP and WMAP+VSA are almost identical when $A_{\\rm X}$ is included in the analysis and the WMAP limits are very similar to when $A_{\\rm X}$ is constrained to be  zero (see Table~\\ref{tab:nsnrun}). We see that there are essentially no limits on $A_{\\rm X}$ when just considering WMAP and that in the case of WMAP+VSA, $A_{\\rm X}$ is compatible with zero suggesting that to within at least  $\\approx 100(\\mu{\\rm K})^2$ the source subtraction procedure has been successful. For the HST prior $A_{\\rm X}<214\\,(\\mu{\\rm K})^2$ at 95\\% confidence and for the 2dF prior $A_{\\rm X}<155\\,(\\mu{\\rm K})^2$.  \\begin{table*} \\caption{Limits on $n_{\\rm S}$, $n_{\\rm run}$ and $A_{\\rm X}$ in the flat $\\Lambda$CDM model with a running spectral index when the parameter $A_{\\rm X}$ is included. The final column is the 95\\% confidence upper limit on $A_{\\rm X}$. The units of $A_{\\rm X}$ are in $(\\mu K)^2$.} \\begin{tabular}{lccccc} \\hline CMB & External & $n_{\\rm S}$ & $n_{\\rm run}$ & $A_{\\rm X}$ & $A_{\\rm X}\\,(2\\sigma)$ \\\\ \\hline \\\\ WMAP &  HST & $0.95^{+0.06}_{-0.06}$ & $-0.059^{+0.039}_{-0.039}$ & Unconstrained & \\\\ \\\\ WMAP+VSA &HST & $0.93^{+0.06}_{-0.06}$ & $-0.061^{+0.038}_{-0.038}$ & $-46\\pm 132$ & $<214$  \\\\ \\\\ \\\\ WMAP &  2dF & $0.96^{+0.06}_{-0.06}$ & $-0.036^{+0.036}_{-0.036}$ & Unconstrained & \\\\ \\\\ WMAP+VSA & 2dF & $0.94^{+0.06}_{-0.06}$ & $-0.043^{+0.035}_{-0.035}$ & $-86\\pm 123$ & $<155$  \\\\ \\\\ \\hline \\\\ \\label{tab:ax} \\end{tabular} \\end{table*}  \\begin{figure} \\includegraphics[width=8cm,height=12cm,angle=0]{ax.ps} \\caption{Marginalized distributions for $n_{\\rm S}$, $n_{\\rm run}$ and $A_{\\rm X}$ when $A_{\\rm X}$ is allowed to vary. The line-styles are as in Fig.~\\ref{fig:lcdm}. but COBE+VSA is excluded The left-hand column is for the HST prior and the right-hand for the 2dF prior.} \\label{fig:ax} \\end{figure}  Assuming that the subtraction is perfect and that the SZ contribution to the power spectrum is $\\propto \\ell^2$ for $\\ell<4000$ then our results on $A_{\\rm X}$ can be used to derive a limit on $(\\Delta T_\\ell)^2$ in a bandpower, $B_{3000}$,  covering $2000\\le\\ell\\le 4000$ as observed by CBI. Under the assumption that $\\Delta T_{\\ell}\\propto \\ell^2$, we find that \\begin{equation} B_{3000}={28\\over 3}A_{\\rm X}\\,. \\label{b3000} \\end{equation} This leads to a limit of $B_{3000}< 1997\\,(\\mu{\\rm K})^2$ using the HST prior and $B_{3000}< 1446\\,(\\mu{\\rm K})^2$ for the 2dF prior. The value quoted by Mason et al. (2003) is $B_{3000}=(508\\pm 168)\\,(\\mu{\\rm K})^2$ which is more stringent than our limit. However, this value and our limit are derived in very different ways. The measurement of Mason et al. (2003) is a direct limit from high resolution imaging of 3 deep fields. While it is direct, the measurement the global power spectrum could be significantly affected by sample variance; one could have observed a field which in there are more clusters than the global average and hence obtain a biased estimation of the global power spectrum. Our limit is indirect, coming from the power spectrum measured over 82~${\\rm deg}^2$ at lower angular resolution and it requires that the power spectrum be $\\propto\\ell^2$ as well as the cosmological model to be correct. It is likely that this represents a reliable upper bound since the power spectrum will grow less rapidly than $\\ell^2$ for $\\ell>1500$ due to the fact that the clusters responsible for the SZ effect are not point sources. Moreover, it is not as sensitive to the Poisson distributed  number of clusters in an individual field. A more realistic modelling of the SZ effect using accurate power spectra could yield a more stringent upper bound. The two methods provide useful complimentary information and it is possible that a much more stringent constraint on $A_{\\rm X}$ will be possible when the VSA observes with higher resolution in the near future.  One important feature of the power spectrum observed by WMAP is the apparent absence of power at very low $\\ell$. This could be due to some as yet unknown physics, or it could be a manifestation of the interaction between the subtle systematic effects caused by the side-lobes, the Galactic cut, and the power spectrum estimation algorithm used by the WMAP team (Efstathiou 2003). It is worth assessing to what extent our result is dependent on the measured anisotropies with $\\ell<10$. By excluding the multipoles with $\\ell<10$ from our analysis we find that $n_{\\rm S}=1.01\\pm 0.07$ and $n_{\\rm run}=0.007\\pm 0.049$, strongly suggestive of an $n_S\\equiv 1$, scale invariant initial power spectrum for WMAP with the 2dF prior, whereas $n_{\\rm S}=0.97\\pm 0.06$ and $n_{\\rm run}=-0.015\\pm 0.047$ for WMAP+VSA and the same prior. The weakening of the constraint on $n_{\\rm run}$ should not be a surprise since excluding multipoles with $\\ell<10$ cuts out nearly a whole power of ten in $k$ and $n_{\\rm run}$ is the coefficient of a power series in $\\log(k/k_{\\rm c})$. However, we see that the inclusion of the VSA tends to prefer a spectral index lower than just WMAP. It is clear from this that the reason for the preference for a negative running of the power spectrum when multipoles with $\\ell<10$ are included is the tension between the measurements at $\\ell<10$ by WMAP and for $\\ell>1000$ by the VSA.  We should comment briefly on one aspect of our analysis of the running spectral index models which is not ideal: the preferred values of $z_{\\rm re}$ and $\\tau$. Most recent analyses of CMB data include an upper bound one of these parameters. In Spergel et al. (2003) a flat prior of $\\tau<0.3$ was used when in some cases the data had a preference for a high value of $\\tau$ by virtue of the low-$\\ell$ TE correlation power spectrum; our analysis is no different and we believe that this is responsible for the differences between our analysis and that of Spergel et al (2003). All the likelihood curves and derived limits have made the not unreasonable assumption that $z_{\\rm re}<30$. However, in some cases, particularly those for which we have included no prior from 2dF, the preferred values of $z_{\\rm re}$ are close to this limit, uncomfortably close in some cases and one might be concerned that our results are sensitive to this. We find that the models with $n_{\\rm run}$ significantly less than zero tend to have larger values of $z_{\\rm re}$ which explains why our derived limit of $n_{\\rm run}=-0.060\\pm 0.037$ from WMAP is larger than the one quoted in Spergel et al (2003). It also suggests that, by excluding $z_{\\rm re}>30$, we have weakened the constraint on $n_{\\rm run}$ rather than artificially modifying the preferred value away from zero. An epoch of reionization with $z_{\\rm re}\\approx 30$ would seem unlikely in the context of early structures being the source of ionization, but it is clear that the data suggest it.  For the general 12-parameter $\\Lambda$CDM model, we find that our marginalized distributions for $n_{\\rm S}$ and $n_{\\rm run}$ are broadened, as one would expect. Nevertheless, even in this case, the addition of VSA data significantly reduces tails of the distributions for $n_{\\rm S}$ greater than unity and for positive $n_{\\rm run}$, as compared with using WMAP as the only CMB data set. Indeed, these effects are reinforced by the use of the {\\sc AllCMB} data set. The inclusion of additional CMB data beyond WMAP also leads to a noticeable reduction in the preferred value of $\\omega_{\\rm b}$ and a corresponding increase in $\\omega_{\\rm dm}$.  To summarize, we find that there is evidence for $n_{\\rm run}<0$ in a limited class of models, but within the general $\\Lambda$CDM model with 12 parameters the evidence is much weaker. Standard models of inflation are generally incompatible with such large negative values of $n_{\\rm run}$, but the data appears to point in that direction, although not totally conclusively. The inclusion of an external prior from 2dF appears to weaken the result by fixing $\\Omega_{\\rm m}\\approx 0.3$ in conjunction with the CMB data. The measurement of $\\Omega_{\\rm m}h$ using the galaxy power spectrum is responsible for this shift. It is an interesting question as to how reliable this measurement is since a slight shift in the results, a preference for $\\Omega_{\\rm m}h\\approx 0.17$ rather than $\\Omega_{\\rm m}h\\approx 0.21$ would bring their preferred value into line with that suggested by the CMB alone and would uphold the possibility of $n_{\\rm run}<0$. Since none of the galaxy redshift surveys have conclusively observed the turnover in the power spectrum on which this determination of $\\Omega_{\\rm m}h$ is based we assert that there is still room for some doubt. We have also investigated the possible systematic effects which could weaken our result. We believe that the assumptions behind the power spectrum measurements presented in Dickinson et al. (2004) are the best available using the observations which we have made and the other information from the literature we have used. For sure, we have shown that measurements of the CMB power spectrum beyond $\\ell=1000$ can have an impact on the estimation of cosmological parameters and that future measurements in this region by the VSA, the PLANCK satellite and others will enable us in the future to make more definitive statements."
    },
    "0402/astro-ph0402520_arXiv.txt": {
        "abstract": "We present results of near-infrared photometry ($J$, $H$, $K_S$) for a hard X-ray selected sample of active galactic nuclei (AGNs) obtained from optical identification of the sources detected in {\\it ASCA} surveys (total $\\sim$75 deg$^2$) with a flux limit of $(1-3) \\times 10^{-13}$ erg s$^{-1}$ cm$^{-2}$ (2--10 keV). The sample covers the AGNs at 0.1 $\\lesssim z \\lesssim$ 1 with $L_{\\rm 2-10keV} \\sim 10^{42} - 10^{46}$ erg s$^{-1}$ with very high completeness. The near-infrared photometric data of the sample are obtained from the Two Micron All Sky Survey (2MASS) and observations with Kitt Peak National Observatory 2.1 m telescope and the University of Hawaii 88 inch telescope. The fraction of red ($J-K_{S}>2$ mag) AGNs in our sample is $\\sim 2\\pm1\\%$, which is comparable to that for optically- or UV-selected quasi-stellar objects (QSOs, i.e. luminous AGNs). The number of red AGNs found in our sample is consistent with that expected from the surface density of red AGNs found in 2MASS by Cutri et al. (2001). We find that the anomalously-small dust-to-gas ratios in circumnuclear gas, which is seen in some AGNs with Seyfert-class luminosity, also occur in the QSOs (AGNs with luminosity of $L_{\\rm 2-10keV} \\ge 10^{44.5}$ erg s$^{-1}$) in the systematically surveyed sample in hard X-rays. For all the QSOs with an X-ray absorption of $N_{\\rm H} \\ge 10^{22}$ cm$^{-2}$ in our sample, the values of $A_V/N_{\\rm H}$ are smaller than the Galactic value by a factor of $\\sim$5 to $\\sim$100. Since a fraction of this population among the QSOs in our sample is about 30\\%, such fraction of optical/UV-selected type 1 QSOs known to date may show type 2 nature in X-ray. ",
        "introduction": "Many candidates of red, presumably moderately absorbed, quasi-stellar objects (QSOs) --- luminous active galactic nuclei (AGNs) --- have recently been found in QSO surveys in radio, X-ray, and near-infrared (NIR) wavelengths. Webster et al. (1995) revealed the presence of many red QSOs among radio-selected QSOs and suggested that the red color originated in an absorption to its nucleus. Since a significant fraction of the radio-selected QSOs has redder colors than optically-selected QSOs ($B-K=2\\sim3$ mag), it is argued that 80\\% of QSO population could be significantly absorbed and escape from traditional optical/UV selections. Subsequent studies revealed that the origin of the red color is not absorption to its nucleus, but a red synchrotron continuum of a jet component in $\\sim 80$\\% of the red radio-loud QSOs (Francis et al. 2001). However, it is still possible that there is a large population of moderately ($E(B-V) \\lesssim 1$ mag) reddened QSOs in radio-selected QSOs (White et al. 2003).  On the other hand, for radio-quiet QSOs, which are the dominant population of QSOs in the optical and X-ray wavelengths, the existence of red broad-line QSOs is reported  in identification of soft X-ray selected {\\it ROSAT} sources (Kim \\& Elvis 1999). Since the QSOs have harder X-ray spectra than optically-selected QSOs and  have large H$\\alpha$ to H$\\beta$ broad-line flux ratios, it is likely that the extinction of $A_V \\sim 2$ mag towards the nucleus makes the colors of the QSOs redder than optically-selected QSOs. The fraction of the red ($B-R>2$ mag) QSO is not large ($\\sim 1\\%-7\\%$) in the soft X-ray selected sample, but a number of absorbed QSOs could be equal to that of the blue optically-selected QSOs (Kim \\& Elvis 1999), because the soft X-ray selection is biased to less absorbed QSOs and heavily absorbed QSOs are hard to be detected in the soft X-ray band. More recently, in the NIR wavelength, Two Micron All Sky Survey (2MASS) also has revealed the presence of many AGNs with NIR colors redder than $J-K_S=2$ mag at $z\\lesssim0.3$ (2MASS red AGNs). The number density of the red AGNs is claimed to be comparable to that of optically-selected AGNs with $K_S < 14.5$ mag (Cutri et al. 2001, 2002).  The discoveries of red AGN population suggest that we have been missing a significant fraction of AGNs in traditional optical/UV selections of AGNs due to the moderate absorption to the nucleus. However the fraction of the red absorbed AGNs in the entire AGN population is not clear. Radio-selected samples are affected by red QSOs with the red synchrotron component. The Soft X-ray selection is biased to less absorbed QSOs and thus not suitable to disclose the fraction of the absorbed QSOs. The 2MASS-selected red AGNs are limited in the low redshift universe, primary because of the effect of the $K$-correction on intrinsic NIR colors of QSOs. Using a hard X-ray selected sample is one of the best ways to examine it because the hard X-ray is less affected by absorption, although very recently Wilkes et al. (2002) and Risaliti et al. (2003) have shown the presence of X-ray faint QSOs.  Another important issue in terms of absorption to an active nucleus is anomaly of dust-to-gas ratio in a circumnuclear region. In some of AGNs, an extinction to a nucleus  estimated in optical light is significantly smaller than that derived from X-ray spectrum (e.g., Simpson 1998; Akiyama et al. 2000a; Risaliti et al. 2001; Akiyama, Ueda, \\& Ohta 2002b; Willott et al. 2003). Since the former extinction is caused by dust, while the latter by gas phase metal, this suggests that a  dust-to-gas ratio is much smaller than the standard value obtained in the Galaxy. Maiolino et al. (2001a) compiled Seyfert galaxies showing evidence for absorption in the X-ray and/or NIR/optical data from literature, and found that the $A_V/N_{\\rm H}$ values of most of them are 10 to 100 times smaller than the Galactic value. They claim that the most likely interpretation of this is the depletion of small grains due to coagulation (Maiolino, Marconi, \\& Oliva 2001b), though there may be other possibilities such as dust depletion due to sublimation by AGN radiation, or a difference in locations of X-ray and optical absorbing matters (Weingartner \\& Murray 2002). In any cases,  the anomaly in dust-to-gas ratio affects not only on the understanding of physical conditions in a circumnuclear region, but also on the view of the unified scheme of AGNs; there may be a population that shows type 1 nature in optical/UV region but show type 2 nature in X-ray region. The sample by Maiolino et al. (2001a) is rather heterogeneous and limited to low luminosity AGNs ($<10^{44.5}$ erg s$^{-1}$ in a 2--10 keV band). We intend to extend the luminosity range higher (QSO regime) using an unbiased systematic sample obtained through  optical follow-up of hard-X-ray source surveys.  To make a sample of AGNs at a higher redshift with a larger luminosity without biasing to less absorbed objects, we constructed a highly complete hard X-ray selected AGN sample covering a large survey area. Using a 2--10~keV hard X-ray emission, we can detect AGNs with an absorption up to $N_{\\rm H}$ of $10^{22.5}\\sim10^{23}$ cm$^{-2}$ with a negligible bias in the redshift range below $\\sim 1$. In fact, a {\\it Beppo-SAX} hard X-ray survey found a red QSO (Vignali et al. 2000), and the presence of many red optical counterparts found in deep {\\it Chandra} and/or {\\it XMM-Newton} surveys (Hasinger et al. 2001) hints that there are many red AGNs. However, the optical identification of these samples is not complete so it has been difficult to draw conclusions about what fraction of all AGN that is absorbed.  In this paper, we present  NIR colors of hard X-ray (2--10 keV) selected AGN samples obtained through optical identification of hard X-ray surveys with {\\it ASCA}. The optical identification is fairly complete and the sample is suitable to estimate the fraction of red AGNs and the distribution of $A_V/N_{\\rm H}$ values in QSOs at intermediate redshifts. Throughout this paper, we use cosmological parameters of $H_0=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m =0.3$, and $\\Omega_\\Lambda=0.7$. ",
        "conclusions": "\\subsection{$K_S$ magnitudes and Hard X-ray Flux} In Figure \\ref{k_fx}, $K_{S}$ magnitude against 2--10 keV flux is shown for our sample with open circles. Over-plotted filled circles, crosses, and open squares are the same as those in Figure \\ref{L_z}. A symbol size of an open circle represents the 2--10 keV luminosity; large, middle, and small symbols indicate the AGNs with the 2--10 keV luminosity larger than $10^{45}$ erg s$^{-1}$, $10^{44-45}$ erg s$^{-1}$, and smaller than $10^{44}$ erg s$^{-1}$, respectively. Dashed lines represent the constant X-ray to NIR flux ratio of $\\log f_{\\rm 2-10~keV}/f_{K_{S}}$ = $+$3, $+$2, $+$1, 0, $-$1, $-$2, and $-$3, from top to bottom. Most of the {\\it ASCA} AGNs distribute between $\\log f_{\\rm 2-10~keV}/f_{K_{S}}$ = 0 and +1, which is consistent with that of the mean spectrum of radio-quiet QSOs by Elvis et al. (1994) ($\\log f_{\\rm 2-10~keV}/f_{K_{S}} \\sim +0.4$). All but one of the optically-faint AGNs in the sample (shown with filled circles) fall above the line of $\\log f_{\\rm 2-10~keV}/f_{K_{S}}$ = +1. It is also notable that AGNs without broad H$\\beta$ emission lines (crosses), most of which are type 2 Seyferts at relatively lower redshift, tend to have relatively small X-ray to $K_S$ flux ratios. If the AGN components of them have the same X-ray to $K_S$ flux ratios as those of the other AGNs in our sample, the excess in $K_S$ flux at a given X-ray flux is considered to be caused by the contribution of their host galaxy light.  \\subsection{Color Distribution of the Hard X-ray Selected Sample} The $J-K_{S}$ colors of the {\\it ASCA} AGNs are plotted as a function of redshift in Figure \\ref{jk_z}. Over-plotted filled circles, crosses, and open squares are the same as those in Figure \\ref{L_z}. Small dots represent the optically-selected QSOs; counterparts in 2MASS (Second Incremental Data Release) of QSOs in Veron-Cetty \\& Veron catalog (2000) excluding radio-loud QSOs (Barkhouse \\& Hall 2001). Solid, dotted, and dashed lines show color tracks for the mean-QSO spectrum by Elvis et al. (1994), for that with an absorption of $A_V=1$ mag, and for a template spectrum of an Sbc galaxy by Coleman, Wu, \\& Weedman (1980) to represent the color of a host galaxy of an AGN, respectively. No evolution is considered. $B-K_{S}$ color against redshift for the {\\it ASCA} sample is plotted in Figure \\ref{bk_z}. Marks and tracks are the same as those in Figure \\ref{jk_z}.  In Figures \\ref{jk_z} and \\ref{bk_z}, the {\\it ASCA} AGNs distribute between the Sbc-galaxy and the mean-QSO color tracks or around them at $z \\lesssim 0.3$, while at $z \\gtrsim 0.3$, they locate around the mean-QSO color track. At the low redshift, most of the sample are low luminous AGNs (Seyfert class); the total colors are mixture of those of host galaxies and nuclei. The objects at the high redshift are mostly luminous AGNs, resulting in that their colors show those of ordinary QSOs.  Since almost all of the optically-selected QSOs have $J-K_{S}<2.0$ mag, we adopt $J-K_{S}>2.0$ mag as a criterion for red AGNs, which is the same as that adopted by Cutri et al. (2001). In our sample, two out of 104 AGNs show red color. Therefore, the fraction of the red AGNs in the {\\it ASCA} sample is $2\\pm1\\%$. If the error bars of $J-K_{S}$ color for the sample AGNs are considered, the maximum of the fraction is $\\sim 8\\pm3\\%$. Thus among the AGNs with  the 2--10 keV flux larger than $\\sim 10^{-13}$ erg s$^{-1}$ cm$^{-2}$, the fraction of red AGNs is not significantly larger than that in the optically-selected sample ($\\sim 5\\%$) (Barkhouse \\& Hall 2001).  On the other hand, the $B-K_{S}$ color distribution of the {\\it ASCA} AGNs seems to be biased toward red compared with the optically-selected QSOs (dots) particularly at $z\\lesssim0.3$. Most of the {\\it ASCA} AGNs showing red $B-K_{S}$ color are type 2 Seyferts, so the origin of their red color is likely to be host-galaxy light.  \\subsection{Individual Notes on Red AGNs in the {\\it ASCA} Sample} There are only two AGNs with $J-K_{S}$ colors redder than 2 mag in our sample; 1AXG J121854$+$2957 (NO07) with $J-K_{S} = 2.98$ mag and 1AXG J172938$+$5230 (NO10) with $J-K_{S} = 2.03$ mag. There are six AGNs at $z>0.6$ showing similarly red ($J-K_{S} \\sim 2$ mag) colors, which are redder than those of the other AGNs in the redshift range. These 8 red AGNs are marked with over-plotted pluses in Figure \\ref{k_fx}. We discuss these objects and the origin of the red colors in this section. In each panel of Figure \\ref{indiv}, the NIR and optical photometric data points of each red AGN are shown by filled and open squares with the error, respectively. Solid and dashed lines indicate the mean-QSO spectrum (Elvis et al. 1994) with the absorption shown in the figure (Galactic extinction curve by Cardelli, Clayton, \\& Mathis (1989) is used) and the Sbc-galaxy template (Coleman et al. 1980) without absorption, respectively. These model spectra are independently fitted to the NIR data points of each AGN with the least-square method. It should be noted that the optical and NIR data are not coeval so that the variability may affect the spectral energy distribution (SED) for each AGN.  \\subsubsection{{\\it ASCA} AGNs with $J-K_{S}>2.0$ mag} 1AXG J121854$+$2957 (NO07) has the reddest $J-K_{S}$ color (2.98 mag) among the sample. It is an AGN at $z = 0.178$ with an intrinsic (i.e., absorption corrected) X-ray luminosity of $L_{\\rm 2-10~keV} = 1.1 \\times 10^{44}$ erg s$^{-1}$. The $H\\beta$ emission line of this AGN is narrow, but the $H\\alpha$ emission line has a faint wing (Fiore et al. 1999). Thus this object seems to be a Seyfert 1.9. The SED of the AGN from $B$ to $K_{S}$ is shown in Figure \\ref{indiv}. The mean-QSO spectrum with $A_{V} = 4.5$ mag can well reproduce the SED from $J$ to $K_{S}$. The $A_{V}$ is considerably smaller than that expected from the X-ray absorption of this AGN; by using the Galactic dust-to-gas ratio of $A_V/N_{\\rm H} = 5.6 \\times 10^{-22}$ mag cm$^2$ (Predehl \\& Schmitt 1995), $A_{V} = 46_{-4}^{+6}$ mag is estimated from the $N_{\\rm H} = 8.2_{-0.7}^{+1.1} \\times 10^{22}$ cm$^{-2}$ derived from the X-ray spectrum of this AGN obtained with {\\it XMM-Newton} (Loaring, Page, \\& Ramsay 2003). If the mean-QSO spectrum with $A_{V} = 4.5$ mag is assumed, $B$ and $R$ -band data show excesses; a contribution from the galaxy may brighten $B$ and $R$ -band magnitudes. It is consistent with a fact that this AGN appears to be nearly point-like in NIR images but extended in optical images. Although the NIR SED can be also reproduced by a galaxy spectrum with an absorption of $A_V = 8.4$ mag and with a luminosity of about 17 $L^*$, it is unlikely because the host galaxy seems to be too bright. In Figure \\ref{k_fx}, this AGN locates near the line of $\\log f_{\\rm 2-10~keV}/f_{K_{S}}=0$, which is close to a location occupied by a part of the 2MASS red AGNs shown with filled triangles.  1AXG J172938$+$5230 (NO10) is a broad-line AGN at $z = 0.278$ with an intrinsic X-ray luminosity of $L_{\\rm 2-10~keV} = 1.7 \\times 10^{44}$ erg s$^{-1}$. It looks like a point source on the NIR and optical images. Its SED from $J$ to $K_S$ is well reproduced with the mean-QSO spectrum with $A_V = 0.8$ mag, but in this case, $B$ and $R$ -band magnitudes are too bright. The mean-QSO spectrum without absorption can reproduce the optical and NIR SED reasonably well. It is consistent with the fact that the deviation of $J-K_{S}$ color from the mean-QSO track is small as seen in Figure \\ref{jk_z}. In Figure \\ref{k_fx}, this AGN also locates close to the line of $\\log f_{\\rm 2-10~keV}/f_{K_{S}}=0$.  \\subsubsection{{\\it ASCA} AGNs with $J-K_{S} \\sim 2.0$ mag at $z > 0.6$} 1AXG J160118$+$0844 (NO53) is an AGN at $z = 0.606$ with an intrinsic X-ray luminosity of $L_{\\rm 2-10~keV} = 5.9 \\times 10^{44}$ erg s$^{-1}$. Its X-ray luminosity and narrow $H\\beta$ emission line suggest that it is a candidate of a type 2 QSO. It is an optical-faint, radio-loud (figure 15 of Akiyama et al.(2003)) AGN. In the optical and NIR images, it appears to be a point source. The SED of this AGN is roughly reproduced with the mean-QSO spectrum with an absorption of $A_{V} = 1.5$ mag, which is very much smaller than that estimated from the X-ray absorption ($N_{\\rm H} = 3.0_{-1.2}^{+1.6} \\times 10^{22}$ cm$^{-2}$ corresponding to $A_{V} = 17_{-7}^{+9}$ mag). The SED can also be reproduced with a galaxy spectrum with an intrinsic luminosity of about 6 $L^*$.  AX J131831$+$3341 (LSS228) is an AGN at $z = 0.653$ with an  intrinsic X-ray luminosity of $L_{\\rm 2-10~keV} = 7.8 \\times 10^{44}$ erg s$^{-1}$. Although its $H\\beta$ emission line is narrow, it shows a broad emission line of Mg~II$\\lambda$2800. The absorption derived from X-ray spectrum is $N_{\\rm H} = 1.1_{-0.4}^{+0.4} \\times 10^{22}$ cm$^{-2}$. This object is an optical-faint object. The SED of this object is roughly consistent with the mean-QSO spectrum with $A_{V} = 1.1$ mag or the galaxy template with a luminosity of about 5 $L^*$. However, its optical and NIR images show that it has a nuclear component and an extended component, and we need to discuss the SED for each component. Akiyama et al. (2000b) and Akiyama and Ohta (2001) discuss this in detail. Its nucleus is very red; $I - K = 4.29$ mag, which corresponds to the absorption of $A_{V} \\sim 3$ mag for the power-law continuum of an index of $\\alpha = -0.5$. The NIR and optical colors of the extended component are consistent with the post-starburst nature of the host galaxy.  1AXG J210738-0512 (NO17) is a broad-line AGN at $z = 0.841$ with an intrinsic X-ray luminosity of $ L_{\\rm 2-10~keV} = 1.2 \\times 10^{45}$ erg s$^{-1}$. It appears nearly a point source in the NIR image but a slight elongation may be seen in the optical image. This is an optical-faint, radio-loud object. Its NIR SED is consistent with the absorbed mean-QSO spectrum with $A_{V} = 1.3$ mag, which is very much smaller than that inferred from the X-ray absorption ($N_{\\rm H} = 4.3_{-1.7}^{+2.1} \\times 10^{22}$ cm$^{-2}$ corresponding to $A_V = 24_{-10}^{+12}$ mag). The red NIR colors are also accountable with a galaxy with a luminosity of about 4 $L^*$. However we prefer the former interpretation, because it has a large X-ray luminosity, shows  the broad emission line, and has a point-like morphology.  AX J131021+3019 (LSS039) is a broad-line AGN at $z=1.152$ \\footnote{ The identification of this X-ray source was not correct in Akiyama et al. (2000a).  Subsequent $XMM-Newton$ observation revealed that object Z by Akiyama et al. (2000a) is the correct identification (See also Ueda et al. 2003).} with an intrinsic X-ray luminosity of $L_{\\rm 2-10~keV} = 1.8 \\times 10^{45}$ erg s$^{-1}$. The object is point like in the optical image, but there is a hint of the presence of an extended feature in the NIR images. Its NIR SED is reasonably reproduced with the mean-QSO spectrum with $A_{V} = 1.4$ mag or the galaxy spectrum with a luminosity of about 17 $L^*$. Again we prefer the former interpretation with the same reasons as those for NO17.  1AXG J233200+1945 (NO18) is an AGN at $z = 1.416$ with an intrinsic X-ray luminosity of $L_{\\rm 2-10~keV} = 3.5 \\times 10^{45}$ erg s$^{-1}$. There is a hint of a broad Mg~II emission line. This is an optical-faint source. In the optical and NIR images, it appears to be a point source. The SED of the AGN is consistent with the mean-QSO spectrum with an absorption of $A_{V} = 1.0$ mag, which is nearly two orders of magnitude smaller than that estimated from the X-ray absorption ($N_{\\rm H} = 1.1_{-0.5}^{+0.6} \\times 10^{23}$ cm$^{-2}$ corresponding to $A_V = 62_{-28}^{+34}$ mag). Although the galaxy spectrum with a luminosity of about 17 $L^*$ can also reproduce the NIR SED, we favor the former interpretation with the same reasons as those for NO17.  1AXG J150423+1029 (NO04) is a broad-line AGN at $z = 1.839$ with an intrinsic X-ray luminosity of $L_{\\rm 2-10~keV} = 1.2 \\times 10^{46}$ erg s$^{-1}$. It is an optical-faint object. It is a radio-loud, bright quasar. In the optical and NIR images, it appears to be a point source. Its SED is consistent with the mean-QSO spectrum with an absorption of $A_{V} = 1.2$ mag, which is smaller than that estimated from the X-ray absorption ($N_{\\rm H} = 5.4_{-1.7}^{+1.8} \\times 10^{22}$ cm$^{-2}$ corresponding to $A_V = 30_{-10}^{+10}$ mag) by more than one order of magnitude. The SED is also consistent with the galaxy spectrum with a luminosity of about 1000 $L^*$, which seems to be exceptionally bright. Therefore, we prefer the former interpretation.  In summary, the relatively-red colors in NIR band of these six AGNs at $z>0.6$ are presumably originated from nuclear light with absorption of $A_V \\sim 1-2$ mag or galaxy light with a luminosity of $\\sim 4-17 L^*$. For many of these AGNs, we favor the former interpretation because of their point-like appearance, high luminosity in X-ray, and/or existence of a broad emission line. 5 out of 6 these AGNs have hard X-ray to $K_S$ band flux ratios larger than $\\log f_{\\rm 2-10 keV}/f_{K_S} = +1$. It suggests that these AGNs are different population from the 2MASS red AGNs, which have hard X-ray to $K_S$ band flux ratios of $\\log f_{\\rm 2-10 keV}/f_{K_S} \\lesssim 0$.  \\subsection{Dust-to-Gas Ratios in QSOs} In this subsection, we estimate the dust-to-gas ratios of QSOs in the {\\it ASCA} sample to investigate whether the anomalous dust-to-gas ratios found in low-luminous AGNs (Seyferts) are also seen in QSOs or not. In Figure \\ref{avnh_L}, $A_{V}/N_{\\rm H}$ values normalized to the Galactic one are plotted against 2--10 keV luminosity for the X-ray absorbed ($N_{\\rm H} \\ge 10^{22}$ cm$^{-2}$) QSOs ($L_{\\rm 2-10~keV} \\ge 10^{44.5}$ erg s$^{-1}$) among the {\\it ASCA} sample. Open circles and open triangles show the radio-quiet and radio-loud QSOs of the sample, respectively. Over-plotted crosses and filled circles are the same as those in Figure \\ref{L_z}. The $N_{\\rm H}$ is estimated from the X-ray spectrum as described in section 2, and the $A_{V}$ is calculated from the observed $J-K_{S}$ color on the assumption that the intrinsic spectrum of a QSO is the same as that of the mean-QSO spectrum by Elvis et al. (1994). Since both X-ray and $K_S$ luminosities are very high (the absolute $K_S$ magnitudes of them are brighter than $\\sim -26$ mag which is 1.5 mag brighter than $M_{K}^{*}$ for local galaxies (Huang et al. 2003)), the contribution to NIR light from a host galaxy is assumed to be negligible. The uncertainty of $A_{V}/N_{\\rm H}$ shown in the figure is calculated from the errors of $J-K_{S}$ color and of $N_{\\rm H}$ for each QSO as well as the dispersion of distribution in $J-K_{S}$ color of the optically-selected QSOs. When no nominal optical absorption is obtained, an upper possible limit on $A_V/N_{\\rm H}$ is calculated by considering the uncertainty, and is shown with an arrow.  Figure \\ref{avnh_L} shows that all the sample QSOs have $A_{V}/N_{\\rm H}$ values smaller than the Galactic value by an order of magnitude or more\\footnote{It should be noted that about half (4/10) of the sample QSOs are radio-loud. Since radio-loud QSOs tend to have flatter X-ray spectrum than radio-quiet ones, their $N_{\\rm H}$ values derived by assuming an intrinsic power-law spectrum of photon index of 1.9 may be over-estimated. However, even if an intrinsic photon index of 1.7 is assumed, their $N_{\\rm H}$ values are reduced by only 20\\% typically.}. This trend is the same as that seen in less luminous AGNs by Maiolino et al. (2001a), and it is indicated that the anomaly of the dust-to-gas ratio also occur in QSO-class AGNs. Furthermore, all the objects in Figure \\ref{avnh_L} show the presence of broad components at least in one of the emission lines of C~IV$\\lambda$1549, Mg~II$\\lambda$2800, H$\\beta$, and H$\\alpha$, which also indicates the lack of the heavy extinction in the optical/UV region. The results show not only the existence of a population that has a type 1 nature in the optical/UV band but show a type 2 nature in the X-ray band, but also show that it is not a rare population. Since a fraction of this population in our sample is about 30\\% among the sample with the X-ray luminosity larger than $10^{44.5}$ erg s$^{-1}$, such fraction of optical/UV-selected type 1 QSOs known to date may have large amount of dust-less circumnuclear gases.  \\subsection{Comparison with 2MASS red AGNs} Cutri et al. (2001) revealed the presence of many red ($J-K_{S} >2$ mag) AGNs at $z \\lesssim$ 0.3, by making optical follow-up for the red AGNs candidates selected with the 2MASS catalog (2MASS red AGNs). The optical follow-up spectroscopy of a part of them indicates that they are mixture of type 1 and type 2 AGNs; the number of type 2 is about one third of that of type 1. They claim that the number density of red type 1 AGNs is comparable to that of known QSOs selected in optical or UV.  How many these red AGNs are predicted to be detected in our surveys with {\\it ASCA} ? The total survey area of ALSS and AMSSn is $\\sim$ 75 deg$^2$. Since the surface number density of red AGNs is $\\sim$ 0.2 deg$^{-2}$ for $K_S \\lesssim 14$ mag (Cutri et al. 2001), the region of our surveys should contain about 15 red AGNs with $K_S \\lesssim 14$ mag. Figure \\ref{k_fx} (filled triangles) shows $K_S$ magnitude and X-ray flux in the 2--10 keV band for a well-defined subsample of the 2MASS red AGNs observed with {\\it Chandra} (Wilkes et al. 2002). About 20\\% of them have X-ray flux larger than the detection limit of AMSSn\\footnote{Strictly speaking, since the fluxes of the 2MASS red AGNs are simply converted from the Chandra 0.3--8 keV count rate assuming a photon index of 2, the true 2--10 keV fluxes could be significantly larger than the plotted values in Figure \\ref{k_fx} if they are heavily absorbed as inferred from their hard spectra (Wilkes et al. 2002). Assuming a typical absorption of $N_{\\rm H} = 10^{22.5}$ cm$^{-2}$ at a redshift of 0.2, we expect that about 35\\% of them have count rates larger than the detection limit of the AMSSn.}. Therefore, about 3 red AGNs are expected to be detected in our surveys with {\\it ASCA}. This is consistent with our observations that two red AGNs with $K_S<14$ mag (or four, if errors in the $J-K_S$ color is taken into account) are detected at $z<0.3$; one is a narrow-line AGN (NO07), and the other is a broad-line AGN (NO10)."
    },
    "0402/astro-ph0402098_arXiv.txt": {
        "abstract": "The genesis of spheroids is central to our understanding of galaxy formation -- they are relatively simple systems, containing about half the stellar mass of the Universe. A major subset of spheroids, massive elliptical galaxies, are preferentially found in clusters where they exhibit old, coeval stellar populations suggesting that they formed synchronously at early epochs.  Here we report SCUBA submillimeter imaging of a region around a $z=1.8$ X-ray selected QSO. The image reveals a remarkable $\\sim400$~kiloparsec long chain of galaxies each with an obscured star-formation rate sufficiently high to build a massive spheroid in less than 1~Gyr.  The large over-density of these galaxies relative to expectations for a random field implies they probably reside in a structure associated with the QSO. We suggest that this star formation is associated with galaxy mergers or encounters within the filament, such as those predicted by the popular hierarchical model of galaxy formation. Our observations suggest that strong absorption in the X-ray spectra of QSOs at high-redshifts may result from a veil of gas thrown up by a merger or merger-induced activity, rather than an orientation-dependent obscuring torus. It is argued that these systems are the precursors of elliptical galaxies found today in the core regions of all rich galaxy clusters. ",
        "introduction": "Hierarchical structure formation models are predicated on the assumption that structures grow via the gravitational collapse of small inhomogeneities in a Gaussian random field.  Within such a model it is expected that galaxies and larger scale structures will form at local maxima of the density field \\citep{kai84}. Inhomogeneities in such over-dense regions, that will go on to become the most massive rich clusters at the present day, tend to collapse earlier than similar inhomogeneities in less-dense regions \\citep{kau96}. Moreover, although the power spectrum of the fluctuations in these models are scale free, detailed numerical simulations show that they naturally create large-scale filamentary structures with the most massive bound systems, such as rich clusters, forming at intersections and junctions of these filaments \\citep{col99}.  There is expected to be considerable cross-talk between structures forming on different scales in these models: with galaxy halos growing through major mergers within the larger-scale filaments and streaming down these into the forming cluster.  These galaxies are most naturally identified as the progenitors of the homogeneous population of old, luminous elliptical galaxies which reside in rich clusters at the present day. These models thus predict that the forming elliptical galaxies should be distributed in a highly anisotropic fashion around the collapsing protocluster \\citep{wes94}.  To test these predictions requires observations in a waveband sensitive to massive star formation at high redshifts.  Recent progress in submillimeter astronomy has shown that star formation in massive galaxies at high redshifts is not a luminous phenomenon at optical wavelengths because of the obscuring effect of dust \\citep{sma02}. The reprocessed starlight peaks at far-infrared wavelengths in the rest frame of the source and is shifted into the submillimeter waveband at $z > 1$. We also require a method of identifying an over-dense region of the early Universe. One method is to target powerful active galactic nuclei (AGN). Submillimeter images of the $\\sim1$\\ Mpc-scale fields around such AGN (high-redshift radio galaxies -- HzRGs) at $z>3$ show extended dust emission associated with the AGN, and luminous submillimeter companions \\citep{ivi00,ste03}, several of which are confirmed to be at the same redshifts as the radio galaxies \\citep{sma03,smi03}.  While these results provide support for models of biased galaxy formation, they relate only to a very rare class of tracer galaxies, powerful HzRGs at $z > 3$. Therefore, if we are to investigate the formation of structures the size of typical clusters then less extreme, lower mass signpost AGN at lower redshifts must be identified and their environments mapped at submillimeter wavelengths. With these points in mind we have initiated a programme to image the fields around $1.5 < z < 3$ QSOs detected in the submillimeter waveband \\citep{pag01}. Here we report the first submillimeter imaging observations of one of these QSOs, RXJ094144.51+385434.8 at $z=1.819$.  A Hubble constant $H_0=70$\\ km\\,s$^{-1}$\\,Mpc$^{-1}$ and density parameters $\\Omega_{\\Lambda}=0.7$ and $\\Omega_{\\rm m}=0.3$ are assumed throughout this Letter. ",
        "conclusions": "The 450- and 850-$\\mu$m images reveal a remarkable structure of submillimeter sources in the field of the QSO, particularly striking at 450\\,$\\mu$m (Fig. 1). For the high redshift ($z > 3$) objects targeted to date \\citep{ste03}, the slightly different negative K-corrections that act on the 850- and 450-$\\mu$m flux densities have resulted in poor quality 450-$\\mu$m images. This is because, while the 850-$\\mu$m datum is shifted along the steep Rayleigh-Jeans tail of the dust emission, effectively canceling out the cosmological dimming, the 450-$\\mu$m datum is shifted close to the peak of this emission giving a marked fall-off of the flux density with redshift. In this respect, the lower redshift of RXJ094144.51+385434.8 has worked to our advantage, providing us with an image of the submillimeter dust emission with unprecedented resolution ($8.5''$, equivalent to $\\sim$70~kpc at $z=1.8$).  We identify six sources with peak signal-noise greater than 3 in the 450-$\\mu$m image; they trace a `chain' of submillimeter galaxies with an extent of at least 400~kpc (if at $z=1.8$).  Source names, coordinates and physical quantities calculated from the images are presented in Table~1.  At 450~$\\mu$m the sources are well enough separated to allow good estimation of their individual flux densities; this is not the case at 850\\,$\\mu$m where we give combined flux densities for Nos 1,2 (the QSO) and 4,5. Quoted errors do not include the calibration uncertainties which are about 10\\% at 850~$\\mu$m and 15--20\\% at 450~$\\mu$m.  Dust masses are calculated from the 450~$\\mu$m flux densities in the standard manner \\citep{hil83}, adopting a value for the dust mass absorption coefficient $\\kappa_{100\\mu m}=5.5$\\,m$^2$\\,kg$^{-1}$ \\citep{dra84}. Far-infrared luminosities ($L_{\\rm FIR}$) are calculated by scaling the 450~$\\mu$m emission with that of Mrk~231 which has $L_{\\rm FIR}=2.0\\times10^{12}\\ L_{\\odot}$ (calculated from a grey-body fit to the far-infrared--millimeter SED - the fitted dust temperature is 42~K). Using Arp220 as a template reduces the $L_{\\rm FIR}$ estimates by a factor $\\sim1.6$. The star-formation rates are calculated from SFR(M$_{\\odot}$/yr)$=L_{\\rm FIR}/(5.8\\times10^9L_{\\odot}$) \\citep{ken98}.  Each galaxy in this filament is itself a very massive ultraluminous infrared galaxy (ULIRG) producing stars at a rate sufficient to build a massive spheroid in $<1$ Gyr. Let us assess the significance of this structure. The density of 450-$\\mu$m sources in this field is 1.4$\\pm$0.6~arcmin$^{-2}$, disregarding the QSO which was already known to be a submillimeter source.  The corresponding density of sources found in `blank field' surveys with 450-$\\mu$m flux densities in excess of 30 mJy is uncertain, but best estimates lie in the range $0.03-0.14$~arcmin$^{-2}$ \\citep{sma02}. These figures thus imply an order of magnitude over-density of luminous star-forming galaxies in the $3.5$~arcmin$^{-2}$ field around RXJ094144.51+385434.8. We can investigate the statistical significance of this over-density by calculating the combined probability that these sources each exist at a distance, $d$, from the QSO given the expectation from the known surface density of blank field objects. The high end of the range given above, i.e. $n=0.14$~arcmin$^{-2}$, gives the most pessimistic estimate of this probability (given by $P=1-\\exp\\{-\\pi n d^2\\}$ for each source).  We find $P=2.9\\times10^{-7}$ that the structure we observe is a chance superposition of field objects ($P=6.5\\times10^{-5}$ if source 2 is excluded). This result provides very strong evidence that the companion galaxies lie at the same redshift and in the same structure as the QSO. Assuming this to be true, the star formation rate density calculated in a 500 kpc cube centered on this structure is $>1000$~M$_{\\odot}$\\ yr$^{-1}$\\ Mpc$^{-3}$, about 3 to 4 orders of magnitude higher than found in unbiased optical, infrared and submillimeter surveys \\citep{ivi02,cha03}.  This field thus provides a striking demonstration of the filamentary distribution of forming galaxies in a high-density region at high redshift.  Can we find further support for the hierarchical model in our data?  In the high resolution 450-$\\mu$m image, the central QSO and at least one of the spatially distinct submillimeter sources have complex morphologies, indicative of either merging or interacting systems. These characteristics are analogous to those observed in local ULIRGs \\citep{jos90} although these have more compact millimeter/submillimeter emission \\citep{dow98} than observed here at higher redshift. Moreover, these mergers must be occurring between similarly massive (Table~1) and gas rich systems, indicating that this activity is arising from major, rather than minor, mergers.  Images of the RXJ094144.51+385434.8 field in the $R$- and $K$-band are shown in Fig.~2. The submillimeter galaxies RXJ094144SMM3, RXJ094144SMM4/5 and RXJ094144SMM6 all have counterparts identified as extremely red objects (EROs) \\citep{els88} -- defined as having $(R - K) > 5.3$ \\citep{poz00}. These objects are often found as counterparts to submillimeter galaxies discovered in blank field surveys \\citep{sma99,ivi02}. The counterpart of RXJ094144SMM6 is particularly interesting; it has the appearance of an advanced merger and is a blue/red composite source, again similar to many counterparts of blank-field submillimeter galaxies \\citep{ivi02}. Similar evidence for a merger origin of submillimeter galaxies has been seen in {\\em Chandra\\/} observations which show them to be coincident with two or three X-ray sources -- interpreted as obscured AGN buried in merging star-forming galaxies \\citep{sma03,ale04}.  If the submillimeter-detected galaxies are proto-ellipticals then at $z\\sim2$ they should contain growing black holes which shine as AGN \\citep{kau00,pag01}. However, the {\\em ROSAT\\/} data for this field show that the QSO is the only luminous X-ray source amongst them. If the other dusty galaxies contain AGN then they have to be of lower luminosity or more highly obscured or both. Observations with the new generation of X-ray telescopes will allow us to probe the evolutionary state of the black holes in these systems \\citep{sma03}.  These images are the first to be made of a high redshift, X-ray absorbed QSO at submillimeter wavelengths. The importance of these objects relative to their better studied X-ray unabsorbed (or optically selected) counterparts is only just becoming clear. For matched ($1 < z < 3$) samples selected close to the break of the X-ray luminosity function, the latter are about one order of magnitude less luminous at submillimeter wavelengths. The most likely interpretation is that the X-ray absorption in these objects is not linked with orientation dependent obscuration - the `unified scheme' \\citep{ant93} - but rather with the evolutionary state of the galaxy \\citep{pag04}. Our submillimeter images indicate that this is indeed the case for RXJ094144.51+385434.8 -- its submillimeter luminosity appears to originate in large scale star formation associated with a major galaxy merger.  Finally, let us consider the fate of this structure. Deep {\\it HST\\/} images of the fields around QSOs \\citep{mcl01} suggest that their typical environments are galaxy clusters of Abell richness class 0 \\citep{abe58}. At the present day such clusters contain within their virial radius of order 10 galaxies with luminosity greater than or equal to L$^*$ \\citep{chr03}. If the 6 submillimeter luminous galaxies in the protocluster around RXJ094144.51+385434.8 continue to form stars at the rates inferred from their submillimeter luminosities (Table~1) then their transformation into L$^*$ ellipticals will be complete within in a fraction of a Gyr. We can thus conclude that the coeval stellar populations of the luminous cluster galaxies will be in place by redshift 1.7. Their subsequent passive luminosity evolution will naturally result in a population of ellipticals with properties that match the core regions of today's rich galaxy clusters."
    },
    "0402/astro-ph0402051_arXiv.txt": {
        "abstract": "We present a theoretical and analytical curve with reproduce essential features of the frequency distributions vs. diameter, of the 42,000 crater contained in the Barlow Mars Catalog. The model is derived using reasonable simple assumptions that allow us to relate the present craters population with the craters population at each particular epoch. The model takes into consideration the reduction of the number of craters as a function of time caused by their erosion and obliteration, and this provides a simple and natural explanation for the presence of different slopes in the empirical log-log plot of number of craters (N) vs. diameter (D). ",
        "introduction": "The present impact crater size frequency distribution, N is the result, on one hand, of a rate of crater formation, $\\phi$, and, on the other hand, the elimination of craters, as time goes by, due to effects like erosion and obliteration. Therefore if we want to understand the crater formation history we will need to know how these forming and erasing factors combine to create $N$. Thus, in this work the above problem is analyzed, and in section 2 we find that N can be expressed in terms of $\\phi$ and the fractional reduction of craters per unit of time, C. Then, a simple model is discussed that describe the crater size distribution in Mars data, collected by Barlow \\citep{bar88}, where it is assumed that $\\phi$ is independent of time. The above model is realistic, since according to several investigations $\\phi$ has remained nearly constant for the last 3 to 3.5 billion years \\citep{har66,neu01,neu83,ryder90}. The simplest interpretation of this model implies that $\\phi$ and $C$ are given as the following inverse power of the diameter, $D$, of the crater: $\\phi\\,\\propto\\,\\frac{1}{D^{4.3}}$, $C\\,\\propto\\,\\frac{1}{D^{2.5}}$. In section 3 the model is applied to craters data on Earths, and it is concluded that also in our planet $C\\,\\propto\\,\\frac{1}{D^{2.5}}$. This result is interpreted to mean that on Mars and Earth we have $C\\,\\sim\\,\\frac{1}{\\rm Volume}$, or equivalently the crater mean life $\\equiv\\,\\,\\frac{1}{C}\\,\\propto\\,{\\rm Volume}\\,\\propto\\,D^{2}\\,h$, with $h\\,\\propto\\,D^{0.5}$. Investigations of geometric properties of Martian impact craters reflect values of the average height $h(D)$ consistent with the above conclusion. ",
        "conclusions": ""
    },
    "0402/astro-ph0402267_arXiv.txt": {
        "abstract": "We present new results from a series of $\\Lambda$CDM simulations of cluster mass halos resolved with high force and mass resolution. These results are compared with recently published simulations from groups using various codes including PKDGRAV, ART, TPM, GRAPE and GADGET. Careful resolution tests show that with 25 million particles within the high resolution region we can resolve to about 0.3\\% of the virial radius and that convergence in radius is proportional to the mean interparticle separation. The density profiles of 26 high resolution clusters obtained with the different codes and from different initial conditions agree very well. The average logarithmic slope at one percent of the virial radius is $\\gamma = 1.26$ with a scatter of $\\pm 0.17$. Over the entire resolved regions the density profiles are well fitted by a smooth function that asymptotes to a central cusp $\\rho \\propto r^{-\\gamma}$, where we find $\\gamma=1.16\\pm 0.14$ from the mean of the fits to our six highest resolution clusters. ",
        "introduction": "A highly motivated and well defined problem in computational astrophysics is to compute the non-linear structure of dark matter halos. This is especially timely given the abundance of new high resolution data that probe the central structure of galaxies (e.g. \\citealt*{deBlok2001a}; \\citealt*{deBlok2001b}; \\citealt*{McGaugh2001}; \\citealt*{Swaters2003}; \\citealt*{deBlok2002}; \\citealt*{Gentile2004}) and clusters (e.g. \\citealt*{Sand2003}). Furthermore, a standard cosmological paradigm has been defined that gives a well defined framework within which to perform numerical calculations of structure formation (e.g. \\citealt{Spergel2003}). This subject has developed rapidly over the past few years, building upon the pioneering results obtained in the early 1990's by \\citet{Dubinski1991} and \\citet{Warren1992}. More recently, the systematic study of many halos at a low resolution led to the proposal that the profile of an `average' cold dark matter halo in dynamical equilibrium could be fit by an universal two parameter function \\citep*{Navarro1996}, with a slope of that asymptotically approaches $-1$ as $r \\to 0$. At the same time, the study of a few halos at high resolution questioned these results (\\citealt*{Fukushige1997}; \\citealt*{Moore1998}; \\citealt*{Moore1999}; \\citealt*{Jing2000}; \\citealt*{Ghigna2000}). These latter authors claimed that of the order a million particles within the virialised region where necessary to resolve the halo structure to 1\\% and the slopes at that radius could be significantly steeper. Just within the last few months, we have seen several groups publish reasonably large samples of halos simulated with the necessary resolution that we can finally determine the scatter in the  density profiles across a range of mass scales (\\citealt*{Fukushige2003}; \\citealt*{Tasitsiomi2003}; \\citealt*{Wambsganss2003}; \\citealt*{Hayashi2003}; \\citealt*{Navarro2003}; \\citealt*{Reed2003}).  Much of the recent controversy in the literature has been due to limited statistics and the lack of agreement over what is a reliable radius to trust a given simulation with a given set of parameters. Several studies have attempted to address this issue (\\citealt*{Moore1998}; \\citealt*{Knebe2000}; \\citealt*{Klypin2001}; \\citealt*{Power2003}; \\citealt*{Diemand2004}). Integration and force accuracy can be understood using controlled test simulations. However, discreteness is probably the most important and least understood numerical effect that can influence our numerical results which is exacerbated due to the lack of an analytic solution with which to compare simulations. Our particle sampling of the nearly collisionless fluid we attempt to simulate can lead to energy transfer and mass redistribution, particularly in the central regions that we are often most interested in.  Collisional effects in the final object or in the early hierarchy of objects can be reduced by increasing the number of particles $N$ in a simulation \\citep{Diemand2004}. The limitation to the phase space densities that can be resolved due to discreteness in the initial conditions can also be overcome by increasing the resolution \\citep{Binney2003}. As we increase the resolution within a particular non-linear structure, we find that the global properties of the resolved structure is retained, including shape, density profile, substructure mass functions and even the positions of the infalling substructures. This gives us confidence that our $N$-body calculations are not biased by using finite $N$ \\citep{Baertschiger2002}. The fact that increasing the resolution allows us to resolve smaller radii is important since the baryons often probe just the central few percent of a dark matter structure - the latest observations of galaxy and clusters probe the mass distribution within one percent of the virial radius, which until recently was unresolved by numerical simulations. Forthcoming experiments, such as VERITAS \\citep{Weekes2002} and MAGIC \\citep*{Flix2004} will probe the structure of dark matter halos on even smaller scales by attempting to detect gamma-rays from dark matter annihilation within the central hundred parsecs ($\\sim 0.1\\%R_{virial}$) of the Galactic halo (e.g. \\citealt*{Calcaneo2000}). These scales are still below the resolution limit of todays cosmological simulations, the estimates of the dark matter densities in these regions are still based on extrapolations which introduce large uncertainties  A simple estimate of the scaling of $N$ with time shows remarkable progress over and above that predicted by Moore's law. The first computer simulations used of the order $10^2$ particles and force resolution of the order of the half mass radii \\citep{Peebles1970}. Today we can follow up to $10^8$ particles with a resolution of $10^{-3}$ of the final structure. The increase in resolution is significantly faster than predicted by Moore's law since equally impressive gains in performance have been due to advances in software.  We are finally at the stage whereaby the dark matter clustering is understood at a level where the uncertainties are dominated by the influence of the baryonic component. It is therefore a good time to review and compare existing results from different groups together with a set of new simulations that we have carried out that are the state of the art in this subject and represent what is achievable with several months of dedicated supercomputer time. For certain problems, such as predicting the annihilation flux discussed earlier, it would be necessary to significantly increase the resolution. This is not possible with existing resources and new techniques should be explored. We begin by presenting our new simulations in Section 2. Section 3 discusses convergence tests and the asymptotic best fit density profiles. In Section 4 we compare our results with recently published results from four other groups mentioned above. ",
        "conclusions": "\\label{Summary}  We have carried out a series of six very high resolution calculations of the structure of cluster mass objects in a hierarchical universe. The clusters contain up to 25 million particles and have force softening as small as $0.1\\% r_{\\rm vir}$.  A convergence analysis demonstrates that for our Treecode with our integration scheme, the radius beyond which we can trust the density profiles scale according to the mean interparticle separation. In the best case we reach a resolution of about $0.3 \\%  r_{\\rm vir}$.  Neither of the two parameter functions, the NFW and M99 profiles, are very good fits over the whole resolved range in most clusters. One additional free parameter is needed to fit all six clusters: The asymptotically flat profile from \\cite{Navarro2003} and an NFW profile with variable inner slope provide much improved fits. The best fit inner slopes are $\\gamma=1.16\\pm 0.14$. {\\it Below} the resolved radius the two fitting formulas used are very different. Future simulations with much higher resolution will show which one (if either) of the two is still a good approximation on scales of $0.1\\% r_{\\rm vir}$ and smaller.  We compare our results with simulations from other groups who used independent codes and initial conditions. We find a good agreement between the cluster density profiles calculated with different algorithms. From $0.03 - 0.5 r_{\\rm Vc max}$ the scatter in the profiles is nearly constant and equal to about 0.17 in logarithmic slope. At one percent of the virial radius (defined such that the mean density within $r_{\\rm vir}$ is $178 \\Omega_M^{0.45} \\rho_{\\rm crit} = 98.4 \\rho_{\\rm crit}$) the slope of the density profiles is $1.26 \\pm 0.16$."
    },
    "0402/astro-ph0402117_arXiv.txt": {
        "abstract": "A sample of 123 radio sources that exhibit significant variations at 1.4 GHz on a seven year base-line has been created using FIRST VLA B-array data from 1995 and 2002 on a strip at $\\delta=0$ near the south Galactic cap. This sample spans the range of radio flux densities from $\\sim2$ to 1000 mJy. It presents both in size and radio flux density range a unique starting point for variability studies of galaxies and quasars harboring lower luminosity Active Galactic Nuclei (AGN). We find, by comparing our variable and non-variable control samples to the Sloan Digital Sky Survey the following: 1) The quasar fraction of both the variable and non-variable samples declines as a function of declining radio flux density levels; 2) our variable sample contains a consistently higher fraction of quasars than the non-variable control sample, irrespective of radio flux; 3) the variable sources are almost twice as likely to be retrieved from the optical SDSS data than the non-variable ones; 4) based on relative numbers, we estimate that quasars are about five times more likely to harbor a variable radio source than are galaxies; and 5) there does not appear to be any significant optical color offset between the two samples, even though the suggestive trend for sources to be bluer when variable has been detected before and may be real. This leads us to conclude that both radio variability and radio flux density levels, in combination with accurate optical information, are important discriminators in the study of (radio) variability of galaxies. The latter start to dominate the source counts below $\\sim 20$mJy. In any case, variability appears to be an intrinsic property of radio sources, and is not limited to quasars. Radio variability at low flux density levels may offer a unique tool in AGN unification studies. ",
        "introduction": "All is flux, nothing stays still; nothing endures but change\\footnote{Heraclitus (540-480 BC)}. This seems to apply particularly well to extragalactic radio sources. Important advances in our understanding of radio source intrinsic and extrinsic variability mechanisms have been made by studying their amplitudes, characteristic time-scales, and cross-correlations with observations at other wavelengths. Radio variability can be classified in three broad classes (Padrielli et al. 1987), based on their frequency dependence: 1) variability only occurs at low frequencies ($\\la 2$GHz), 2) variability occurs both at low and high ($\\ga 5$GHz) frequencies, but is not correlated in time, and 3) low and high frequency variability, but the variations are correlated (although they need not occur simultaneously). This classification also conveniently separates the underlying variability mechanisms into ``most likely extrinsic'' (low frequency variations only), and ``most likely intrinsic'' (correlated broadband variability).  The accepted model for the extrinsic variability of radio sources is refractive interstellar scintillation (ISS), which predominantly induces variations at low frequencies ($\\la 1$GHz), with low amplitudes ($\\sim 2\\%$), and with typical timescales of a few days \\citep{blandford86, rickett86}. Variations of this type have been detected in multi-frequency studies (e.g., Mitchell et al. 1994; Lazio et al. 2001).  Perhaps the strongest arguments for the external origin of these variations are correlations with Galactic latitude (e.g., Gregorini et al. 1986; Gaensler et al. 2000), and an annual modulation due to earth's orbital motion (e.g., Lainela \\& Valtaoja 1993; Bondi et al. 1994; 1996). Low-frequency variations with much longer timescales ($\\sim 10$ years), and/or larger amplitudes are thought to be source intrinsic, however (e.g., R\\'ys \\& Machalski 1990 or, more recently, Gaensler \\& Hunstead 2000).  Broadband radio variability, in which the variations usually appear first and more strongly at higher frequencies, are thought to be intrinsic to the radio source. A model describing this behavior, in which shocks propagate along the radio jets, has been proposed by \\citet{marscher85}. The model further implies that sources viewed close to the line of sight (i.e., flat-spectrum quasars and BL Lacs) should be more variable and have shorter variability timescales. For objects very close to the line of sight, relativistic beaming will become important, amplifying the variations and shortening the timescales to $\\la 1$ day (e.g., Lister 2001). Indeed, intraday variability has been readily detected in samples of compact, flat-spectrum sources. Some of the variability occurs on such short timescales that the apparent source brightness temperature (e.g., Gopal-Krishna et al. 1984) can be as high as $T\\approx 10^{21}$K (e.g., Quirrenbach et al. 2000), well beyond the inverse Compton limit of $10^{12}$K (Kellermann \\& Pauliny-Toth 1969). Doppler boosting factors of up to a few hundred are needed to reconcile these numbers (the apparent T is proportional to the Doppler factor to the third power).  These Doppler factors are actually at the high end of the observed range, and it has been suggested that these velocities are pattern speeds that do not reflect the Lorentz factors of the jet itself (e.g., Lister et al. 2001). Nonetheless, intraday variability has been uniquely identified with the blazar class of objects (e.g., Quirrenbach et al. 1992; Lister et al. 2001).  \\begin{figure*}[th] \\epsscale{2.0} \\plotone{devries_aph.fig1.ps} \\caption{FIRST Zero-Dec survey area. Note the non-uniform source density close to the survey boundaries (due to the very different axis scales this is not obvious for the RA edges).} \\label{area} \\end{figure*}  Comparatively little is known about (intrinsic) radio variability in non-blazar sources at low flux density levels ($\\sim 10$ mJy). Most of the variability research has focused on either bright radio samples (e.g., Gregorini et al. 1986, $S > 0.4$ Jy ; Lister et al. 2001, $S > 0.4$ Jy; Aller et al. 2003, $S > 1.3$ Jy), or small samples at much lower flux levels (a few 10's of mJy -- e.g., R\\'ys \\& Machalski 1990; Bondi et al. 1996; Riley et al. 1998). Recently, Carilli et al. (2003) used deep multi-epoch VLA\\footnote{The Very Large Array is a facility of the National Radio Astronomy Observatory (NRAO), which is operated by the Associated Universities, Inc., under cooperative agreement with the National Science Foundation.}  observations of the Lockman Hole area to study variability at the 0.1 mJy level, resulting in just 9 variable sources.  In this paper, we will use two epoch observations of the FIRST Zero-Dec strip (Sect.~\\ref{observations}) of about 9000 radio sources to expand significantly the number of variable sources down to flux density levels of $\\sim 1$ mJy. Unlike much brighter radio samples, the typical radio source population at these levels is dominated by star-forming galaxies and steep-spectrum systems, with only a small percentage of flat-spectrum AGN (e.g., Windhorst et al.  1999; Richards et al. 1999; de Vries et al. 2002). With our sample, therefore, we are in an excellent position to investigate the radio variability properties of lower radio luminosity galaxies.  Interstellar scintillation is not expected to contribute significantly to the variability in our sources, since our observations were carried out at 1.4 GHz, near the minimum in expected ISS (e.g., Padrielli et al.  1987; Mitchell et al. 1994). Furthermore, given our variability criterion (see Sect.~\\ref{sample}), only sources that vary by significant amounts, especially toward the lower flux levels, are included. These variation levels exceed the typical ISS variations of a few percent by a large factor.  The paper is outlined as follows. In the next two sections, the radio observations and sample selection are described. In Sect.~\\ref{optid} we discuss the optical properties of the variable sample, and compare these against a non-variable control sample. ",
        "conclusions": "We have created a sample of 123 variable radio sources using two epoch observations of a zero-dec strip toward the south Galactic cap.  This sample spans the range of radio flux densities from $\\sim2$ to 1000 mJy. It presents both in size and radio flux density coverage a unique starting point for variability studies of more normal, less AGN-dominated galaxies, especially toward the lower flux density limits. We compared both our variable and non-variable samples to the Sloan Digital Sky Survey optical data.  We found that the quasar fraction of the sample sharply declines as a function of declining radio flux density levels. This is consistent with earlier findings that the radio source population demographics change as one samples at progressively lower flux density levels: AGN-dominated systems tend to be found at the brighter radio flux density levels ($> 10$ mJy), whereas star-forming and normal galaxies dominate the counts at sub-mJy flux density levels (at 1.4 GHz, e.g.,  Windhorst 2003).  Our variable sample contains a consistently higher fraction of quasars than the non-variable control sample, independent of radio flux density. While this explains part of our almost $2\\times$ higher optical matching rate of the variable sample compared to the non-variable one (quasars are easier to detect at a given brightness limit than galaxies), it does imply that our variable sample contains on average slightly brighter sources (though not significantly so, see Sect.~\\ref{apmid}). Based on relative number statistics, we estimate that quasars are about 5 times more likely to harbor a variable radio source than galaxies. However, at flux density levels $< 20$ mJy, the majority of (radio) variable sources are identified as galaxies. And finally, galaxies and quasars that harbor a variable radio source exhibit, on average, bluer optical colors than hosts of non-variable sources.  All of this underlines the fact that both galaxies and quasars can harbor variable radio sources, albeit at different occurrence rates. Some of this is obviously due to the beamed nature of the (variable) quasars, enhancing the variability both by boosting their brightnesses and shortening the variability timescales\\footnote{Even though blazars represent a very small fraction of AGN, they are much more common in variable samples.}. Nonetheless, especially toward lower radio flux density limits, a statistically significant study of variability is possible provided one starts with large enough samples. Large scale optical surveys like the SDSS provide the crucial radio source optical ``environmental'' information. As more SDSS data become available on the zero-dec strip, we will investigate this further."
    },
    "0402/astro-ph0402392_arXiv.txt": {
        "abstract": "We present calculations of magnetic field evolution by the Hall effect and Ohmic decay in the crust of neutron stars (NSs). In accreting NSs, Ohmic decay is always the dominant effect due to the large resistivity. In isolated NSs with relatively pure crusts, the Hall effect dominates Ohmic decay after a time $t_{\\rm switch}\\simeq 10^4\\ \\mathrm{yr}\\ B_{12}^{-3}$, where $B_{12}$ is the magnetic field strength in units of $10^{12}\\ \\mathrm{G}$. We compute the evolution of an initial field distribution by Ohmic decay, and give approximate analytic formulas for both the surface and interior fields as a function of time. Due to the strong dependence of $t_{\\rm switch}$ on $B_{12}$, early Ohmic decay can alter the currents down to the base of the crust for $B\\sim 10^{11} \\mathrm{G}$, neutron drip for $B\\sim 10^{12}\\ \\mathrm{G}$, and near the top of the crust for $B\\gtrsim 10^{13}\\ \\mathrm{G}$, respectively.  We then discuss magnetic field evolution by the Hall effect. Several examples are given to illustrate how an initial field configuration evolves. Hall wave eigenfunctions are computed, including the effect of the large density change across the crust. We estimate the response of the crust to the magnetic stresses induced by Hall waves, and give a detailed discussion of the boundary conditions at the solid-liquid interface. Finally, we discuss the implications for the Hall cascade proposed by Goldreich \\& Reisenegger. ",
        "introduction": "Our picture of neutron star magnetic field evolution comes from studies of radio pulsars, accreting neutron stars, and most recently, magnetars. The known radio pulsars fall into two general classes: young ($\\lesssim 10^7\\ \\mathrm{yrs}$) pulsars with spin periods $\\sim 1\\ \\mathrm{s}$ and $\\sim 10^{12}\\ \\mathrm{G}$ magnetic fields, and the millisecond radio pulsars, which have inferred magnetic fields as low as $10^8\\ {\\rm G}$. Whereas most radio pulsars are isolated objects, the millisecond pulsars are predominantly in binaries and have undergone a period of accretion. This has led to the suggestion that millisecond pulsars form by accretion, which both spins up the neutron start to short rotation periods, and also acts to reduce the magnetic field (see Bhattacharya 1995 for a review).  This ``recycling'' scenario is consistent with observations of accreting neutron stars. Young accreting neutron stars in high mass X-ray binaries (HMXBs) are observed as X-ray pulsars with inferred $\\sim 10^{12}\\ {\\rm G}$ magnetic fields, confirmed by detection of cyclotron lines in their X-ray spectra (Tr\\\"umper et al.~1978). Old accreting neutron stars in low mass X-ray binaries (LMXBs) are generally not X-ray pulsars, implying no magnetic disruption of the accretion disk in those systems, and magnetic fields $\\lesssim 10^9$--$10^{10}\\ {\\rm G}$ (see Cumming, Zweibel, \\& Bildsten 2001 for a recent discussion). Recently, 5 millisecond accreting X-ray pulsars have been discovered in transient LMXBs, with spin periods between $180$ and $410\\ \\mathrm{Hz}$, and inferred magnetic fields $10^8$---$10^9\\ \\mathrm{G}$, exactly in the range expected for millisecond pulsar progenitors (Chakrabarty et al.~2003).  An additional possibility is that neutron star magnetic fields decay with age. Field evolution in isolated radio pulsars has been studied using their space and velocity distributions, as well as the statistical distribution in the $P$-$\\dot P$ diagram.  This probes field evolution on the timescale of the ages of radio pulsars ($\\lesssim 10^7\\ {\\rm yrs}$). At first, it was argued that exponential field decay on a few million year timescale was required to understand the $P$--$\\dot P$ distribution (Ostriker \\& Gunn 1969; Gunn \\& Ostriker 1970; Narayan \\& Ostriker 1990), but the work of Bhattacharya et al.~(1992) and others in the 90's argued that in fact no field decay occurs over the lifetime of a radio pulsar. However, this issue remains somewhat open: recently, several authors have again raised the suggestion that neutron star magnetic fields evolve on a timescale something like 10 million years (Cordes \\& Chernoff 1998; Tauris \\& Manchester 1998; Tauris \\& Konar 2001; Gonthier et al.~2002).  The discovery of magnetars has complicated this picture, but also provided a new opportunity to study magnetic field evolution on relatively short timescales. Observed as the young $\\lesssim 10^4\\ \\mathrm{yr}$ soft gamma repeaters (SGRs) and anomalous X-ray pulsars (AXPs), these sources are believed to be directly powered by decay of ultrastrong $\\sim 10^{14}$--$10^{15}\\ \\mathrm{G}$ magnetic fields. Long term release of magnetic energy may act to delay the cooling of the neutron star (Thompson \\& Duncan 1996; Heyl \\& Kulkarni 1998; Colpi, Geppert, \\& Page 2000; Arras, Cumming, and Thompson 2004), while rapid release of magnetic stresses that build in the neutron star crust (Duncan \\& Thompson 1992; Thompson \\& Duncan 1995, 1996) may power the frequent gamma-ray bursts observed from SGRs, and now also AXP 1E 1048.1-5937 (Gavril, Kaspi, \\& Woods 2002).  These ideas and observations have motivated a large range of theoretical work on neutron star magnetic field evolution (e.g.~see Bhattacharya \\& Srinivasan 1995 for a review).  In this paper, we discuss the evolution of currents in the neutron star crust, which evolve due to Ohmic decay, and by the Hall effect. The importance of currents in the crust is that they decay much more quickly than currents in the core. They therefore play an important role in models of accretion-induced field evolution: either the currents are assumed to flow only within in the crust, for example produced after the birth of the neutron star by thermomagnetic effects (Blandford, Applegate, \\& Hernquist 1983), or magnetic flux may be expelled from the core into the crust due to vortex-fluxoid interactions (e.g.~Srinivasan et al.~1990; Ruderman, Zhu, \\& Chen 1998). At higher field strengths, Hall effects most likely are responsible for building up stresses in magnetar crusts which then are released during SGR and AXP bursts.  Previous investigations into Ohmic decay have focused mainly on the decay of the surface field, and involved calculating Ohmic decay eigenmodes for the crust (Sang \\& Chanmugam 1987), self-similar solutions near the surface of the star (Urpin, Chanmugam, \\& Sang 1994), or direct time-integrations of the Ohmic diffusion equation for an initial distribution of currents (Urpin \\& Muslimov 1992; Urpin \\& Konenkov 1997; Page, Geppert, \\& Zannias 2000). For accreting stars, the evolution due to Ohmic decay for fields supported by crustal currents has been calculated in great detail (e.g., Geppert \\& Urpin 1994; Urpin, Geppert, \\& Konenkov 1998; Konar \\& Bhattacharya 1997). We revisit Ohmic decay in this paper, calculating both Ohmic decay modes and self-similar solutions, and show that they are related in a simple way. As well as constraining the ``initial condition'' for the Hall effect, this allows a general way to understand the distribution of currents in the crust at a given age, as well as the decline of the surface field.  The evolution by the Hall effect is much less well-understood, being non-linear and therefore much more complex. Although the Hall effect itself is non-dissipative, Goldreich \\& Reisenegger (1992, hereafter GR) proposed that it leads to dissipation through a turbulent ``Hall cascade'', magnetic energy cascading from large to small scales where it dissipates by Ohmic decay in the Hall time associated with the largest lengthscale. Recent studies by Biskamp and collaborators (Biskamp et al.~1999) of Whistler turbulence indeed show a clear cascade of energy over more than an order of magnitude in lengthscale. Other authors have also demonstrated transfer of magnetic energy between different scales (Muslimov 1994; Naito \\& Kojima 1994; Shalybkov \\& Urpin 1997; Urpin \\& Shalybkov 1999; Hollerbach \\& R\\\"udiger 2002). In a different approach, Vainshtein, Chitre, \\& Olinto (2000) included the effect of the density gradient in a simple model. Most recently, Rheinhardt \\& Geppert (2000), Geppert \\& Rheinhardt (2002), and Rheinhardt, Konenkov, \\& Geppert (2003) discuss a ``Hall drift instability'' which can lead to non-local transfer of energy to small scales.  In this paper, we study the implications of realistic crust models for evolution of crustal currents. In \\S 2, we calculate the timescales associated with Ohmic decay and the Hall effect for realistic crust models. We discuss the relative importance of Ohmic decay and the Hall effect at different stages of a neutron star's life. In \\S 3, we consider Ohmic decay in detail. We show that previous studies of Ohmic decay can be understood in a simple way. In \\S 4, we discuss the evolution due to the Hall effect. We start with a simple discussion of the initial evolution of a dipole field which gives a simple physical picture. We then calculate the properties of linear Hall waves in a realistic crust model. We clarify the physical nature of the ``Hall drift instability'' and comment on its relevance for the overall field evolution. We summarize our main results in \\S 5. In the Appendix, we discuss the behavior of Hall waves at the boundaries of the crust. ",
        "conclusions": "  (i) The relative importance of Ohmic decay and the Hall effect is summarized in Figure \\ref{fig:bt} which shows the $\\Omega\\tau=1$ contour in the $B$-$T$ plane. In isolated neutron stars with $B\\lesssim 10^{13}\\ \\mathrm{G}$, the dominance of the Hall effect depends sensitively on the impurity parameter $Q$: recent calculations by Jones (2001) indicate $Q\\gtrsim 1$ for densities greater than neutron drip (see also de Blasio 2000 for densities lower than neutron drip), in which case the Hall effect never dominates the evolution of crustal currents, considerably simplifying the evolution. If $Q\\sim 10^{-3}$ however, as in the original Flowers \\& Ruderman (1977) estimate, Ohmic decay dominates the Hall effect only for a time  $t_{\\rm switch}\\approx 2\\times 10^4\\ {\\rm years}/B_{12}^3$, after which the conductivity is set by impurity scattering, giving $\\Omega\\tau\\approx B_{12}/Q$. In accreting neutron stars with $B\\lesssim 10^{13}\\ \\mathrm{G}$, Ohmic decay always dominates, since the crust is heated ($T\\gtrsim 10^8\\ {\\rm K}$, e.g.~Brown 2000) and may also be impure $Q\\gtrsim 1$ (Schatz et al.~1999).  (ii) Figure \\ref{fig:ohmictime_vs_depth} summarizes the evolution of currents in the crust while Ohmic decay dominates ($t<t_{\\rm switch}$). The electrical conductivity at these early times is set by phonon scattering. This figure relates previous work on either Ohmic decay eigenmodes, or on self-similar solutions, in a simple way. For a particular time $t$, Figure \\ref{fig:ohmictime_vs_depth} gives the density $\\rho_c$ down to which the currents have significantly decayed. For $\\rho>\\rho_c$, no significant evolution of the currents has occurred. For $\\rho<\\rho_c$, the current distribution is self-similar, being dominated by the Ohmic decay mode with turning point at $\\rho=\\rho_c$. Because $t_{\\rm switch}\\propto 1/B^3$, the allowed distribution of currents when the Hall effect begins to dominate is very sensitive to $B$. For $B\\approx 10^{13}\\ {\\rm G}$, $\\rho_c$ lies at or above the top of the crust; for $B\\approx 10^{11}\\ {\\rm G}$, $\\rho_c$ lies at the base of the crust. At late times, when $\\rho_c$ is greater than the initial location of the currents, the current distribution evolves in a self-similar manner and the surface field decays as a power law (eq.[\\ref{eq:fielddecay}]).  (iii) The basic physics of the Hall effect is that the field in the crust is advected by the electron flow. This is summarized in Figure \\ref{fig:shear}, which considers a variant of the usual laboratory demonstration of the Hall effect. Figure \\ref{fig:twist} shows the evolution of an initial poloidal dipole field, in which the toroidal currents initially \"twist\" the field; the resulting poloidal currents then generate a quadrupole poloidal field via a \"pinch\". These considerations led us to a simple time-independent solution: a poloidal dipole field, in which the electrons rotate rigidly, does not evolve due to the Hall effect.  (iv) We discussed the properties of Hall waves in the crust. The simplest example is a background magnetic field which is uniform and vertical (for example, emerging from the core). The vertical wavelength of the wave grows towards the top of the crust (as pointed out by GR), and we showed that the turning point location is $\\rho_{\\rm turn}\\approx 10^{12}\\ {\\rm g\\ cm^{-3}}/n^{4/3}$. Below the turning point, the wave is able to propagate and its properties are well-described by WKB scalings. Above the turning point, the wave evanesces towards the top of the crust.  (v) A horizontal field (e.g.~as may be produced if the currents supporting the field are predominantly in the crust) has similar wave-like solutions at short wavelengths. However, if there is a background current (the background $B$ varies with height), long wavelength perturbations become unstable. The physics of this instability is that the shear in the background electron velocity drives unstable growth, overcoming the internal shear of the wave which would otherwise drive an oscillation. Therefore we agree with the recent work of Geppert and collaborators that there is an instability, and we have clarified its basic physics. However, we argued that the relevance of this instability for the field evolution in the crust, and particularly for the observable surface field, is not clear.  (vi) We have made a first attempt to address the question of whether Hall waves strain the crust beyond its yield point. We find that the maximum strain is at or above the turning point for the mode (eq.~[\\ref{eq:theta_max}]). However, our calculation is limited to the special case of vertically-propagating waves on a vertical background field; we leave further calculations of the response of the crust to non-vertical background fields and wavevectors for future work. In particular, the force balance is important to consider for magnetar strength fields, for which $B^2>4\\pi\\mu$, directly coupling the induction and momentum equations in the crust (see the recent discussion in Arras et al.~2004).  (vii) Finally, we briefly discussed the non-linear evolution. Recent numerical work on Whistler turbulence (Biskamp et al.~1999) has demonstrated GR's conjecture that the Hall effect would create a \"cascade\" of energy to small scales. However, the scalings of the cascade imply strong, anisotropic turbulence rather than weak, isotropic turbulence as assumed by GR. The Hall wave eigenfunctions calculated in \\S 4 (again, for vertically propagating waves and a vertical field) show that crust breaking is probably not important for standard radio pulsar strength fields."
    },
    "0402/astro-ph0402671_arXiv.txt": {
        "abstract": "Radio emission has been detected from all the stages of stellar evolution across the HR Diagram. Its presence reveals both astrophysical phenomena and stellar activity which, otherwise, would not be detectable by other means. The development of large, sensitive interferometers has allowed us to resolve the radio structure of several stellar systems, providing insights into the mass transfer process in close binary systems. I review the main characteristics of the radio emission from several kinds of stars, paying special attention to those cases where such an emission originates in relativistic jets. ",
        "introduction": "If the stellar radio emission was restricted to solar-like activity levels, our knowledge of the Galaxy in radio wavelenghts would be significantly lesser. The radio flux density $S_{\\nu}$ at a frequency $\\nu$ of a star, measured at the earth, can be expressed in terms of the Rayleigh-Jeans approximation to the black body formula \\begin{equation} S_{\\nu} = {2 k \\nu^{2} T_{\\rm B} \\Omega_{\\rm S} \\over c^{2}}, \\end{equation} where $k$ is the Boltzmann's constant, $T_{\\rm B}$ is the brightness temperature and $\\Omega_{\\rm S}$ is the solid angle subtended by the star. The maximum distance at which a star like the Sun would be observable at 5~GHz with a radio telescope, assuming a typical minimum flux density of 1~mJy, is 0.2~pc for the quiet Sun, 0.3~pc for the slowly varying component, and 7.6 pc for the very strong solar radio bursts. If we take into account that the nearest star, Proxima Centauri, is 1.3~pc away, the solar-like emission would be observed only from a few nearby stars. In fact, the first {\\it radio stars} detected were not really stars emitting at radio wavelengths but strong extragalactic sources.  Fortunately, the situation is much better. Since $S_{\\nu} \\propto T_{\\rm B} \\Omega_{\\rm S}$, it is possible to detect radio stars that have either unusually large emitting surfaces, generally associated to stars having massive stellar winds, or enhanced brightness temperatures, associated to stars exhibiting  powerful non-thermal radio emission.  Early efforts to observe radio stars concentrated on the nearby red supergiants. Kellermann \\& Pauliny-Toth (\\cite{kellerman66}) reported a signal of 0.11$\\pm$0.03~Jy while observing Betelgeuse at 2.8~cm, although no signal was detected on the next 11 nights. Seaquist (\\cite{seaquist67}) reported the detection of Betelgeuse, with a flux density of 0.023$\\pm$0.006~Jy at 2.8~cm, and $\\pi$Aurigae, with a flux density of 0.031$\\pm$0.009~Jy. However, none of these stars has been detected at such flux levels since then. The results obtained in these pioneering observations could have been affected by different factors, such as unknown problems with the equipment, the intrinsic variability of most radio stars, and confusion with background sources.  The fact that most of the solar radio emission occurs in a large variety of flares of varying lifetimes inspired the study of radio emission from non-solar type stars. Red dwarf stars were known to exhibit flares at optical wavelengths, and became obvious candidates for radio flare searches. In 1969 Bernard Lovell, at Jodrell Bank, pioneered these observations. The first indisputable observations of radio flare star emission were the detections of a radio flare in YZ~CMi and in UV~Ceti at 408~MHz, obtained using long baseline interferometers (Davis \\etal\\ \\cite{davis78}). The construction of large aperture synthesis arrays in the mid 1970s allowed to perform reliable measurements of radio emission and marked the beginning of the stellar radio astronomy.  The advent of radio interferometry prompted a significant advance in the knowledge of radio stars. Thanks to this technique it was possible to obtain a) accurate positions, which allow to see if the position of a radio source coincides with that of a star, b) higher sensitivity, because all the data collected for long observing times can be added together, and c) a better determination of the variability, since a time-variable event is most certain if independent subsets of an array are detecting the same time-variable event. Since then, the field of stellar radio astronomy has advanced explosively. We are now able to image radio emission from stellar sources with unprecedented resolution and sensitivity. Thanks to these observational advances, the number and classes of stars known to be sources of radio emission have increased dramatically. As we will see, radio emission has been detected from all the stages of stellar evolution, from birth to death, and has revealed astrophysical phenomena and stellar activity that are not detectable by other means.  To illustrate the progress made in the field of stellar radio astrophysics, it is worth noting that two dedicated conferences were conducted in the past two decades. In 1984, a workshop on stellar continuum radio astronomy was held in Boulder (Hjellming \\& Gibson \\cite{hjellming85}). That workshop reviewed the observational and theoretical work on radio stars from the first decade of the interferometer era. In 1995, a workshop was held in Barcelona, providing an overview of the rich and powerful sets of new radio observations accumulated over the past decade and new theoretical ideas and models arising from them (Taylor \\& Paredes \\cite{taylor96}). ",
        "conclusions": "The study of radio emission from the stars has advanced significantly over the last decades, both theoretically and observationally, thanks to the development of large and sensitive interferometers. We are now able to measure the mass loss rates in hot stars, to estimate the magnetic fields through polarization measurements, to map stellar structures at milliarcsecond level, and to study the processes involved in the particle energization. The mechanisms for producing the stellar radio emission are now better understood. Different energy sources have been invoked to explain the stellar radio emission. For RS~CVn systems the energy source is the magnetic field, whereas for X-ray binaries and microquasars it is the mass accretion onto the compact companion. The study of radio stars also contributes to issues not directly related to the astrophysical phenomena that characterize these sources, as the matching between the radio and the optical reference systems.  Some aspects of the interpretation of the radio emission in active close binaries remain still unclear or need additional observational support. Further advances in the understanding of stellar radio emission are limited by sensitivity and resolution. Furthermore, only a few hundred stellar radio sources are currently known. In this sense, the development of the Square Kilometer Array (SKA) and the Expanded Very Large Array (EVLA) in the near future, interferometers with much better sensitivity and spatial resolution as well as larger frequency coverage than current ones, will undoubtedly prompt a tremendous advance in stellar radio astronomy. For instance, the number of stellar radio sources is expected to increase at least by about four orders of magnitude, providing thus a much larger census of properties that will allow us to better address fundamental questions in astronomy. Also, the discovery of new classes of radio stars will likely bring unexpected phenomena to our attention."
    },
    "0402/astro-ph0402209_arXiv.txt": {
        "abstract": "We present the full spectroscopic white dwarf and hot subdwarf sample from the SDSS first data release, DR1.   We find 2551 white dwarf stars of various types, 240 hot subdwarf stars, and an additional 144 objects we have identified as uncertain white dwarf stars. Of the white dwarf stars, 1888 are non-magnetic DA types and 171, non-magnetic DBs. The remaining (492) objects consist of all different types of white dwarf stars: DO, DQ, DC, DH, DZ, hybrid stars like DAB, etc., and those with non-degenerate companions. We fit the DA and DB spectra with a grid of models to determine the \\Teff and \\logg for each object. For all objects, we provide coordinates, proper motions, SDSS photometric magnitudes, and enough information to retrieve the spectrum/image from the SDSS public database.  This catalog nearly doubles the known sample of spectroscopically-identified white dwarf stars.  In the DR1 imaged area of the sky, we increase the known sample of white dwarf stars by a factor of 8.5. We also comment on several particularly interesting objects in this sample. ",
        "introduction": "The Sloan Digital Sky Survey \\citep[SDSS:][]{yor00} is a continuing imaging and spectroscopic survey of some seven to ten thousand square degrees in the north Galactic cap.  Although its main focus is extra-galactic, there are many Galactic spin-off projects resulting from the survey. The SDSS's principal science objectives focus on obtaining redshifts of distant galaxies and quasars by first imaging the sky in 5 passbands, then selecting potential targets for spectroscopic follow-up based on the 5-band photometry. This spectroscopic selection process is referred to as ``targeting'' and many different targeting categories, each with different criteria and priorities, are used to fill all the fibers available on each 640-fiber spectroscopic plate.  Where there are not enough primary targets (ie., galaxies, QSOs, etc.) to fill  a given plate, the lower priority targeting categories (which include various stellar and serendipity categories) get to allocate fibers. SDSS obtained spectra are of high enough quality to allow accurate object and line identifications well beyond those necessary for redshift determinations. We thus end up with high-quality stellar spectra from objects directly targeted as interesting Galactic objects as well as those thought to be, but which ultimately were not, extra-Galactic objects.  Complementary to \\citet{har03} which presented white dwarf stars from  a sample of early SDSS data, this paper reports on the white dwarf stars found in the spectroscopic data contained within the SDSS Data Release 1 \\citep[ DR1: see also \\url{http://www.sdss.org/dr1}]{dr1}.   In a 190 deg$^2$ area of sky, \\citet{har03} found 260 white dwarf stars. In an area of sky 7.15 times larger, we find a factor of 9.85 more white dwarf stars, or an increased density of approximately 38\\% compared to that of \\citet{har03}. We suspect this difference is simply due to random fluctuations in how we target and acquire white dwarf spectra in the SDSS. QSO target selection code changes, for example, can have significant effects on the ultimate yield of SDSS white dwarf spectra.  Since the DR1 spectroscopic coverage on the sky contains the area analyzed in \\citet{har03}, all those stars are included here, although perhaps with a different spectrum than was analyzed in that work. The white dwarf sample presented here is not meant to be a statistically complete or even well-defined sample of white dwarf stars; it merely represents the white dwarf stars that happen to have spectra in DR1.  There are undoubtedly many more white dwarf stars contained within DR1 for which we only have photometry.  We make no attempt to report on those stars here (except see the Appendix for data on previously known white dwarf stars in DR1 for which we do not have spectra).  We also include spectra for hot subdwarf stars, the sdB and sdO stars, since they overlap the hot white dwarf stars in color-color space and need to be identified in order to find the white dwarf stars (besides being interesting in their own right, of course).  For the  broader context of white dwarf stars in the SDSS and particular notes of some unusual objects, see the \\citet{har03} paper and the earlier simulations of \\citet{fan99} who discusses where the white dwarf stars should appear in the SDSS photometric color space.  The SDSS spectra are well-described in \\citet{edr} and \\citet{dr1}.  Briefly, they cover a wavelength range of approximately 3800--9200\\AA\\ with a resolution $\\sim 1800$ and are spectrophotometrically calibrated to within about 10\\% on average. The average S/N of a {\\it g}=20.2 spectrum is $\\sim 4$ per pixel and the redshift accuracy is of order 30 \\kms, as determined for the SDSS galaxy sample.  The \\citet{mcc99} catalog lists 2249 white dwarf stars while the online updates (at {\\url http://www.astronomy.villanova.edu/WDCatalog/index.htm}) now include a total of 3066 white dwarf stars (as of June, 2003), some of which are included in the first SDSS white dwarf catalog paper by \\citet{har03}.  Here, we present 2551 certain white dwarf stars, 240 hot subdwarf stars, and another 144 possible, but uncertain white dwarf and hot subdwarf stars from the 1360 deg$^2$ of DR1.  We find 108 of our white dwarf stars are already present in the McCook \\& Sion (1999) catalog. We present a more complete description of the overlap between the two catalogs in the Appendix.  Another comprehensive spectroscopic survey, the 2dF QSO Redshift Survey \\citep{boy00}, has also produced a sizeable catalog of white dwarf stars \\citep{ven02}.  The 2dF survey went a bit deeper than the SDSS, but covered a smaller solid angle: 740~deg$^2$ with $18.4 \\leq B \\leq 21.0$ \\citep{ven02}.  Their results include 942 spectroscopically-identified DA white dwarf stars, providing another large increase in the number of known white dwarf stars. ",
        "conclusions": "By selecting candidate spectra based mostly on a variety of color and proper motion cuts, we have found 2791 spectroscopically-confirmed white dwarf and subdwarf stars in the first data release of the SDSS.  We currently see no reason why this detection rate should change in future SDSS data releases,  and thus can look forward to an additional 10,000 or so white dwarf stars from the SDSS by the time it is finished.  This and future catalogs will allow us to expand dramatically our sample of particularly unique, interesting, and rate white dwarf stars. We have already pointed out some of the work going on in this regard with the DQ, DH, pulsating, and hot DO white dwarf stars.  We expect to make updated catalogs available as more data are released by the SDSS and welcome contact with fellow white dwarf researchers with regards to SDSS objects.  An ultimate goal of white dwarf research is a vastly improved luminosity function over what is known at present. The vagaries of how SDSS white dwarf spectra are obtained makes building a full luminosity function a difficult task, but one which should be ultimately possible.  However, a luminosity function concentrating on the hot end of the white dwarf cooling sequence suffers from fewer targeting complications and thus should be doable on much shorter timescales than the entire luminosity function will require. We are currently beginning this work in addition to preparing the next public catalog of white dwarf stars from the SDSS.  The data tables presented here are available online at the SDSS DR1 value added catalog site: {\\url http:///www.sdss.org/dr1/products/value\\_added/wdcat/dr1/index.html}."
    },
    "0402/astro-ph0402401.txt": {
        "abstract": "Using photometric redshifts we determine the galaxy population of the clusters of galaxies Cl0016+16 at $z~=~0.55$, Cl1600+41 at $z~=~0.54$, Cl1601+42 at $z~=~0.54$ and MS1008--1224 at $z~=~0.31$. Comparing the clusters, we find no evidence for a universal shape of the total luminosity function (LF) at these redshifts. When dividing the LFs into spectral types, we find that the LF of the early-type galaxies alone can be described by a Gaussian, while the LF of the late-type galaxies is well fitted by a Schechter function, suggesting that the separate LFs for different populations may be universal. The difference in the total LFs can mainly be attributed to the varying relative normalisation of these populations, implying that clusters with an abundant population of late-type galaxies also have steeper faint-end slopes. In MS1008--1224 we detect a faint blue population that dominates over a population with colours consistent with dwarf ellipticals, opposite to clusters at lower redshift.  Compared to low-redshift clusters, we find that a general fading of the late-type population by $\\sim~2$ mag and the early-type population by $\\sim$ 1 mag describes the evolution from $z~=~0.55$ to $z~=~0$ well.  As a consequence of the different early-type and late-type LFs and their dependence on cluster radius, the fraction of blue cluster galaxies, as measured by the Butcher--Oemler effect, differs between the clusters and depends on limiting magnitude and radius.  We find a correlation between the dwarf-to-giant ratio and the surface density, indicating that the high density environment in the cluster cores is hostile to dwarf galaxies. ",
        "introduction": "Ever since the results of Butcher \\& Oemler (1984, hereafter BO84), showing that the fraction of blue galaxies in clusters increases rapidly between $z~=~0$ and $z~\\sim~0.5$, it has been clear that there is strong evolution in cluster galaxy populations with redshift. These results have been confirmed by subsequent surveys (e.g. Rakos \\& Schombert 1995; Ellingson et al. 2001; Kodama \\& Bower 2001; Margoniner et al. 2001). A general picture explaining the over-all features of cluster formation and evolution is given within the cold dark matter (CDM) hierarchical clustering scenario (e.g. Bower 1991; Kauffmann \\& White 1993; Kauffmann 1995; Baugh, Cole \\& Frenk 1996). A large number of observations and N-body simulations now support this scenario. The blueing with redshift of the cluster galaxies is explained by a higher accretion rate of star forming field galaxies at higher redshift (Bower 1991), possibly combined with a general increase in star formation in field galaxies at higher $z$ (Diaferio et al. 2001). A consequence of the CDM model is that high-$z$ clusters assemble during a much shorter time interval than comparably rich clusters today (Kauffmann 1995). This general scenario is also supported by the observed increase of mergers in high-redshift clusters (van Dokkum et al. 1999). Furthermore, N-body simulations (e.g., Dubinski 1998) show that the brightest cluster galaxies (BCG) naturally form via mergers in a hierarchical scenario.  The difference in radial distribution between early-type and late-type galaxies in clusters can be explained by the hierarchical model if infalling field galaxies, which are predominantly blue star-forming galaxies, have their star formation truncated as they fall deeper into the cluster potential (Balogh, Navarro \\& Morris 2000; Diaferio et al. 2001). Mergers and interactions in the high-density cluster environment also affect the star formation, transforming the spectral type of the galaxies.  One of the most important probes of cluster evolution is the luminosity function (hereafter LF), describing the number of galaxies per magnitude bin in a cluster. Binggeli, Sandage \\& Tammann (1988) showed that the total LF for the Virgo cluster rises as a Gaussian at bright magnitudes and then flattens to a \"plateau\" (or even a slight decrease), before it turns steep at faint magnitudes. LFs with similar shapes are found in other nearby clusters, e.g. Coma (Trentham 1998a). The existence of a numerous population of faint dwarf galaxies in clusters, more common than in the field, is verified by a number of observations (e.g., Smith, Driver \\& Phillipps 1997; Wilson et al. 1997; Phillipps et al. 1998; Trentham 1998a; Yagi et al. 2002). There are, however, also clusters that do not show this steep increase at faint magnitudes (Trentham 1998b, 1998c). Most likely, this is a consequence of different evolution depending on e.g., richness and epoch of formation.  Conselice, Gallagher \\& Wyse (2001) discuss different scenarios that could lead to the dwarf population in local clusters. By comparing models with the observed kinematic and spatial properties of a number of Virgo cluster galaxies, they suggest that a transformation of spirals into dwarf ellipticals (dEs) by \"galaxy harassment\" (Moore, Lake \\& Katz 1998) is occurring. Some dEs may, however, have formed outside the cluster in galaxy groups, which were later accreted. These dEs may therefore be as old as the cluster ellipticals. The fading of these cluster galaxies into todays dwarf population is discussed by Wilson et al. (1997) and Smail et al. (1998).  Besides measuring the total LF, it is of high interest to investigate the LFs of different galaxy populations and their evolution with redshift. While previous studies of type-dependent LFs have focused on {\\it morphological types}, e.g., Binggeli et al. (1988) and Andreon (1998), here we divide galaxies into {\\it spectral types} characterised by their colours, as described in Section 3.3. In this paper, when we refer to galaxies as early-types and late-types we mean their spectral type, while morphological types are denoted by their type in the Hubble sequence, e.g., ellipticals and spirals.  A consequence of the hierarchical scenario is that clusters observed at high redshift ($z~\\gsim~0.5$), where the infall of field galaxies is assumed to peak (Bower, Kodama \\& Terlevich 1998), should have a higher fraction of late-type galaxies compared to local clusters. This should lead to a differential evolution between the late-type and early-type cluster LFs with redshift, with the late-type LF shifted towards brighter magnitudes at increasing redshift, as compared to the early-types.  Because it is observationally costly, if at all possible, to obtain redshifts for the numerous faint population, most observational efforts to study galaxy populations in high-$z$ clusters have concentrated on the brighter galaxies ($M_B~\\lsim~-19$). Recently, however, photometric redshifts have been demonstrated to be a powerful tool for studying especially the faint population of cluster galaxies. Even though the accuracy of photometric redshifts can not be compared to spectroscopic redshifts, they can provide a reliable determination of clusters membership, and therefore significantly reduce the necessary amount of field galaxy subtraction. In a previous paper (Dahl\\'{e}n, Fransson \\& N\\\"{a}slund 2002, hereafter DFN02), we demonstrated the use of photometric redshifts in a study of the population of the intermediate rich cluster Cl1601+42 at $z~=0.54$. The photometric redshift selection minimizes the amount of background subtraction needed. Internal properties of the cluster, such as radial distribution and luminosity function, both for the total cluster population, as well as for different populations separately, can therefore be determined. Here we extend our previous study of Cl1601+42 to include two additional rich clusters, Cl0016+16 at $z~=~0.55$ and MS1008--1224 at $z~=~0.31$. We also include observations of a fourth poor cluster, Cl1600+41 at $z~=~0.54$.  Throughout this paper we assume a Hubble constant $H_0~=50$ km s$^{-1}$ Mpc$^{-1}$, and a cosmology with $\\Omega_M~=~0.3$ and $\\Omega_{\\Lambda}~=~0.7$, unless anything else is assumed. Magnitudes are given in the Vega based system. ",
        "conclusions": "\\subsection{Variations in the cluster luminosity function} The different shapes of the LFs in our sample suggest that there is no universal form of the total cluster LF at $z~\\gsim~0.3$ (Fig. \\ref{Figure5}). To quantify this we construct an average LF from the three individual best-fitting LFs. We then calculate the reduced chi-square when fitting each LF to the average LF. We also calculate the reduced chi-square when fitting the individual LFs to a Gaussian+Schechter function. Results are shown in Table 8. The chi-squares are significantly higher for the fits to the average LF, supporting our claim for a non-universal LF. We also find that a single Schechter function gives a poor representation of the cluster LF, which agrees with e.g. Driver et al. (1994) and Wilson et al. (1997). This is especially evident for MS1008--1224, where the best Schechter fit to $M_B~<~-17.7$ yields a faint-end slope that is decreasing, while the data-points clearly show an increase in the number of faint galaxies with $M_B~>~-19$. As in the case for Virgo, a better representation is given by the sum of a Gaussian and Schechter function. The reduced chi-square for for a Gaussian+Schechter function is $\\chi^2/\\nu~=~0.98$ , while a single Schechter function yields $\\chi^2/\\nu~=~5.8$.  From the separate late-type and early-type LFs (right-hand panels of Fig. \\ref{Figure5}), it is clear that the relative abundance of the early-type and the late-type populations determines the over-all shape of the LF. In Cl1601+42 late-type galaxies dominate at $M_B~\\gsim~-20$. This results in a total LF that increases over the whole magnitude range, and has a steep faint-end slope. MS1008--1224 is similar to Cl1601+42 in that it contains a numerous population of faint late-type galaxies, which dominates the total LF at faint magnitudes. The magnitude where this cross-over takes place is, however, slightly fainter in MS1008--1224 ($M_B~\\sim~-19.5$) compared to Cl1601+42 ($M_B~\\sim~-20$), which leads to a total LF where both an intermediate-magnitude Gaussian part, as well as a steep faint-end slope of late-type galaxies are distinguishable. Note also that MS1008--1224 has a population of faint {\\it early-type} galaxies, which adds to this.  In Cl0016+16 there are relatively few faint late-type galaxies, compared to the number of early-type galaxies at $-21~\\lsim~M_B~\\lsim~-19$. This leads to a total LF that has a Gaussian shape at intermediate magnitudes, and only at the faintest bin is there an indication of a rise.  As discussed in DFN02, the differences between the shapes of the cluster LFs can be explained by the fact that clusters in the hierarchical clustering scenario at high redshift have a larger fraction of newly accreted galaxies with ongoing star formation. A general increase in the star-formation rate in field galaxies with redshift (Diaferio et al. 2001), also contributes to an increase of star-forming galaxies in clusters at high $z$.  After accretion, and a possible period of enhanced star formation, the galaxies could have their gaseous envelopes removed by tides or ram pressure stripping, leading to a fading over time-scales of Gyrs, as the remaining gas reservoir is exhausted ('strangulation') (Balogh et al. 2000; Diaferio et al. 2001). Such fading of the faint blue galaxies is discussed by Wilson et al. (1997), who find that a fading of dwarf irregulars by $\\sim~3$ mag can explain the difference between the LF in clusters $z~\\sim~0.2$ and the local Virgo cluster. Further, Conselice et al. (2001) show that spirals accreted at high $z$ can be transformed into the faint dE population seen in local clusters by \"galaxy harassment\" (Moore et al. 1998).  It is of obvious interest to compare the LFs of the different populations of these clusters with those of nearby clusters like Coma and Virgo. We then have the paradoxical situation that studies of nearby cluster are limited to one, two or at most three colours, or alternatively to a morphological separation of different classes. A direct comparison with our spectral classification is therefore difficult. As a first step we therefore compare the {\\it total} LF of our clusters with that of Coma at $z~=~0.02$ (Trentham 1998a) and Virgo $z~=~0.003$ (Trentham \\& Hodgkin 2002).  In Fig. \\ref{Figure6} we plot the rest-frame $B$--band LF for Cl1601+42 (open circles), Cl0016+16 (filled circles), and MS1008--1224 (triangles) together with the LF for the rich Coma cluster (squares) and the poor Virgo cluster (diamonds). For clarity, the LFs have arbitrary off-sets in the $y$-direction. As suggested by our discussion of the individual LFs, we represent the total LF as a sum of a Gaussian and a Schechter function. Because the bright and faint populations have distinctly different properties, we discuss them below separately. \\begin{table} \\caption{Reduced chi-squares when comparing the LFs of Cl0016+16, Cl1601+42 and MS1008--1224 to the best-fitting Gaussian+Schechter function, as well as an LF which is an average of the three individual LFs.} \\begin{tabular}{lcc} \\bf Cluster & $\\chi^2/\\nu$  & $\\chi^2/\\nu$ \\\\ ~ & (individual fit)  & (average fit) \\\\ Cl0016+16 & 0.45 & 1.25\\\\ Cl1601+42 & 0.46 & 1.49\\\\ MS1008-1224 & 0.98 & 1.46\\\\ \\end{tabular} \\label{Table8} \\end{table} \\subsubsection{The bright population} For Coma the Gaussian peaks at $M_B^G~=-19.3~\\pm0.1$, indicating a fading of the bright population by $\\Delta M\\sim~1$ mag, compared to the high-redshift clusters (Table \\ref{Table6}). The total LF of the Virgo cluster only has a marginal Gaussian peak. However, the population of morphologically classified elliptical galaxies, as well as giant spirals, has a clear Gaussian shape peaking at $M_B^G~\\sim~-19$ (Binggeli et al. 1988; Ferguson \\& Sandage 1991), supporting a fading of this population similar to that suggested for Coma.  These results are consistent with Smail et al. (1997), who find that the characteristic magnitude $M_V^*$ for early-type galaxies fade by $\\Delta M\\sim~0.7$ mag between $z~=~0.54$ and $z~=~0$. (We have here converted the results in Smail et al. to our adopted cosmology.) Smail et al. find a weaker trend for bright late-type galaxies, which fade by $\\Delta M\\lsim~0.4$ mag. An evolution is also found by Kodama \\& Bower (2001), who estimate that the bright blue galaxy population fades by $\\Delta M~\\sim~1$ mag between $z~\\sim~0.4$ and $ z~=~0$. De Propris et al. (1999) investigate the $K$--band luminosity evolution of the bright galaxy population in clusters at $0.1~<~z~<~1$, and find that $K^*(z)$ is consistent with passive luminosity evolution. We have compared our luminosity evolution with stellar synthesis models from Bruzual \\& Charlot (2003). We model the Gaussian population with an early-type population characterized by a single burst stellar population of age 8.0 Gyr and a Salpeter IMF. This combination results in colours that match the early-type spectral energy distribution in Coleman et al. (1980). For passive evolution, we find that this population should fade by $\\sim~0.7$~mag since $z\\sim~0.55$. This is slightly less than the $\\sim~1$~mag evolution that our results suggests. This difference will diminish if we either assume a Scalo IMF, or that the Gaussian population includes a fraction of galaxies with residual star formation. \\subsubsection{The faint population} In DFN02 we argued that the different shapes of the LF at $M_B~\\gsim~-20$ between Coma and Cl1601+42 can be understood as a result of the dynamically younger age of the latter cluster. The steep blue end of the LF in this cluster should then consist of recently accreted field galaxies, which is supported by the fact that this part of the LF is almost exclusively made up of late-type galaxies. The rapidly increasing fraction of late-type galaxies in the outer parts of the cluster is consistent with this accretion scenario, and is also reflected in the rising blue fraction at large radii in this cluster, as shown in Section 5.4. A fading of the blue star forming galaxies in Cl1601+42 by $\\sim~2$ mag, would transform the LF of Cl1601+42 into a LF similar to Coma.  The steepening of the faint-end in MS1008--1224 is clearly shown in Fig. \\ref{Figure6}. The slope is somewhat steeper than for Coma, and is shifted towards brighter magnitudes. From Fig. \\ref{Figure5} we see that the faint-end of MS1008--1224 is made up of both blue and red galaxies. The slope of the faint blue population is similar to the blue population in Cl1601+42, and a fading by $\\sim~1$ mag of this population would result in a LF similar to the one in MS1008--1224. A subsequent fading by an additional magnitude will make both Cl1601+42 and MS1008--1224 similar to Coma. The fading of the late-type population between $z~=0.5$ and $z~=0$ for these two clusters could therefore be described by a relation $\\Delta m~\\simeq~-4z$ mag.  \\subsubsection{The red cluster Cl0016+16} The LF of Cl0016+16 at $z~=0.55$ differs clearly from Cl1601+42 at the same redshift. The LF of the former is similar to Coma in the plateau region, having a Gaussian shape, but shifted to brighter magnitudes by $\\sim~1$ mag. From Fig. \\ref{Figure5} we see that the relatively low abundance of late-type galaxies in Cl0016+16, as compared to Cl1601+42, is responsible for the different shapes of the LF at intermediate magnitudes.  Previous investigations of Cl0016+16 show that the bright galaxy population is dominated by red early-type galaxies. Koo (1981) found that this cluster has an unusually high fraction of red galaxies compared to other high-redshift clusters, and suggested that the star formation must have ended a few Gyrs before the observed epoch. From the absence of blue galaxies and the small scatter in the colours of the red population, Smail et al. (1995) conclude that the cluster is old, despite its high redshift. Of the ten clusters in the MORPHS sample (Smail et al. 1997) with morphological classification from HST in the redshift range $0.37~<~z~<~0.56$, Cl0016+16 has the lowest fraction of spiral galaxies, $f_{sp}$ = 21 per cent, compared to a mean of $f_{sp}$ = 44$\\pm7$ per cent for the remaining nine clusters.  The low spiral fraction is, not unexpectedly, related to the low blue fraction found in Cl0016+16, which indicates that the cluster does not follow the general blueing described by the BO-effect. This may be explained by a higher dynamical age for Cl0016+16, which is supported by the high X-ray luminosity and velocity dispersion (Table \\ref{Table3}), indicating that Cl0016+16 is more relaxed than Cl1601+42.  Cl0016+16 has a very high fraction of post-starburst galaxies. These galaxies have no current star formation, but were forming stars $\\sim$1--2 Gyr before we observe them. Poggianti et al. (1999) found that 32$\\pm9$ per cent of the galaxies brighter than $M_V\\sim-20.5$~in Cl0016+16 are spectroscopically consistent with this type of galaxies. This suggests that a large fraction of the galaxies in Cl0016+16 has experienced star formation since $z~\\sim~0.8$, but have subsequently faded to become red at $z~=~0.54$. Post-starburst galaxies are less frequent in local clusters (Poggianti et al. 1999), and although not actively forming stars, Cl0016+16 is therefore not similar to a local cluster placed at high redshift. This is also supported by the brighter early-type population in this cluster compared to local clusters.  The difference between rich clusters at high and low redshift can qualitatively be explained within the hierarchical clustering scenario. Kauffmann (1995) shows that high-$z$ rich clusters assemble over a shorter time interval than low-$z$ clusters of similar richness. In general, this naturally leads to a higher fraction of blue galaxies in high-$z$ clusters. The properties of Cl0016+16 can then be understood if it formed during a short time interval at $z~\\gsim~0.8$, leaving a high fraction of post-starburst galaxies. This suggests that the cluster was blue at $z~\\sim~0.8$, while the red colours at $z~=~0.55$ indicates a very low infall of field galaxies during the last $\\sim~1$ Gyr before it is observed.  It is possible that the special environment of Cl0016+16 could have an influence on the blue fraction and the cluster LF. Connolly et al. (1996) and Hughes \\& Birkinshaw (1998) provide strong evidence that Cl0016+16 is part of a supercluster structure, with two associated clusters at similar redshift with projected distances of 5 and 13 Mpc. The formation of this giant structure could have depleted the number of field galaxies surrounding Cl0016+16, and therefore the accretion rate of late type galaxies. \\subsubsection{A universal cluster LF?} To summarise our discussion, we find no evidence for a universal shape of the total cluster LF, in agreement with e.g., Binggeli et al. (1988) and Driver, Couch \\& Phillipps (1998). The LFs for the spectroscopic early-type and late-type galaxies separately have, however, similar shape in the different clusters, but with varying relative strengths. This suggests an universality for type-specific LFs. This is analogous to the claim of type-specific LFs for different morphological types by Binggeli et al. (1988) and Andreon (1998). A comparison with Binggeli et al. (1988) is especially interesting, since they find that the bright part of the LF mainly consists of elliptical galaxies with a Gaussian LF, and that the steep faint-end consists of irregulars and dEs. This suggests, as expected, that the early-type population mainly consists of ellipticals, while many faint late-type galaxies are irregulars. To explain the different shapes of the high-$z$ cluster LFs, we argue that these reflect different dynamical states of the clusters. As the clusters get dynamically older, we expect the late-type population to fade relative to the early-type population, and LFs to become more similar. \\begin{figure} \\centerline{ \\psfig{figure=Figure6.eps,height=9.8cm}} \\caption{Luminosity function in the rest-frame $B$--band for Cl1601+42 at $z~=~0.54$ (open circles), Cl0016+16 at $z~=~0.55$ (filled circles),  MS1008--1224 at $z~=~0.31$ (triangles), Coma at $z~=~0.02$ (squares) and Virgo at $z~=~0.003$ (diamonds). Data for Coma and Virgo are taken from Trentham (1998a) and Trentham \\& Hodgkin (2002), respectively. The LFs are arbitrary off-set in the $y$-direction.} \\label{Figure6} \\end{figure} \\subsection{The dwarf population in MS1008--1224} The faint limit in absolute magnitude we reach for MS1008--1224 at $z~=~0.31$ allows us to study the dwarf population $\\sim~1.5$ mag deeper in this cluster compared to the clusters at $z~\\sim~0.55$.  Studies by Trentham (1998c) and Boyce et al. (2001) have shown that dwarf irregular (dIrr) galaxies have $B-R~\\sim~0.9$, while dEs have colours in a broader range $1.3~\\lsim~B-R~\\lsim~2.0$. To compare the colours of the dwarf population in MS1008--1224 with other nearby clusters, with more limited colour information, we therefore plot in Fig. \\ref{Figure7} the rest-frame $B-R$ colours of galaxies with $-17.7~<~M_B~<~-16.2$.  In the $B-R$ histogram of MS1008--1224 in Fig. \\ref{Figure7} there are two peaks, suggesting that there indeed are two distinct populations representing dIrrs and dEs. In the figure we also show the distribution of $B-R$ colours for the faint galaxies, which on the basis of the full $BVRI$ photometry we have classified as early-type and late-type galaxies, respectively. The fact that this classification closely follows the two-colour division into the two peaks shows that we reliably can use the $B-R$ colour to broadly distinguish between early-type and late-type galaxies.  From the $B-R$ colours Trentham (1998c) finds that all dwarf galaxies with $-18.9~<~M_R~<~-16.9$ in Abell 665 at $z~=0.18$ are consistent with being dEs (called dSphs by Trentham). In Abell 963 at $z~=~0.21$ the majority of the dwarf galaxies with $-19.2~<~M_R~<~-18.2$ also have colours consistent with dEs. There is, however, in this cluster also a population with colours intermediate between dEs and dIrrs, which Trentham proposes may represent a transitional stage between these types.  In a recent study Boyce et al. (2001) investigate the $B-R$ colour of galaxies with $-19~<~M_R~<~-16.5$ in Abell 868 at $z~=0.154$. They find a distribution in colour with two peaks, showing the presence of a dominant population of dEs, but also a population of dIrrs. For the faintest galaxies ($-17.5~<~M_R~<~-16.5$) there are similar numbers of dEs and dIrrs.  The main difference between MS1008--1224 and lower-redshift clusters is the relative fraction of dIrrs to dEs. To make a quantitative comparison between the faint populations, we use the $B-R$ index to calculate the fraction of dIrrs for MS1008--1224, Abell 963, Abell 665 and Abell 868. We assume that galaxies with rest-frame $B-R~<~1.2$ are dIrrs, which is the criterion used by Boyce et al. (2001).  A complication is that MS1008--1224 is not observed in standard Johnson--Cousins filters, but in Bessel filters. To correct for this we use galaxy templates from Coleman et al. (1980) and filter transmission curves to calculate the transformation between the standard $B-R$ colours and the $B-R$ colours for the Bessel filters. We then find that the division between dIrrs and dEs at $B-R$ = 1.2 in standard Johnson--Cousins filters corresponds to $B-R$ = 1.34 in the Bessel filters.  Using this criterion, we find that the fraction of faint galaxies that belong to the dIrr population is $0.72~\\pm0.03$ for MS1008--1224, to be compared with $0.07~\\pm0.03$ for Abell 963, $0.01~\\pm0.01$ for Abell 665 and $0.22~\\pm0.02$ for Abell 868. A further problem when comparing these numbers is that the different investigations use different magnitude intervals for defining the dwarf populations. Irrespective of this, however, we reach the important conclusion that MS1008--1224 is dominated by dIrrs, while the lower redshift clusters are dominated by dEs.  The fraction of dIrrs in Coma, classified by the same colour criterion as above, for galaxies with $-15.3~<~M_R~<~-13.3$ is $0.11~\\pm0.01$ (calculated from fig. 7 in Trentham 1998a). The fraction of {\\it morphological} dIrrs in Virgo ($-17.8~<~M_B~<~-15.8$), calculated from fig. 6 in Trentham \\& Hodgkin (2002), is 0.25$\\pm0.07$ The results emphasize the conclusion that dEs dominate the faint population in nearby clusters, while our study shows that the opposite is true for MS1008--1224 at $z~=~0.31$. It would obviously be of great interest to study more clusters to these limits to see if this is a general property.  Galaxy harassment could here be a mechanism that transforms the faint late-type population, dominating medium redshift clusters such as MS1008--1224, into dEs, as shown by Conselice et al. (2001). The similarity between the faint-end slopes of the late-type population ($\\alpha_f~=~-1.84\\pm0.12$) and the early-type population ($\\alpha_f~=~-1.69\\pm0.14$) at faint magnitudes ($M_B~>~-19$) is consistent with this picture. Based on these observations alone, we cannot, however, rule out that slower mechanisms such as 'strangulation' (Balogh et al. 2000), are at least partly responsible for the luminosity and colour evolution.  \\begin{figure} \\centerline{ \\psfig{figure=Figure7.eps,height=8.5cm}} \\caption{The rest-frame $B-R$ colour distribution for galaxies in MS1008--1224 with $-17.7~<~M_B~<~-16.2$. The solid line shows the total number of galaxies, while the dotted and dashed lines show the distribution of early-type and late-type populations, as classified by multi-band photometry, separately.} \\label{Figure7} \\end{figure}  \\subsection{Dwarf-to-giant ratio} Previous studies have shown that the ratio of faint to bright galaxies increases at large radii, where the surface density decreases. Driver et al. (1998) find that five out of seven clusters at $z~\\sim~0.15$ show an increasing dwarf-to-giant ratio with radius. A $K$--band study of AC 118 at $z~=~0.3$ by Andreon (2001) shows the same trend. In DFN02 we found that Cl1601+42 has a steeper faint-end slope of the LF in the outer part of the cluster, which is equivalent to an increasing dwarf-to-giant ratio.  Here, we calculate the dwarf-to-giant ratio by defining galaxies with $-19.5~<~M_B~<~-17.7$ as dwarfs, and brighter galaxies as giants. The separation between dwarfs and giants is chosen to match the cross-over in the cluster LFs between early-type and late-type galaxies (Fig. \\ref{Figure5}). This division is somewhat brighter than used in other papers (e.g., Driver et al. 1998), and is here adopted in order to get sufficient statistics in the faint bin. We further divide each cluster into a core region with projected radius $<~0.5$ Mpc, and an outer region between 0.5 and 1.1--1.5 Mpc, depending on cluster.  In Fig. \\ref{Figure8} we plot the resulting dwarf-to-giant ratios as function of surface density of the giant galaxies for these two regions in all clusters. Cl0016+16 is represented by circles, Cl1601+42 by squares and MS1008--1224 by triangles. Filled symbols show results for the core area, while open symbols represent results from the outer area. The horizontal lines represent the field value and the corresponding 1$\\sigma$-errors. \\begin{figure} \\centerline{ \\psfig{figure=Figure8.eps,height=8.5cm}} \\caption{Dwarf-to-giant ratio as a function of bright galaxy surface density for Cl0016+16 (circles), Cl1601+42 (squares) and MS1008--1224 (triangles). Filled symbols represent the inner region of the cluster ($R~<~0.5$ Mpc), while open symbols represent the outer region. The horizontal lines represent the field value (solid line) and 1$\\sigma$-errors (dotted lines).} \\label{Figure8} \\end{figure}  Fig. \\ref{Figure8} shows that Cl0016+16 and Cl1601+42 have a clear trend of an increasing dwarf-to-giant ratio when going from the inner high-density region, to the outer lower-density region, which is consistent with the trend found by Phillipps et al. (1998) and with Andreon (2001). As an explanation of this these authors suggest that the high-density environment in cores of clusters is hostile to the dwarf galaxies, while the outer, low-density regions do not affect this population, which therefore have a dwarf-to-giant ratio and faint-end slope similar to that of the field.  For the low redshift cluster MS1008--1224, there is no obvious trend. We note, however, that the error in the dwarf-to-giant ratio is substantial. \\subsection{The Butcher--Oemler effect as function of radius and limiting magnitude} The BO effect is of major observational importance as a probe of the cluster galaxy population. Although much of the discussion below to a large extent is a consequence of the early-type and late-type LFs in Fig. \\ref{Figure5} and their dependence on radius (e.g., Fig. \\ref{Figure4}), we will for this reason illustrate the consequences of this with two important implications for the BO effect.  Already in BO84, and more recently in DFN02, it was pointed out that the blue fraction $f_B$ increases at larger radii. This can be understood as a result of the fact that early-types are in general more centrally concentrated than late-types. In Fig. \\ref{Figure9} we show the dependence of $f_B$ on radius, where the radius is given in fractions of $R_{30}$ (as defined in Section 4.3). \\begin{figure} \\centerline{ \\psfig{figure=Figure9.eps,width=8.5cm}} \\caption{The blue fraction, $f_B$, as a function of radius for Cl0016+16 (filled circles), Cl1601+42 (open circles) and MS1008--1224 (triangles). $R_{30}$ is the fiducial value used by BO84 for calculating $f_B$.} \\label{Figure9} \\end{figure} Also in this respect Cl0016+16 deviates from the other clusters. Inside of 1.25$R_{30}$, $f_B$ decreases rapidly, while outside this radius $f_B$ stays almost constant, at a value well below that of the other clusters. Cl1601+42 on the other hand shows a steep radial dependence. The higher blue fraction in Cl1601+42 is consistent with the higher fraction of late type galaxies in this cluster (Fig. \\ref{Figure4}). MS1008--1224 has a trend in between the two higher redshift clusters. From a very low $f_B$ in the core region, there is a rapid increase to $\\sim~1.25~R_{30}$, followed by a slower increase in the outer regions. The behaviour of $f_B$ in these clusters reflects, as expected, the radial dependence of the early-type fraction shown in the right panel of Fig. \\ref{Figure4}.  In DFN02 we showed that $f_B$ in Cl1601+42 depends strongly of the limiting magnitude. This dependence is a direct consequence of the different shapes of the early-type and late-type LFs in this cluster. At fainter magnitudes the total LF becomes more dominated by late-type galaxies, which leads to a higher $f_B$. In Fig. \\ref{Figure10} we plot the dependence of $f_B$ on the limiting magnitude for the three clusters. As expected, the increase of $f_B$ is much weaker in Cl0016+16 than in Cl1601+42, which is consistent with the LF for this cluster being dominated by early-type galaxies to fainter magnitudes (Fig. \\ref{Figure5}). MS1008--1224 again shows an intermediate behaviour. The strong increase in $f_B$ at faint magnitudes in MS1008--1224 is caused by the dominating population of blue dwarf galaxies.  \\begin{figure} \\centerline{ \\psfig{figure=Figure10.eps,width=8.5cm}} \\caption{The fraction of blue galaxies, $f_B$, as a function of limiting magnitude for Cl0016+16 (filled circles), Cl1601+42 (open circles) and MS1008--1224 (triangles).} \\label{Figure10} \\end{figure}"
    },
    "0402/astro-ph0402386_arXiv.txt": {
        "abstract": "We have constructed a sample of bright near-infrared sources which are detected at radio wavelengths but undetected on the POSS I plates in order to search for a population of dust-obscured quasars. Optical and infrared spectroscopic followup of the sample has led to the discovery of seventeen heavily reddened quasars ($B-K>6.5$), fourteen of which are reported here for the first time.  This has allowed us to define a region in the $R-K, J-K$ color plane in which $50\\%$ of the radio-selected objects are highly reddened quasars. We compare the surface density of this previously overlooked population to that of UVX- radio-selected quasars, finding that they make up $\\sim 20\\%$ of the total quasar population for $K\\lesssim 15.5$. ",
        "introduction": "Spectroscopy of our candidate quasars was carried out in the optical and near-infrared at four different observatories.  The majority of the optical spectra were obtained at the 10m Keck telecopes using the ESI and, to a lesser extent, LRIS spectrographs. Optical spectra were also obtained at the 3 m Shane telescope at Lick Observatory.  Near-infrared spectroscopy was performed at the NASA Infrared Telescope Facility (IRTF) using SpeX \\citep{spex1} and at the MDM Observatory on Kitt Peak using the TIFKAM infrared camera in spectroscopic mode \\footnote{TIFKAM was funded by the Ohio State University, the MDM consortium, MIT, and NSF grant AST-9605012. NOAO and USNO paid for the development of the ALADDIN arrays and contributed the array currently in use in TIFKAM.}.  We also obtained seven spectra from the Sloan Digital Sky Survey Spectral database.  The LRIS data were reduced and spectra extracted using standard IRAF procedures.  ESI spectra were produced using a combination of IRAF tasks and customized software developed by the authors.  More details on the algorithms used in the ESI software are given in \\citet{white03a}.  The near infrared spectra obtained at IRTF with SpeX were reduced using Spextool \\citep{spex3}.  We flux calibrated and corrected the data for telluric absorption using spectra of nearby A0V stars observed immediately before or after each object spectrum, following the procedure outlined in \\citet{spex2}.  The near infrared spectra obtained with TIFKAM at MDM were extracted and reduced using standard spectral reduction tasks in the IRAF package ONEDSPEC.  We used both A0V and G2V stars (to calibrate the hydrogen absorption lines in the A0V star) when correcting these spectra for telluric absorption, following \\citet{hanson96}.  We classified our spectra using algortihms in the radial velocity analysis (RV) package of IRAF.  We used the FXCOR task to cross-correlate the spectra with a set of galaxy and star templates, implementing the technique illustrated in \\citet{td79}.  Our templates consisted of elliptical, S0, Sa, Sb, Sc and starburst galaxies from \\citet{kin96} and \\citet{cal94}\\footnote{We chose the starburst galaxy template with $0.61 < E(B-V) < 0.70$.  Of the six available templates, this spectrum had the highest reddening;  it was chosen as most consistent with our survey's selection criteria.}.  This task was most useful in classifying and assigning redshifts for the absorption line galaxies.  We used the Tonry \\& Davis $R$ value to determine the strength of the correlation, rejecting identifications with $R<3.0$.  Some spectra with emission features were misidentified as starbursts, but were assigned correct redshifts; others did not correlate at all.  In the latter cases we used the RV task RVIDLINES to measure radial velocities from spectral lines using emission features present in the quasar composite spectrum from \\citet{broth01} as well as the first three Paschen lines.  We then fit gaussian profiles to the emission lines and classifed objects with line widths $> 1000$ km/s as quasars.  Objects with narrow lines were classified as starbursts or narrow-line AGN based on [O III]/H$\\beta$ and [N II]/H$\\alpha$ line diagnostics \\citep{oster89}.  Table 2 summarizes all the data collected on our quasar candidates.  Columns (1) and (2) list the candidates' J2000 coordinates from the FIRST catalog.  In columns (3) and (4) we list the $B$ and $R$ magnitudes, respectively, from the complete Guide Star 2 catalog which reaches the POSS-II plate limits.  We chose to present these magnitudes despite our selection criteria which required that the objects be fainter than the APM limit, in order to allow for near-infrared to optical color studies which we discuss in the following sections.  In columns (5), (6) and (7) we present 2MASS $J$,$H$, and $K_s$ photometry, followed by the peak and integrated FIRST 20 cm flux densities in columns (8) and (9).  In columns (10) and (11) we present $J-K$ and $R-K$ colors.  We used the $R$ magnitude from GSC2 for $R-K$.  We argue in \\S 4 that these near-infrared and optical to near-infrared colors can be useful diagnostics for efficient selection of red quasars.  We note, however, that the non-simultaneity of the POSS-II and 2MASS observations make our colors provisional at best.  We report the redshift of the candidate when available in column (12), followed by the classification we determined using the aforementioned analysis in column (13); any other references in the literature are provided in column (14). ",
        "conclusions": "We have conducted a survey to search for optically obscured, infrared-bright quasars, detected in both the FIRST and 2MASS surveys.  We required that the objects be undetected by the APM scans of the POSS-I plates.  Using these criteria we obtained a list of 69 candidates.  Of the 51 spectroscopically observed candidates, 17 were red quasars. From the optical and near infrared colors of the candidates  we have shown that red quasar detection efficiency rises to $\\sim 50\\%$ when adopting the color criteria $R-K>4$ and $J-K>1.7$.  While the near-infrared is helpful in finding red quasars because of its weaker sensitivity to dust extinction, its effectiveness diminishes with redshift and increased obscuration.  To evaluate the significance of red quasars to the whole quasar population it is necessary to characterize the distribution of $E(B-V)$ as a function of intrinsic luminosity and redshift.  With a sample of seventeen quasars this cannot yet be accomplished.   We have made some estimates of the percentage of red quasars which range from a firm lower limit from directly detected and classified optically faint objects of $\\sim 3\\%$, to an extinction- and survey-bias-corrected best estimate of $\\sim 20\\%$.  Future work will include a wider survey for red quasars taking advantage of the all-sky coverage of the latest 2MASS release.  We will also eliminate optical magnitude constraints and will select candidates based solely on their optical and near-infrared colors.   We intend to study the reddening of these objects in greater detail by calculating Balmer decrements and comparing unreddened spectra to quasar composites made from optically selected samples to search for any intrinsic differences between the two populations.  We thank Dave Deutsch of the Manhattan Center for Science and Mathematics for his careful and consistent inspection of the Digitized Sky Survey images to weed out optical identifications of red quasar candidates.  E. G. and D. J. H. acknowledge support from National Science Foundation grant AST-00-98259.  R. H. B. acknowledges support of the Institute of Geophysics and Planetary Physics (operated under the auspices of the US Department of Energy by the University of California, Lawrence Livermore National under contract W-7405-Eng-48) as well as additional support from NSF grant AST00-98355 (University of California, Davis).  This publication makes use of data products from the Two Micron All Sky Survey, which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center/California Institute of Technology, funded by the National Aeronautics and Space Administration and the National Science Foundation.  The Guide Star Catalogue-II is a joint project of the Space Telescope Science Institute and the Osservatorio Astronomico di Torino. Space Telescope Science Institute is operated by the Association of Universities for Research in Astronomy, for the National Aeronautics and Space Administration under contract NAS5-26555. The participation of the Osservatorio Astronomico di Torino is supported by the Italian Council for Research in Astronomy. Additional support is provided by European Southern Observatory, Space Telescope European Coordinating Facility, the International GEMINI project and the European Space Agency Astrophysics Division.  Funding for the Sloan Digital Sky Survey (SDSS) has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Aeronautics and Space Administration, the National Science Foundation, the U.S. Department of Energy, the Japanese Monbukagakusho, and the Max Planck Society. The SDSS Web site is \\url{http://www.sdss.org/}.  The SDSS is managed by the Astrophysical Research Consortium (ARC) for the Participating Institutions. The Participating Institutions are The University of Chicago, Fermilab, the Institute for Advanced Study, the Japan Participation Group, The Johns Hopkins University, Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, University of Pittsburgh, Princeton University, the United States Naval Observatory, and the University of Washington."
    },
    "0402/astro-ph0402453_arXiv.txt": {
        "abstract": "We present {\\it HST} STIS observations of two extreme NLS1s, IRAS~13224$-$3809 and 1H~0707$-$495.  The spectra are characterized by very blue continua, broad, strongly blueshifted high-ionization lines (including \\ion{C}{4} and \\ion{N}{5}), and narrow, symmetric intermediate- (including \\ion{C}{3}], \\ion{Si}{3}], \\ion{Al}{3}) and low-ionization (e.g., \\ion{Mg}{2}) lines centered at their rest wavelengths. The emission-line profiles suggest that the high-ionization lines are produced in a wind, and the intermediate- and low-ionization lines are produced in low-velocity gas associated with the accretion disk or base of the wind.  In this paper, we present the analysis of the spectra from these two objects; in a companion paper we present photoionization analysis and a toy dynamical model for the wind.  The highly asymmetric profile of \\ion{C}{4} suggests that it is dominated by emission from the wind, so we develop a template for the wind from the \\ion{C}{4} line. We model the bright emission lines in the spectra using a combination of this template, and a narrow, symmetric line centered at the rest wavelength.  We also analyzed a comparison sample of {\\it HST} spectra from 14 additional NLS1s, and construct a correlation matrix of emission line and continuum properties.  A number of strong correlations were observed, including several involving the asymmetry of the \\ion{C}{4} line. ",
        "introduction": "In 1992, it was demonstrated by Boroson \\& Green that the optical emission line properties in the region of the spectrum around H$\\beta$ are strongly correlated with one another.  A principal components analysis allowed the largest differences among optical emission line properties to be gathered together in a construct commonly known as ``Eigenvector 1''.  The strongest differences hinged on the strength of the \\ion{Fe}{2} and [\\ion{O}{3}] emission, and the width and asymmetry of H$\\beta$.  This set of strong correlations is remarkable, as it involves correlations among the dynamics and gas properties between emission regions separated by vast distances. Furthermore, the same set of correlations are observed in samples selected in many different ways.  This pervasiveness leads us to believe that we are observing the manifestation of a primary physical parameter. A favored explanation is that it is the accretion rate relative to the black hole mass onto the active nucleus (e.g., Boller, Brandt \\& Fink 1996; Pounds, Done \\& Osborne 1995; Comastri et al.\\ 1998; Leighly 1999ab;  Mathur 2000;  Mineshige et al.\\ 2000; Kuraszkiewicz et al. 2000; Pounds et al.\\ 2001; Puchnarewicz et al.\\ 2001; Ballantyne, Iwasawa, \\& Fabian 2001; Boroson 2002; Crenshaw, Kraemer \\& Gable 2003).  Narrow-line Seyfert 1 galaxies (NLS1s\\footnote{In this paper, the term NLS1s is used to refer to all objects fulfilling the optical criteria regardless of their luminosity.  Likewise, the term broad-line Seyfert 1 is used to refer to objects of both Seyfert and quasar luminosities.}) are identified by their optical emission line properties.  They typically have narrow permitted optical lines (FWHM of H$\\beta<2000\\rm\\, km\\,s^{-1}$), weak forbidden lines ([\\ion{O}{3}]/H$\\beta<3$; this distinguishes them from Seyfert 2 galaxies), and frequently they show strong \\ion{Fe}{2} emission (Osterbrock \\& Pogge 1985; Goodrich 1989).  These are the same properties that define the Boroson \\& Green (1992) Eigenvector 1, and indeed, NLS1s fall at one end of Eigenvector 1.  Thus, the study of NLS1s offers an attractive research opportunity: if we can understand the origin of the emission-line properties of NLS1s, then we may be in a position to understand AGN emission in general.  Furthermore, because as least some of the lines in NLS1s are narrow, identification of the frequently strongly-blended lines is less ambiguous than it is in broad-line objects.  NLS1 research dramatically increased after the publication of a seminal paper in 1996 by Boller, Brandt \\& Fink describing their discovery that NLS1s have systematically steeper {\\it ROSAT} spectra than AGN with broad optical lines.  Subsequently it was shown from {\\it ASCA} observations that the hard X-ray spectrum is also steeper (Brandt, Mathur \\& Elvis 1997; Leighly 1999b), and that the X-ray emission shows higher fractional amplitude of variability (Leighly 1999a).  These results were generally interpreted as evidence that the black hole is systematically smaller in NLS1s, implying a higher accretion rate relative to the Eddington value compared with Seyfert 1 galaxies with broad optical lines.  Once the high accretion rate paradigm was established, it could be used to understand why NLS1s have narrow optical lines. Boller, Brandt, \\& Fink (1996) mention that if the black hole mass in NLS1s is smaller, and the emission lines are produced at the same radius as in broad-line objects (due to the same radiating flux)  the dynamical velocities will be smaller, and narrower lines will be produced.  This theme was expanded upon by Wandel \\& Boller (1998), who consider the fact that the steeper soft X-ray spectrum of NLS1s, when extrapolated into the UV, implies a larger ionizing flux relative to their luminosity.   This causes the emission region to be larger where the velocities are  smaller.  However, the optical emission lines are only one side of the coin; UV emission lines, generally characterized by a higher degree of ionization, may also provide valuable information.  This is because the optical and UV emission lines are generally not produced in the same gas.  The best evidence for this comes from so-called reverberation mapping experiments (e.g., Peterson 1993).  Also, the profiles of the emission lines suggest this as well, as the higher ionization lines are observed to be generally broader than the lower ionization lines (e.g., Baldwin 1997).  Therefore, we may learn more about NLS1s by studying their UV emission lines.  Several groups have investigated the UV emission line properties of NLS1s.  The first paper investigating a sample of IUE spectra from NLS1s was published by Rodr\\'iguez-Pascual, Mas-Hesse \\& Santos-Lle\\'o (1997).  They reported that the strongest UV emission lines have broad wings, unlike the optical emission lines, and they speculate that the UV line-emitting gas may be optically thin in the continuum.  Kuraszkiewicz et al.\\ (2000) present a detailed study of {\\it IUE} and {\\it HST} spectra from NLS1s.  They found that, compared with broad-line quasars, the equivalent widths of \\ion{C}{4} and \\ion{Mg}{2} are low.  They found that, in the cases where the lines could be deblended, the \\ion{Si}{3}]/\\ion{C}{3}] ratio was high.  In addition, the \\ion{Si}{4}+\\ion{O}{4}] feature near 1400 \\AA\\ was found to be strong compared with \\ion{C}{4}, and the \\ion{Al}{3} equivalent width was high.  They also found that although the UV emission lines in NLS1s are broader than the Balmer lines, they still tend to be narrower than those lines in broad-line quasars.  They then model many of the emission line ratios using a 1-zone {\\it Cloudy} model (Ferland 2001).  They infer that the line ratios imply that the ionization parameter is lower and the density is higher in NLS1s than in broad-line Seyfert 1 galaxies.  We comment on their inferences in the companion paper to this one (Leighly 2004a; hereafter referred to as Paper II).  Wills et al.\\ (1999), studying the {\\it HST} spectra from a small sample of PG quasars, found that Eigenvector 1 is also associated with a high \\ion{Si}{3}] to \\ion{C}{3}] ratio, stronger UV low-ionization lines, weaker \\ion{C}{4}, but stronger \\ion{N}{5}.  They infer from these spectra that NLS1s are typified by high densities, low ionization parameters, and nitrogen enhancement, perhaps from nuclear starbursts.  All the work done so far on the UV spectra of NLS1s has been generally confined to studies of the properties of the integrated lines, to a greater or lesser extent, depending on the analysis.  Specifically, the conditions of a single-zone BLR are modelled or discussed more or less quantitatively, and evidence for a radially-extended BLR is inferred when particular lines are found to not be consistent with the one-zone model (e.g., Kuraszkiewicz et al.\\ (2000) infer that, since the observed \\ion{Mg}{2} is underpredicted by the one-zone model, it must be produced in a different region than the other UV emission lines).  However, we know that generally the emission line profiles are not all the same, and higher ionization lines tend to have broader profiles.  This implies that the gas properties are different among regions with larger or smaller line-of-sight velocities, and therefore a study of the average properties could be misleading.  In this paper, we present the observations and analysis of {\\it HST} observations of two NLS1s that have truly remarkable UV spectra. In Section 2, we introduce the targets, describe our observations and discuss the analysis of the emission lines.  In Section 3, we compare the results from these two NLS1s with spectra from 14 NLS1s drawn from the {\\it HST} archive, and perform a correlation analysis and a principal components analysis on the results.  In a companion paper II, photoionization analysis, and a toy dynamical model are presented, and the results are discussed.  Partial results from this study have been published previously in Leighly (2000) and Leighly (2001).  We use $H_0=50\\rm\\, km\\,s^{-1}\\,Mpc^{-1}$ and $q_0=0.5$ unless otherwise noted. ",
        "conclusions": "\\begin{itemize} \\item We present a detailed analysis of the UV spectra from two Narrow-line Seyfert 1 galaxies, IRAS~13224$-$3809 and 1H~0707$-$495. We find that their continua are as blue as that of the average quasar.  We observe that the high-ionization emission lines (including \\ion{N}{5} and \\ion{C}{4}) are broad (FWHM$\\approx\\rm 5000\\rm \\,km\\,s^{-1}$), and strongly blueshifted, peaking at around $2500\\rm \\,km\\,s^{-1}$ and extending up to almost $\\sim 10,000\\rm \\, km\\,s^{-1}$.  In contrast, the intermediate- and low-ionization lines (e.g., \\ion{C}{3}] and \\ion{Mg}{2}) are narrow (FWHM 1000--1900$\\rm\\,km\\,s^{-1}$) and centered at the rest wavelength. \\ion{Si}{3}] is prominent, and other low ionization lines (e.g., \\ion{Fe}{2} and \\ion{Si}{2}) are strong.  Based on these observations, the working model that we adopt considers that the blueshifted high-ionization lines come from a wind that is moving toward us, with the receding side blocked by the optically thick accretion disk, and the intermediate- and low-ionization lines are emitted in the accretion disk atmosphere or low-velocity base of the wind.  \\item The strongly blueshifted \\ion{C}{4} profile suggests that it is dominated by emission in the wind.  We used the \\ion{C}{4} profile to develop a template for the wind.  We then used this template, plus a narrow and symmetric component representing the disk emission to model the other bright emission lines. We inferred that the high-ionization lines \\ion{N}{5} and \\ion{He}{2} are also dominated by wind emission, and a part of Ly$\\alpha$ is emitted in the wind. The intermediate- and low-ionization lines \\ion{Al}{3}, \\ion{Si}{3}], \\ion{C}{3}], and \\ion{Mg}{2} are dominated by disk emission, and a part of Ly$\\alpha$ is disk emission as well.  The 1400\\AA\\/ feature, comprised of \\ion{O}{4}] and \\ion{Si}{4} was difficult to model; however, it appears to include both disk and wind emission.  \\item IRAS~13224$-$3809 and 1H~0707$-$495 have distinctive X-ray properties among NLS1s; in order to determine whether their distinctive properties carry over to the UV, we analyzed {\\it HST} archival spectra from 14 other NLS1s with a range of two orders of magnitude in UV luminosity.  We find indeed that these two objects are extreme in this sample in the following properties: strongly blueshifted \\ion{C}{4} line, the low equivalent widths of many of the lines, particularly \\ion{C}{4} and \\ion{He}{2}, high \\ion{C}{3}]/\\ion{C}{4}, \\ion{Si}{3}]/\\ion{C}{3}], \\ion{Al}{3}/\\ion{C}{3}], and \\ion{N}{5}/\\ion{C}{4} ratios, steep $\\alpha_{ox}$, and blue UV continuum.  Correlation analysis finds a number of strong correlations.  An anticorrelation between \\ion{C}{4} asymmetry and equivalent width suggests that the line is generally composed of a highly asymmetric wind component and a narrower symmetric component.  The anticorrelation between \\ion{C}{4} asymmetry and $\\alpha_{ox}$ and $\\alpha_{u}$, and with \\ion{He}{2} suggests that UV-strong and X-ray weak continua may be associated with a wind, as would be expected if the acceleration mechanism is radiative-line driving.  The dominance of \\ion{Al}{3} and \\ion{Si}{3}] over \\ion{C}{3}] suggests possibly that the continuum is transmitted through the wind before it illuminates the intermediate-ionization line emitting gas.  \\item Paper II investigates the physical conditions in the line-emitting gas through {\\it Cloudy} modeling of the disk and wind emission lines.  The results are used, in combination with a toy dynamical model for the wind, to estimate the distance of the wind from the continuum emitting region.  Further discussion of the results is given, including a comparison with previous results. \\end{itemize}"
    },
    "0402/astro-ph0402335_arXiv.txt": {
        "abstract": " ",
        "introduction": "The discovery of the first extra-solar planet in orbit around a main-sequence star was announced in 1995. Since then, the number of known extra-solar planets did not cease to grow. As the observations are accumulating, planetary systems with 2 and 3 planets are being discovered. More than ten are, presently, known\\footnote{For an up-to-date list see the web page ``Extra-Solar Planets Encyclopaedia'', by J.Schneider at www.obspm.fr/planets and links therein.}. As the discoveries are recent and many of the discovered planets are at the edge of observational capabilities, the uncertainties on their orbital elements and masses are large. It is worth recalling that one of the 2-planet systems previously announced, HD 83443, vanished from the lists after new observations failed to show the radial velocity variations previously identified with a second planet. Another important example of the current uncertainties is the ``jump'' suffered by the determined eccentricity of HD 82943b. During long time, it was listed as $\\sim 0.4$, while a new determination using observations over a long span of time gives only $0.18$. By the same occasion, the mass of HD 82943c became twice bigger than believed before (Mayor et al., 2004). These discrepancies should be enough to show us how hazardous is the task of getting conclusions from the present data and that we should avoid conclusions critically depending on the available data.  In the current state of art, we are just capable of discovering big planets with not too large periods. Therefore, the planets so far discovered are big and most of them have orbits close to the central stars. Another characteristic is the large eccentricities of many of them. Even if large eccentricities favors discovery, this characteristic is not only due to observational bias and needs an explanation. (See Perryman 2000 for a review of the existing hypotheses.) Large eccentricities are considered as the result of early migration processes. It is generally believed that the planets did not form at their present observed locations, but were driven by a migration process due to tidal interaction of the planets with the discs where they were formed (see Papaloizou, 2003). Whether this orbital drift is still at work or not is a matter of debate, although it is more plausible to assume that it stopped after the end of the planetary formation stage. These early processes were also responsible for having driven the (surviving) systems to very stable conditions in which orbit periapses appear close to alignment or anti-alignment. This condition is observed in several systems.  Periapses alignment (or anti-alignment) may occur in resonant and non-resonant systems alike. In resonant systems, they are the natural states after the system is trapped into a mean-motion resonance (see section \\ref{capture}). At variance, in non-resonant system, they are a consequence of the angular momentum  variations during resonance crossings without capture (Ferraz-Mello et al. in preparation). However, and independently on how they reached this condition, an important consequence of this type of configuration is that they constitute a stabilizing mechanism for planetary orbits, especially if they have large eccentricities.  Four extra-solar systems seem to satisfy the resonance condition: {Gliese 876}, { HD 82943}, {55 Cnc} and { 47 UMa}. The first two have planets with periods in a 2/1 commensurability, the third in a 3/1, and the later close to a 7/3. With regard to { Gliese 876}, numerical simulations (Laughlin and Chambers 2001, Lee and Peale 2002) seem to indicate that these bodies are actually deeply trapped in an apsidal corotation (see section \\ref{apsidal}): They exhibit a libration of both resonant angles $\\sigma_i = 2\\lambda_2 - \\lambda_1 - \\varpi_i$, and also an alignment of their major axes. Apsidal corotation seems to be the natural issue of a capture in resonance in the case of two planets with initially low eccentricities (Ferraz-Mello et al. 2003; {\\it cf.} this paper, section \\ref{capture}). The alignment (or anti-alignment) of periapses has not yet been confirmed in the case of the other planetary systems above mentioned.  The most conspicuous non-resonant system showing nearly aligned periapses is $\\upsilon$ Andromedae. This system has been the object of many numerical and analytical studies (for references, see Michtchenko and Malhotra, 2004). The orbit of the planets $c$ and $d$ in this system are such that the distance between their periapses oscillates about zero with half-amplitude $\\sim 60$ degrees and period $\\sim 7260$ years. ",
        "conclusions": "The contents of this paper include with variable emphasis, the topics of a series of lectures whose main title was ``Routes to Order: Capture into Resonance''. This was indeed the subject of the last section above. The study of this subject has however shown that differently from the restricted three-body problem, the capture into resonance drives the system immediately to stationary solutions known as ``apsidal corotations''. The whole theory of these solutions was also included in the paper from the beginning, that is, from the formulation of the Hamiltonian equations of the planetary motions and the expansion of the disturbing function in the high-eccentricity planetary three-body problem. The secular theory of non-resonant systems was also given. Motions with aligned or anti-aligned periapses, resonant or not, resulting from non-conservative processes (tidal interactions with the disc) in the early phases of the system life, seem to be frequent in extra-solar planetary systems."
    },
    "0402/astro-ph0402429_arXiv.txt": {
        "abstract": "The recent analysis of the cosmic microwave background data carried out by the WMAP team seems to show that the sum of the neutrino masses is $<0.7$~eV. However, this result is not model-independent, depending on precise assumptions on the cosmological model. We study how this result is modified when the assumption of perfect lepton symmetry is dropped out. ",
        "introduction": "The anisotropy spectrum of the cosmic microwave background is without doubt one of the most important observables quantity in cosmology. Recently, its precise measurement by the WMAP experiment, has lead to a precise determination of the values of the parameters describing our Universe \\citep{Spe03}. In particular, the WMAP team finds a very tight bound on the sum of neutrino mass, representing an improvement of roughly one order of magnitude with respect to the laboratory limit. It seems that finally the exciting possibility of performing precision tests in particle physics models using cosmological data, is within our reach.  Unfortunately, parameter estimation is plagued with the problem of parameter degeneration. With this it is meant that different parameters have similar effect on the microwave background spectrum; the latter cannot be used, then, to distinguish between them. A well known example is the degeneracy between the spectral index $n$ and the reionization redshift $z_{nr}$. Put in another way, we could say that there is not a one-to-one relationship between the parameters and the spectrum. The simplest way to break such a degeneracy is to measure some other quantity; for example, the matter power spectrum.  Another difficulty in dealing with the degeneration of parameters is due to the fact that we do not have an explicit relationship between the parameters and the spectrum. This means that probably there are some degeneracies we are not aware of, and that, even for kwown degeneracies, we do not know exactly how much they influence our results. Thus we are forced to use numerical methods, that are in some cases very time expensive.  Our work deals with the possibility to extract from the microvave background informations on the degree of lepton asymmetry existing in our Universe. In particular, we study how the estimation of cosmological parameters is affected by the existence of such an asymmetry. The paper is structured as follows: in section 2, we review the basis of neutrino physics; in section 3, we briefly recall the basis of the microwave background physics ; in section 4, we describe the techniques for parameter estimation; in section 5 we speak about the role of neutrinos in cosmology; in section 6 we present and discuss our results; finally in section 7 we draw our conclusions. ",
        "conclusions": "The analysis described in the preceding section points to the fact that the assumptions usually made on the lepton asymmetry, do really influence the estimation of cosmological parameters. In particular, the parameters that are more affected are the cosmological constant, the neutrino mass and the spectral index. Values of the cosmological constant between 0 and 0.7 cannot be distinguished, having nearly the same probability. The neutrino mass is peaked at 1~eV, this being not in agreement with the WMAP results. Finally, higher values of the spectral index are favoured. We think then that is worth studying in more depth this topic. It is without doubt necessary to apply our method to the WMAP data, using a larger parameter space, to be explored using MCMC methods."
    },
    "0402/astro-ph0402103_arXiv.txt": {
        "abstract": "{In this review talk I cover some recent developments in understanding the role that magnetic fields play in cataclysmic variables. I discuss the recent DNO--QPO unification models; the disk--magnetosphere boundary; some issues concerning the soft blackbody component and the nature of the X-ray spectra in MCVs; whether the SW Sex stars are magnetic, and finally I mention the weird behavior of FS~Aur and HS\\,2331+3905.}  \\resumen{Spanish abstract -- to be supplied by the editors.}  \\addkeyword{Accretion, accretion disks} \\addkeyword{Novae, cataclysmic variables}  \\begin{document} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0402/astro-ph0402273_arXiv.txt": {
        "abstract": "Time resolved photometry of the long-period intermediate polar V1062~Tau confirmed the presence of the previously reported orbital and spin periods,  and revealed the presence of a third one, corresponding to the beat of the two. While the orbital periodicity was present throughout our data, only one of the shorter periods was detectable at any given time. On a time-scale of $\\sim$90 days, the short-period modulation in the light curve of the star changed three times between the spin period and the beat. On longer time-scales, we report two outbursts of V1062~Tau (the first to be recorded for this object) -- both of which were probably short, low-amplitude ones ($\\Delta{m}\\sim1.2$~mag). Our data also suggest a brief low state phase in 2002 January. Thus, this system joins two exclusive groups of intermediate polars: those undergoing short outbursts, and those having low states.  We propose that the alternations between the short periods that modulate the light curve  were caused by changes in the accretion mode, from disc-fed accretion, to disc-overflow accretion. We further  suggest that these changes may have been triggered by changes in mass-transfer rate, which were manifested by the low-state/outburst activity of the system. ",
        "introduction": "Intermediate polars (IPs) are a subclass of Cataclysmic Variables (CVs) having a white dwarf, the magnetic field of which is strong enough to control (at least partially) the accretion flow from the secondary star, but is not strong enough to synchronize the spin of the primary star with the orbital revolution (for a review of IPs see e.g. Patterson 1994; Warner 1995; Hellier 1995; Hellier 1996).  A prominent observational characteristic of IPs is multiple periodicities in the light curve, which may be modulated by the spin and orbital periods, and a myriad of combinations of the two.  The intermediateness of the IPs, between the non-magnetic CVs and the strongly-magnetic polars, makes this class a diverse one, which includes stars with a wide range of asynchronism and magnetic-field strengths. The diversity of the class is further enhanced by the few accretion modes found in IPs. While most systems are predominantly disc-fed, some accrete predominantly in the disc-overflow mode (Hellier 1995) and one object is a discless, stream-fed IP.  V1062~Tau (1 H 0459+248) was discovered with the {\\it HEAO 1} Scanning Modulation Collimator (A3 experiment) and the Large-Area Sky Survey (A-1 experiment). A low-resolution spectrum with EXOSAT indicated that the source is a cataclysmic variable (Remillard et al. 1994, hereafter R94). The optical counterpart was identified as a UV-bright star, with spectral features characteristic of a cataclysmic variable. Time-resolved observations revealed two periodicities: a $P=9.952$~h period detected in optical photometry was suggested to be the orbital period of the underlying binary system ($P_{orb} \\equiv \\Omega ^{-1}$; we note, however that confirmation by radial velocity measurements is required). Another periodicity, $P\\sim62$~min, which was detected both in X-ray and in the optical band, was naturally interpreted as the spin period ($P_{spin} \\equiv \\omega^{-1}$) of the magnetic white dwarf in an intermediate polar system (R94). Further X-ray observations with ASCA and RXTE confirmed the suggested classification of the system, and refined the spin period to $P_{spin}=61.73\\pm13$~min (Hellier, Beardmore \\& Mukai 2002).  Here we report on the results of a photometric campaign of V1062~Tau in 2002 and 2003, during which the object exhibited unusual activity both on short time scales and on long ones ",
        "conclusions": "\\subsection{Long-term variations} Our data suggest a low state of V1062~Tau in 2002 January, only the ascent of which was observed (Sec.~\\ref{ObsLongTerm}). In the (sparse) published material of V1062~Tau we find three additional occasions in which V1062~Tau was possibly observed in low-state. The object was observed by R94 in five different occasions between 1988 February 22 -- 1990 December 11. In one (1988 February 22), the object was 3-4 times fainter. Munari, Zwitter \\& Bragaglia (1997) present a spectrum of the star from 1995 November 4, which yielded an integrated $V$ magnitude of $16.6$ -- about a magnitude fainter than the star's normal brightness level which was reported by R94. Finally, we  note that R94 reported that V1062~Tau is UV bright (indeed this was the criterion by which they identified the optical counterpart of the  X-ray source). A blank spectrum which was obtained after four hours of integration with IUE on 1994 November 3 (Szkody \\& Silber 1996) may imply a low state of the star during this epoch.  Low state are observed in all well-studied polars, as well as in many non-magnetic CVs (see Warner 1999 for a recent review of low states in CVs). However only two IPs (both with orbital periods $P_{orb}\\la$4~h) have been observed in low states. Low states appear to come in two types. 'slightly low' states are brightness reduction with an amplitude $\\la$1~mag, usually having recurrence time-scales which are characteristic for each object (typically 5--10 yr). This type is observed only in non-magnetic CVs. In contrast 'very low' states are more extreme drops in brightness, with no characteristic recurrence time-scales, occurring both in magnetic and non-magnetic CVs.  Although commonly attributed to reduced mass-transfer from the secondary star, the underlying mechanism of low states is still poorly understood. 'Very low' states have been proposed to be the result of migration of star-spots through L1, reducing the  Roche-lobe overfill of the secondary at L1 (Livio \\& Pringle 1994). 'Slightly low' states may be the effect of magnetic cycles in the secondary star (see e.g. Warner 1988).  Garnavich and Szkody(1988) found in a sample of magnetic and non-magnetic CVs that for systems with orbital periods $P_{orb}\\ga200$~min, the magnitude difference between low- to high-state increases with $P_{orb}$. For V1062~Tau, low states are expected to be $\\ga$6~mag fainter than the normal state. That out of three recorded low-states in V1062~Tau, two were observed only $\\sim$1~mag fainter than its mean brightness may suggest that low states of V1062~Tau do not follow the expected trend. Alternatively, the time V1062~Tau spends in low states may be short relative to the ascension/descension times.  There are a number of IPs which undergo dwarf-nova like outbursts. However, two systems exhibit shorter, low amplitude outbursts, which are suggested to be the result of increased mass transfer from the secondary (Hellier et al. 2000; Hellier, Mukai \\& Beardmore 1997): V1223~Sgr (only one observed outburst, van Amerongen \\& van Paradijs, 1989), and TV~Col (a few short outbursts, e.g. Hellier \\& Buckley, 1993). Another IP, EX~Hya, shows outbursts on intermediate time scales ($\\sim2-3$~d; Hellier et al. 2000)). It is yet unclear whether these events are the result of disc-instabilities, or mass-transfer enhancements (Hellier et al. 2000).  The 2002 February 19 outburst of V1062~Tau, which was only  partially recorded (Sec.~\\ref{ObsLongTerm}),  should probably be classified as a short burst. The 0.14~mag decline at ending section of the outburst light curve probably marked the start of the decline from maximum brightness (Fig.~\\ref{figLC}, inset). Additionally, the upper limit for the duration of the 2002 February outburst is $T\\le4.9$~d -- less that half the expected duration for a dwarf nova-like outburst of a $P_{orb}\\sim$10~h CV (Warner, 1995). The outburst of 2002 August 13, which was recorded only by a single snapshot, is also consistent with a short-type burst (Sec.~\\ref{ObsLongTerm}). Therefore, V1062~Tau is probably a new member in the exclusive group of IPs having short outbursts.  Finally, the different characteristics of the outburst of 2002 February 19, and of the 2002 Jan low state suggest that the two are caused by different instability mechanisms in the secondary star. The superposition of the short outburst on the recovering branch of the low state present a remarkable interplay between the two different instabilities.  \\subsection{Short-term signals} During 2001 December -- 2002 March V1062~Tau oscillated with either the spin, or the orbital sideband period, switching three times between the two. The transition time-scale between these two modes was $\\sim$2~d. The characteristic time scale for maintaining a period was $\\sim$5--$\\sim$25~d (Sec.~\\ref{ObsFourier}). The amplitude of both periods was correlated with the orbital phase, with different phase lags in the different subsets.  Modulations at the spin period are the signature of a disc-fed IP system. Possible sources of such variations are aspect variation of the accretion zone (e.g. Hellier et al. 1987), or reprocessing of X-ray emission in axisymmetric regions of the disc, either the inner edge of the magnetically-truncated disc, and/or the outer rim of a concave disc (Warner 1986). Sideband modulations may also occur in disc-fed systems, arising from reprocessing of X-ray light in fixed locations in the binary reference frame: the secondary star (Patterson \\& Price 1981), and/or the bright spot bulge on the outer rim of the accretion disc (Hassall et al. 1981). In predominantly stream-fed IPs, optical sideband modulations may also arise from the accretion stream flipping between the magnetic poles twice every synodic cycle.  The detection of both $P_{spin}$ and $P_{OSB}$ in the $I$-band light curve of V1062~Tau, and the absence of both in $V$ (see Sec.~\\ref{ObsFourier} and Fig.~11 of R94), suggest that the source of the two signals is more likely to be reflection rather than emission from the hotter accretion region. The strong variation of the amplitude of both signals over the orbital cycle (sec.~\\ref{SecObsCorr}, Fig.~\\ref{figAmp}) also supports reflection as the source of the modulations. Had these variations emanated from the accretion region of the CV system, strong variability of the amplitudes of the short periodicities would require the obscuration of most of this inner region during some part of the orbital cycle. However the orbital wave form is hardly an eclipsing one. If the accretion region is not obscured, any flux modulations it may exhibit should not be affected by the orbital cycle. Thus, measured in magnitude units, changes in the amplitude of such modulations over orbital cycle should only reflect the varying relative contribution of the fixed signal the to the changing total brightness of the system. This is hardly the case in V1062~Tau (Fig.~\\ref{figAmp}). Finally, the changes in the characteristics of the short signals between different data sets (Fig.~\\ref{figAmp}) would be difficult to account for if these variation are emission from the accretion region.  On the other hand, a simple reflection model also seems insufficient to account for the properties of the short-term modulations. In particular, the changes between the different data sets and the form of the amplitude dependence on the orbital phase introduce similar difficulties when ascribing the short periodicities to reflection.  The transitions between $P_{spin}$ and $P_{OSB}$ imply changes in the accretion mode and/or in the geometry of the system. A similar effect had been observed in the X-ray light of TX~Col -- a long period IP ($P_{orb}=5.72$~h), in which both the spin and the sideband are observable. The relative intensities of the two signals were found to vary by an order of magnitude over a typical time-scale of $\\la$1~month (Buckley \\& Tuohy 1989; Norton et al. 1997; Wheatley 1999). Changes in the relative power of the sideband signal were also observed in FO~Aqr (Beardmore et al. 1998), AO~Psc, V1223~Sgr and BG~CMi (e.g. Hellier 1998). Such changes are usually attributed to variable amounts of disc- and disc-overflow accretion. The cause of such transitions is not clear. X-ray observations of TX~Col suggest that periods of enhanced disc-overflow are related to decreased $\\dot{M}$ (Wheatley 1999). However changes in the accretion mode may also be driven by various mechanisms in the binary system  which are not related to changes in $\\dot{M}$ (see, e.g., Norton et al. 1997).  The period switching which we observed in the optical light of V1062~Tau may also reflect changes in the accretion mode, between predominantly disc-accretion and predominantly disc-overflow accretion. The exclusive appearance of the two signals may imply that during periods of predominantly disc-overflow accretion, disc-accretion became negligible or ceased completely, making the disc a 'nonaccretion-disc' during these periods (see King \\& Lasota 1991). The average brightness of V1062~Tau was not considerably different during periods in which the system's light  was modulated by different a short-period modulation (see Fig.~\\ref{figLC}; particularly, the average brightness of set {\\bf C} and that of set {\\bf D} are indistinguishable). However, we note the close coincidence of period switching events with the low-state/outburst activity of the star: one transition from $P_{OSB}$ to $P_{spin}$ coincided with the apparent low state of 2002 Jan, and another occurred shortly after the short outburst of 2002 February 19.  Period switching may also be caused as a result of changes in scale-height of the accretion disc, if both short signals emanate from reprocessed light (e.g. the beat is reflection from the  bright spot bulge and/or from the secondary star, and the spin  is reflection from the outer rim of a concave disc). Orbital sideband modulations would occur during periods when the accretion disc is thin. Conversely, when the disc rim thickens considerably  (e.g. following changes in mass-transfer rate from the secondary), it may obscure the secondary star, and the axi-symmetric outer rim of the disc would become the main source of reflection, varying at the spin period. It would be hard, however to reconcile such a change in the accretion disc with the stable average magnitude of V1062~Tau over the periods of changing short modulation."
    },
    "0402/astro-ph0402045_arXiv.txt": {
        "abstract": "It is embarrassing that 95\\% of the universe is unaccounted for. Galaxies and  larger-scale cosmic structures  are composed mainly of  `dark matter'  whose nature is still unknown. Favoured candidates  are weakly-interacting particles that have survived from the very early universe, but more exotic options  cannot be excluded. (There are strong arguments that the dark matter  is not composed of baryons).  Intensive experimental searches  are being made for the `dark'  particles (which pervade our entire galaxy), but we have indirect clues to their nature too. Inferences   from galactic dynamics and gravitational lensing allow astronomers to `map' the dark matter distribution; comparison with numerical simulations of galaxy formation can constrain (eg) the particle velocities and collision cross sections. And, of course, progress in understanding the extreme  physics of the ultra-early universe could offer clues to what particle might  have existed then, and how many would have survived.   The mean cosmic density of dark matter (plus baryons) is  now pinned down  to be only about 30\\% of the so-called  critical density corresponding to a `flat' universe. However, other recent  evidence -- microwave background anisotropies, complemented by data on distant supernovae -- reveals  that our universe actually is `flat', but that its dominant ingredient  (about 70\\% of the total mass-energy) is something quite unexpected ---  `dark energy' pervading all space,  with negative pressure.   We now confront two mysteries:  (i) Why does the universe have three quite distinct  basic ingredients -- baryons, dark matter and dark energy -- in the proportions (roughly)  5\\%, 25\\% and 70\\%?  (ii) What are the (almost certainly profound)  implications of the  `dark energy' for fundamental physics? ",
        "introduction": "\\subsection{Baryons}  The inventory of cosmic baryons is readily compiled. Stars and their remnants, and gas in galaxies, contribute no more than 1\\% of the critical density (i.e. they give $\\Omega_b   < 0.01$). However  several percent  more could be contributed by diffuse material pervading  intergalactic space: warm gas (with $kT \\simeq 0.1$ keV) in groups of galaxies and loose clusters, and cooler gas pervading intergalactic space that  manifests itself via the `picket fence' absorption lines in quasar spectra. (Rich clusters are rare, so their conspicuous gas content, at several KeV, is not directly significant for the total inventory,  despite its  importance as a probe)  These baryon estimates are concordant with those inferred by matching the He and D abundances at the birth of galaxies with the predicted outcome of  nucleosynthesis in the big bang, which is sensitive to the primordial baryon/photon ratio, and thus to   $\\Omega_b$.  The  observational estimates have firmed up, with improved measurements of  deuterium  in high-$z$ absorbing clouds. The best fit occurs for $\\Omega_b \\simeq 0.02h^{-2}$ where $h$ is the Hubble constant in units of 100 km s$^{-1}$ Mpc$^{-1}$. Observations favour $h \\simeq 0.7$.  $\\Omega_b$  is now pinned down by  a variety of argument to be $0.04 - 0.05$. This corresponds to only $\\sim 0.3$ baryons per cubic metre, a value so low that it leaves little scope for dark  baryons. (It is therefore unsurprising that the MACHO/OGLE  searches should have found that compact objects do not make a substantial  contribution to the total  mass of our own galactic halo.)  \\subsection{How much dark matter?} An important recent development is that $\\Omega_{DM}$ can now be constrained to a value  around 0.25 by several independent lines of evidence:  (i)   One of the most ingenious and convincing arguments  comes from noting that baryonic matter in clusters -- in galaxies, and in intracluster gas -- amounts to  $0.15-0.2$ of the inferred virial mass (White {\\it et al.} 1993) . If clusters were a fair sample of the universe, this would then be essentially the same as the cosmic ratio of baryonic to total mass.    Such an argument could not be applied to an individual galaxy,  because  baryons segregate towards the centre. However, there is no such segregation on the much larger scale of clusters: only a small correction is necessary to allow for  baryons expelled during the cluster formation process.  (ii)   Very distant galaxies appear  distorted, owing to gravitational lensing by intervening galaxies and clusters.   Detailed modelling of  the mass-distributions needed to cause the observed distortions yields a similar estimate. This is a straight measurement of $\\Omega_{DM}$ which (unlike (i)) does not involve assumptions about $\\Omega_b$, though it does depend on having an accurate measure of the clustering amplitude.  (iii) Another argument is based on the way density contrasts grow during the cosmic expansion: in a low density universe, the expansion kinetic energy overwhelms gravity, and the growth of structure saturates at recent epochs. The existence of conspicuous clusters of galaxies with redshifts as large as $z=1$ is hard to reconcile with the rapid recent growth of structure that would be expected if $\\Omega_{DM}$ were unity. More generally, numerical simulations based on the cold dark matter (CDM) model model are a better fit to the present-day structure for this value of $\\Omega_{DM}$(partly because the initial fluctuation spectrum has too little long-wavelength power if $\\Omega_{DM}$ is unity).  Other methods will soon offer independent estimates. For instance, $\\Omega_{DM}$ can be estimated from  the  deviations from the Hubble flow induced by large-scale irregularities in the  mass distribution on supercluster scales.  \\subsection{What could the dark matter be?}  The  dark matter  is not  primarily baryonic. The amount of deuterium calculated to emerge from the big bang  would  be  far lower than observed if the average  baryon density  were $\\sim 2$ (rather than $\\sim 0.3$)   per cubic metre. Extra exotic particles that do not participate in nuclear reactions, however, would not scupper the concordance.  Beyond the negative statement that it is non-baryonic, the nature of the dark matter  still eludes us.  This key question may yield to a three-pronged attack:  1. $\\underline{\\hbox {Direct detection}}$. As described by other contributors to this meeting, several groups are developing cryogenic detectors for supersymmetric particles and axions This is an exciting quest. Of course, not even optimists can be confident that the actual dark matter particles have parameters within the range that these experiments are yet sensitive to. But the stakes are high: detection of most of the gravitating stuff in the universe, as well as a new class of elementary particle. So it seems well worth committing to these experiments  funding that is equivalent to a small fraction of the cost of a major accelerator  2. $\\underline{\\hbox {Progress in particle physics}}$. Important recent measurements suggest that neutrinos have non-zero masses; this result has crucially important implications for physics beyond the standard model. The inferred neutrino masses seem, however, too low to be cosmologically important.  If the masses and cross-sections of supersymmetric particles were known, it should be possible to predict how many survive, and their contribution to $\\Omega$, with the same confidence with which we can compute the nuclear reactions that control primordial nucleosynthesis. Associated with such progress, we might expect a better understanding of how the baryon-antibaryon asymmetry arose, and the consequence for $\\Omega_b$.  Optimists may hope for progress on still more exotic options.  3. $\\underline{\\hbox {Simulations of galaxy formation and large-scale structure}}$.  When and how galaxies form, the way they are clustered, and the density profiles within individual systems, depend on what their gravitationally-dominant constituent is. A combination of better data and better simulations is starting to set generic constraints on the options. The CDM model works well. But there are claimed discrepancies, though many of us suspect these may ease when the galaxy formation process is better understood. For instance the centre of a halo would, according to the simulations, have a `cusp' rather than the measured uniform-density core: this discrepancy has led some authors to explore modifications where the particles are assumed to have significant collision probabilities, or to be moving with non-negligible velocities (i.e. `warm' not cold.). These calculations are in any case offering interesting constraints on  the properties of heavy supersymmetric particles. (Also, straight astronomical observations can rule out a contribution to $\\Omega$  of more than 0.01 from neutrinos -- this is compatible with current experimental estimates.) ",
        "conclusions": ""
    }
}